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Guandique, M. Segrura) for the support they offered in data retrieval. Without them this work would hâve never existed. I would also like to thank S. Vinatier and M. Hirtzig from Paris Observatory for their help in scientific interprétation, G. Kollias and S. Stamogiorgos for helping in software, D. Couritis and T. Hadzistergos for helping in future missions chapter, as well as R. Kirk for the fruitful discussions on RADAR image processing. It is also a great pleasure to acknowledge E. Bratsolis for his guidance in despeckle filtering technique and A. Hayes for providing RADAR data and interprétations of software. I owe particular thanks to my ffiends and colleagues A. Solomonidou for her valuable help during my Thesis and E. Mitsakou for her tolérance

The story of Titan begins on the night of March 25, 1655, when the young Dutch astronomer Christiaan Huygens (1629-1695), with his brother Constantijn, tumed their own designed and constructed 4 m long telescope to Satum. As he described it in his report entitled "De Saturni luna observatio nova", which was published a year later in the Hague, Huygens wanted to understand the strange protrusions referred to as ansae or handles appearing at both sides of the giant planet.

Huygens not only observed these odd extends of Satum but also by using his telescope with a magnification of 50 he noticed the existence of a bright star -as he initially thoughtwhich he eventually realized that it was a satellite. Huygens gathered ail his observations in Systema Saturnium which was published in 16591. Based on his observations, Huygens calculated the orbital period of Titan at 15 days, 23 hours, 13 minutes not far from the modem value of 15 days, 22 hours, 41 minutes and 24.35 seconds. Except for the discovery of Titan, Huygens explained that Satum's strange extends were in fact its rings.

Huygens did not name Titan with its contemporary name but he called it Luna Saturni.

In 1847, Sir John Frederick William Herschel in his publication Results of Astronomical Observations made at the Cape of Good Hope proposed to give the names of legendary Titans to Satum's satellites, inspired from the Greek Mythology. Following Herschel's suggestion, the biggest moon of the Satumian System is now known as Titan.

Catalan astronomer José Comas Solà assumed the existence of a strongly absorbing atmosphère surrounding Titan, when he claimed he observed limb darkening in his observations in 1907 [START_REF] Comas Solâ | Observations des satellites principaux de Jupiter et de Titan[END_REF]. Sir James Jeans included Titan [START_REF] Jeans | The dynamical theory ofgases[END_REF] in the 4th édition of his theoretical work of escaping atmospheric processes, perhaps being influenced by Solà's observations but he did not refer Solà in his sources list (p. 348). Pioneer's follow-ups were the legendary Voyager 1 and 2 both launched in 1977 (Fig. The scientific importance of the Voyager missions' measurements is enormous, as they hâve visited ail the giant planets in the Solar System retuming to the Earth a wealth of data and images, which hâve greatly enhanced our understanding of the Solar System. Apart from the Voyagers no other human device has visited Uranus and Neptune or travelled so far away from our planet. Both spacecraft are still in operation at the edge of the Solar System.

Nowadays, Voyager 1 and 2 crossed the termination shock of the Heliosphere in December 2004 and in August 2007 respectively, where the velocities of the plasma significantly decrease due to the interaction with the interstellar medium and they are now exploring the Heliosheath. They will enter into the interstellar space possibly before 2020.

The spacecraft are still détectable through the Deep Space Network (DSN). From the above payload, the UVS, the MAG, the LECP, the CRS and the PWS are still in operation after 35 years.

Voyager 1 encountered Titan from north to south, having an orbital angle at 8.7°in respect to Titan's orbital plane with a relative speed of 17.3 km/s [START_REF] Coustenis | L' atmosphère de Titan a partir des observations infrarouges de Voyager[END_REF]. Both

Voyagers flew by Titan close to its northem spring equinox at solar longitude Ls=9°. The Voyagers' encounters of Titan gave a wealth of data for the nature of the biggest Kronian satellite. Their observations confirmed the presence of Titan's optically thick scattering haze and determined an optical depth larger than 5 with a peak haze brightness at 240 km above its surface [START_REF] Smith | Encounter with Satum: Voyager 1 Imaging Science Results[END_REF].

The haze, which covered the Southern hemisphere, was significantly brighter compared to the northem one. The northem polar régions in particular showed darker haze colors, this was called the northem polar hood. This North-South asymmetry was explained by variations in density, particle size and composition due to different solar heating. Indeed, the northem hemisphere was 10 K colder than the southem one at the VI era [START_REF] Flasar | Titan's atmosphère: température and dynamics[END_REF].

The other feature that Voyager missions recorded was the detached haze layer which was discovered at 100 km above the peak of the main haze [START_REF] Smith | Encounter with Satum: Voyager 1 Imaging Science Results[END_REF][START_REF] Mcdonald | Mid-ocean ridges: fine scale tectonic, volcanic and hydro thermal processes within the plate boundary zone[END_REF]. A global distribution of the detached haze with a latitudinal variation southem of 45°N at altitudes 300-350 km was reported from Voyager 2 recordings [START_REF] Rages | Vertical distribution of scattering hazes in Titan's upper atmosphère[END_REF].

Above 60 N, the detached haze merged with the main haze.

The Voyager 1 remote data determined the bulk composition and the vertical température profile of Titan's atmosphère. Both Vl/RSS and Vl/IRIS estimated the mean mass ratio to be within the range of 27.8 and 29.3 amu (Lindal et al., 1983;[START_REF] Lellouch | Titan's atmosphère and hypothesized océan: A reanalysis of the Voyager 1 radio-occultation and IRIS 7.7-pm data[END_REF]. In combination with the Vl/UVS observations [START_REF] Broadfoot | Extrême ultraviolet observations from Voyager 1 encounter with Satum[END_REF][START_REF] Strobel | EUV émission from Titan's upper atmosphère -Voyager 1 encounter[END_REF] and solar occultation data [START_REF] Smith | A New Look at the Satum System: The Voyager 2 Images[END_REF] it was then determined that nitrogen is the dominant species in Titan's atmosphère, while methane is a minor constituent.

Additionally, only upper limits were estimated for argon [START_REF] Strobel | Upper limit on Titan's atmospheric argon abundance[END_REF]. VI data also detected hydrogen but failed to detect ammonia [START_REF] Hanel | Infrared observations of the Satumian System from Voyager 1[END_REF]. As far as the methane stratospheric abundance is concemed, [START_REF] Lellouch | Titan's atmosphère and hypothesized océan: A reanalysis of the Voyager 1 radio-occultation and IRIS 7.7-pm data[END_REF] gave a methane mixing ratio up to 3.4% assuming that it is not supersaturated in the stratosphère.

During the radio occultation experiment, the satellite occulted the spacecraft. The phase shift as well as the atténuation of the VI radio signal derived two vertical profiles [START_REF] Smith | A New Look at the Satum System: The Voyager 2 Images[END_REF]. The Vl/RSS also indicated the Earth-like thermal structure of the atmosphère. It estimated the surface température at 94±0.7K at 1496 ±0.02 mbar and a radius of 2575±0.5 km. The tropopause was located at 42 km (71.4±0.5K, 130 mbar) while the stratopause at about 200 km (0.75 mbar) with a highest value at 170K. The exact altitude and the température of the stratopause were not determined. Vl/UVS gave also constraints for the upper atmosphère (186±20K at 1265 km) [START_REF] Smith | A New Look at the Satum System: The Voyager 2 Images[END_REF] and the interval was calculated by interpolation techniques. However, recent reanalysis of Voyager data showed the location of homopause higher than 1000 km [START_REF] Vervack | New perspectives on Titan's upper atmosphère from a reanalysis of the Voyager 1 UVS solar occultations[END_REF], while [START_REF] Smith | A New Look at the Satum System: The Voyager 2 Images[END_REF] showed it at 925 km.

Vl/IRIS spectra confirmed the existence of a rich organic environment in Titan's atmosphère. Indeed, the émissions of acetylene (C2H2), ethane (C2H6), ethylene (C2H4), propane (C3H8), methylacetylene (CH3C2H), diacetylene (C4H2), hydrogen cyanide (HCN), cyanoacetylene (HC3N), cyanogen (C2N2) and carbon dioxide (CO2) hâve been recorded by IRIS in the stratospheric layers [START_REF] Kunde | C4H2, HC3N and C2N2 in Titan's atmosphère[END_REF][START_REF] Maguire | C3FL and C3H4 in Titan's atmosphère[END_REF][START_REF] Samuelson | Mean molecular weight and hydrogen abundance of Titan's atmosphère[END_REF]Samuelson, 1983;Coustenis et al., 1989a;[START_REF] Coustenis | Titan's atmosphère from Voyager infrared observations. II -The CH3D abundance and D/H ratio from the 900-1200/cm spectral région[END_REF].

The VI encounters probed Titan's atmosphère in several latitudes (70°N, 50°N, 30°N, 9°N, 8°N, 5°N, 7°S, 30°S and 53°S) and a latitudinal distribution of trace gases abundances was constructed (Fig. 1.3) while the longitudinal variation was not significant [START_REF] Coustenis | Titan's atmosphère from Voyager infrared observations. IV: Latitudinal variations of température and composition[END_REF]. Except for CO2 and ethane, ail the trace constituents showed an enhancement in the high northem latitudes. [START_REF] Coustenis | Titan's atmosphère from Voyager infrared observations. IV: Latitudinal variations of température and composition[END_REF].

Moreover, IRIS limb data indicate that the mixing ratios of most of the species increase with altitude above their condensation level (Coustenis et al., 1991). V2/IRIS observations, which probed Titan few months later, derived the same results [START_REF] Letoumeur | Titan's atmospheric structure from Voyager 2 infrared spectra[END_REF]. The D/H ratio in Titan's atmosphère was determined to be 3-12 times higher than the protosolar value [START_REF] Kim | The abundance of CH3D in the atmosphère of Titan, derived from 8-to 14micron thermal émission[END_REF][START_REF] Coustenis | Titan's atmosphère from Voyager infrared observations. II -The CH3D abundance and D/H ratio from the 900-1200/cm spectral région[END_REF].

1.1.4

Ground-based and space-borne Titan observations A full Titan révolution of the Sun takes about 30 years and therefore long-term observations are needed to hâve a complété picture of its temporal atmospheric évolution. Taking as the initial point of our temporal studies the VI encounter in 1980, the mid-2009 stands as an important time spot in Titan's atmospheric science.

Unfortunately, due to the hazy nature of its atmosphère, distant direct observations of its lower atmosphère and of its surface are impeached. Methane bands strongly absorb the outgoing radiation and it is impossible to see through them except only through narrow methane spectral "windows" (régions of weak absorption) in the near infrared (NIR). These

Windows are centered at 0.83, 0.94, 1.07, 1.28, 1.58, 2.0, 2.9 and 5.0 pm and the observations are limited.

Several ground-based and space-borne observations hâve been performed in the time interval between the Voyager encounter of Titan (1980) and the arrivai of the Cassini-Huygens mission (2004). The advent of adaptive optics and the space-borne télescopes (Hubble Space Telescope -HST, Infrared Space Observatory -ISO) made the surface and the lower atmosphère of Titan reachable from the Earth.

The adaptive optics technique retums a more accurate image of the object through the atmosphère by correcting the aberrations of the light sources. Aberrations mainly originated from atmospheric turbulence, thermal blooming and other non-atmospheric sources [START_REF] Tyson | Principles ofAdaptive Optics[END_REF]. This technique allows resolving Titan's disk and obtaining some spatial resolution on the surface of Titan by using both spectroscopy and imaging. Distant observations of Titan can be performed by radio télescopes as well. Developments in signal processing by using heterodyne receivers amplify the echoes from Titan providing high resolution spectroscopy (Coustenis & Taylor, 2008).

Finally, information about Titan's atmosphère can be provided by stellar occultations of Satum and Titan when they pass in front of a star. Six such events hâve been recorded so far: one on 3 July 1989, one on 21 August 1995, two on December 2001 and two on 14

November 2003 (see [START_REF] Marty | Erratum: Kronos: exploring the depths of Satum with probes and remote sensing through an international mission[END_REF] and references therein).

Both Earth-and space-based observations hâve revealed the heterogeneous nature of the surface and showed that dark and bright régions, consisting of different ice, dominate it.

Such observations came from the HST in 1994 (Smith et al., 1996) and 1997-1998(Meier et al., 2000), the speckle interferometry from Keck Telescope within 1996-1998 (Gibbard et al., 1999;[START_REF] Gibbard | Speckle imaging of Titan at 2 microns: surface albedo, haze optical depth, and tropospheric clouds 1996-1998[END_REF] and adaptive optics from Canada-France-Hawaii Telescope (CFHT) in 1998 (Coustenis et al., 2001), from Very Large Telescope in Chile (VLT) in 2002 [START_REF] Gendron | VLT/NACO adaptive optics imaging of Titan[END_REF], from combination of CFHT and VLT data during 2001-2004[START_REF] Coustenis | Maps of Titan's surface from 1 to 2.5 pm[END_REF] and from W.M. Keck II telescope within 2001-2003(Roe et al., 2004a).

As far as the composition of the surface is concemed, observations from the Infrared Telescope Facility (IRTF) on Mauna Kea, Hawaii, in 1989 showed that the surface consists of dirty ice bedrock (Griffith et al., 1991) fiill of organics and tholins3, a laboratory analogue to Titan's surface material [START_REF] Sagan | Organic chemistry in the atmosphère[END_REF]Sagan & Khare, 1979;[START_REF] Sagan | Organic matter in the Satum System[END_REF]. A mixture of water ice and tholins, which covered parts of the surface, was also discovered by albedo observations taken by the CFHT from 1991 to 1996 [START_REF] Negrao | Titan's surface albedo variations over a Titan season from near-infrared CFHT/FTS spectra[END_REF]. Additionally, observations from the VLT in 2000 also assumed also ammonia ice on Titan's surface [START_REF] Lellouch | Titan's 5-pm window: observations with the Very Large Telescope[END_REF].

More recent observations in 2005 assumed the presence of methane ice mixed with water and tholins ice by the use of the VLT measurements [START_REF] Hirtzig | Titan: Atmospheric and surface features as observed with Nasmyth Adaptive Optics System Near-Infrared Imager and Spectrograph at the time of the Huygens mission[END_REF][START_REF] Negrao | The 2-pm spectroscopy of Huygens probe landing site on Titan with Very Large Telescope/Nasmyth Adaptive Optics System Near-infrared Imager and Spectrograph[END_REF], while VLT measurements in 2004 also gave firm limits of carbon dioxide ice [START_REF] Hartung | Limits to the abundance of surface C02 ice on Titan[END_REF]. However, higher resolution observations needed to distinguish and identify surface features.

Ground-based RADAR observations also contributed significantly to the study of the surface of Titan. After the Voyager encounters, the existence of a global océan was suggested to cover the surface (Flasar, 1983;Lunine et al., 1983). However, earth-based RADAR echoes recorded at Very Large Array in New Mexico indicated that Titan's surface possesses a very high RADAR cross section and latitudinal variation, which is inconsistent with the global océan hypothesis (Muhleman et al., 1990). Lemmon et al. using

data from Steward

Observatory and the Multiple Mirror Telescope in 1992 also showed the spatial heterogeneity of Titan [START_REF] Lemmon | Titan's Rotation: Surface Feature Observed[END_REF], while Cambell et al. noted RADAR observation areas of specular reflections, similar to small lake-like features from Arecibo echoes in 2001 and 2002 [START_REF] Campbell | Radar Evidence for Liquid Surfaces on Titan[END_REF].

Carbon monoxide was detected in Titan's atmosphère at 1.57 pm (Lutz et al., 1983), while its existence was established by several measurements at 5 pm, determining that CO is not uniformly mixed in the atmosphère [START_REF] Marty | Erratum: Kronos: exploring the depths of Satum with probes and remote sensing through an international mission[END_REF] (and references therein).

Using IRAM-30 data taken between April 1996 andDecember 1999, acetonitrile (CH3CN) was firstly detected [START_REF] Bezard | Détection ofAcetonitrile on Titan[END_REF][START_REF] Marten | New Millimeter Heterodyne Observations of Titan: Vertical Distributions of Nitriles HCN, HC3N, CH3CN, and the Isotopic Ratio 15N/14N in Its Atmosphère[END_REF]. These measurements provided 3 Tholins are named after the greek word "0oA6ç" (=mud) by Cari Sagan.

compositional studies of the atmosphère and the first vertical profiles of nitriles, which increased with altitude in the stratosphère, outside the polar région. Moreover, these IRAM-30 recordings gave a strongly enhanced ratio of 15N/14N by a factor of 4 in HCN. The vertical stratospheric HCN profile in Titan was reported from earlier IRAM-30 observations on September 1986 and May 1987 [START_REF] Tanguy | Stratospheric profile ofHCN on Titan from millimeter observations[END_REF] as well as in May 1995 [START_REF] Hidayat | Millimeter and Submillimeter Heterodyne Observations of Titan: Retrieval of the Vertical Profile of HCN and the 12C/13C Ratio[END_REF].

By using heterodyne spectroscopy at the NASA Infrared Telescope Facility (IRTF) in August 1993, October 1995and September 1996, the abundance of ethane was retrieved [START_REF] Kostiuk | Ethane abundance on Titan[END_REF][START_REF] Livengood | NOTE: Improved Détermination of Ethane (C2H6) Abundance in Titan's Stratosphère[END_REF].

The higher resolution of echelle technique applied in IRTF (TEXES) compared to

Voyager, allowed for the first time the séparation of propane V26 band at 748 cm'1 in December 2002, which was blended with the acetylene P-wing in the Vl/IRIS spectra [START_REF] Roe | Propane on Titan[END_REF]. Observations taken during the 1999-2002 with W.M. Keck I telescope showed seasonal variation of ethylene mixing ratio in the atmosphère at 8-13 pm (Roe et al., 2004b).

Keck II/NIRSPEC observations in November 2001 provided abundances for HCN, C2H2, CH4 and CH3D at 3 pm [START_REF] Geballe | High-Resolution 3 Micron Spectroscopy of Molécules in the Mesosphere and Troposphère of Titan[END_REF][START_REF] Kim | Clouds, haze, and CH4, CH3D, HCN, and C2H2 in the atmosphère of Titan probed via 3 pm spectroscopy[END_REF]. Ground based measurements were available during the Huygens descent in January 2005 being obtained at National Astronomical Observatory (NAO) providing ethane abundances [START_REF] Livengood | High-resolution infrared spectroscopy of ethane in Titan's stratosphère in the Huygens epoch[END_REF].

IRTF/IRHS data mentioned above, were used for detecting winds in Titan's atmosphère by studying ethane in two limb positions along with December 2003 measurements at NAO and January 2005 recordings at NAO/HIPWAC, noting an équatorial wind [START_REF] Kostiuk | Titan's stratospheric zonal wind, température, and ethane abundance a year prior to Huygens insertion[END_REF][START_REF] Kostiuk | Stratospheric global winds on Titan at the time of Huygens descent[END_REF]2010). Similarly, European Southern Observatory Very Large Telescope with UVES measured global circulation wind velocities at various altitudes at the time of the Huygens landing in January 2005 [START_REF] Luz | Characterization of zonal winds in the stratosphère of Titan with UVES[END_REF].

ISO focused on Titan for 20 h in 1997. By viewing Titan only as a small disk, it performed disk-averaged recordings. Having much higher resolution than the Vl/IRIS, it reported the same organic inventory as IRIS did (Coustenis et al., 2003;Coustenis & Taylor, 2008). Additionally, ISO detected water vapor for the first time in the atmosphère (Coustenis et al., 1998). It also detected benzene for the first time at 674 cm'1 (Coustenis et al., 2003), whose presence on Titan was later confirmed in Cassini/CIRS observations.

The stellar occultation of Satum and Titan by the star 28 Sagitarii allowed retrieving information about Titan's stratosphère within a range of 250 to 500 km. During this event, two haze layers were detected, having a reverse asymmetry comparing to the Voyager 1 observations [START_REF] Sicardy | Probing Titan's atmosphère by stellar occultation[END_REF]).

1.2

The Cassini-Huygens Mission

Mission development

The need for a space mission focused on the Satumian System, its rings, its satellites and its magnetosphere, led scientists to propose the Cassini-Huygens mission, an ambitious program of scientific in situ observations. Without doubt, the Cassini-Huygens mission is one of the largest, heaviest and most complex interplanetary spacecraft ever built in a cooperative effort between National Aeronautics and Space Administration (NASA), European Space Agency (ESA) and the Agenzia Spaziale Italiana (ASI). The formai beginning of the Cassini mission was with the mission proposai in 1982 to ESA and in 1983 to NASA. After its approval by both space agencies, the mission was developed between 1987 and 1996 [START_REF] Matson | The Cassini/Huygens Mission to the Saturnian System[END_REF]. It launched in 1997 and arrived in the Satumian System in July 2004.

The mission concept consists of two spacecraft: the Cassini orbiter and the Huygens probe, named after Giovanni Domenico Cassini (1625-1712) and Christiaan Huygens (1629-1695) respectively, honoring them for their contribution in the early stages of astronomical observation of the Satumian System.

The Cassini orbiter

The Cassini orbiter is 6.8 m high and 4 m at its widest in diameter. It functions as an artificial satellite of Satum and was the carrier of the Huygens probe which was successfully delivered to Titan. It also operated as the link between the Huygens capsule and the Earth during the Huygens mission phase in 2005.

The orbiter's scientific objectives are to deeply investigate Satum's atmosphère, its rings, its icy satellites, the radiation belts and its magnetosphere. Because of Satum's distance from the Sun, at circa 10 AU, Cassini cannot rely on solar panels as power sources and is equipped with three Radioisotope Thermoelectric Generators (RTG), using the heat from the natural decay of plutonium to generate direct current electricity powers.

Cassini is a State of the art spacecraft hosting on-board science instrumentation capable of performing both remote and in situ observations (Fig. and its energy sources [START_REF] Henry | An Introduction to the Design of the Cassini Spacecraft[END_REF].

Orbiter in situ investi eations [START_REF] Young | Cassini Plasma Spectrometer Investigation[END_REF].

1. The Cassini Plasma Spectrometer (CAPS) explores the nature of the Satumian plasma within and near Satum's magnetic field. CAPS help us understand the complexity of the Satumian magnetosphere and its interaction with its satellites, the rings and the solar wind. It is consisted of Electron Spectrometer (ELS), the Ion Beam Spectrometer (IBS) and the Ion Mass Spectrometer (IMS). The ELS measures électron velocity distribution ranging from 0.6 eV to 28,250 eV, while the IBS records distributions of ion velocities within the range of .1 eV to 49,800 eV. IMS measures the composition of the hot, diffuse plasma from the Satumian magnetosphere and low concentration ion species from 1 eV to 50,280 eV [START_REF] Young | Cassini Plasma Spectrometer Investigation[END_REF]).

The Ion and Neutral Mass Spectrometer (INMS) detects the neutral and charged

particles located near Titan, Satum and the other moons to leam more about their extended atmosphères and ionosphères. Thus, the INMS focuses on the upper part of Titan's atmosphère, between the altitudes of 900 and 1000 km, where the complex photochemistry of nitrogen and methane begin to build more complex organic molécules, which descend towards its surface. Moreover, INMS also studies the interaction of Titan's atmosphère with Satum's magnetospheric plasma. It consists of two ion sources, one closed and one open. The former opérâtes with the non-reactive neutrals, the nitrogen and methane, while the latter functions for reactive neutrals like the atomic nitrogen and for positive ions with energies less than 100 eV. The instrument is able to detect heavier hydrocarbon molécules such as benzene and détermine their molecular mass [START_REF] Waite | The Cassini Ion and Neutral Mass Spectrometer (INMS) Investigation[END_REF].

3. The Magnetosphere Imaging Instrument (MIMI) images Satum's magnetosphere, studies its dynamics and measures the interactions between the magnetosphere and the solar wind and the magnetosphere with Satum's atmosphère, Titan and the other icy INCA studies the interaction of Titan's exosphere with the Satumian magnetosphere every 90 s [START_REF] Krimigis | Magnetosphere Imaging Instrument (MIMI) on the Cassini Mission to Satum/Titan[END_REF].

4. The Dual Technique Magnetometer (MAG) studies Satum's internai magnetic field and its interactions with the solar wind, the rings and the moons of Satura. MAG also maps the magnetic State of Titan and its atmosphère as well as Satum's ring and dust interactions with the electromagnetic environment. MAG investigates the structure of the magnetotail and the dynamic processes therein. It consists of a fluxgate magnetometer and a vector/scalar hélium magnetometer mounted both on the spacecraft's boom (see Fig. 1.4). Hence, MAG senses small changes in fields spanning four orders of magnitude with high sensitivity, recording the strength and direction of the magnetic fields [START_REF] Dougherty | The Cassini Magnetic Field Investigation[END_REF].

5. The Radio and Plasma Wave Spectrometer (RPWS) studies radio émissions plasma waves, thermal plasma and dust. It consists of three orthogonal electric field antennas which focus on the détection of electric fields ranging from 1 Hz to 16 MHz, and three orthogonal search coil magnetic antennas to detect magnetic fields in the range within 1 Hz to 12 MHz. A Langmuir probe is used for measuring the électron density and température. The electric antennas, the search coils and the Langmuir probe are mounted on the orbiter's main body. RPWS measures the électron density and température in Titan's ionosphère and study the escape thermal plasma from Titan's wake région [START_REF] Gumett | The Cassini Radio and Plasma Wave Investigation[END_REF]).

6. The Cosmic Dust Analyzer (CDA) measures the ice and dust grains in and near the Satum System with masses ranging from 10'19 and 10"9 kg. CDA studies the physical properties of the ice and dust grains and their Chemical composition, as well as their interaction with the magnetosphere of Satum, its rings and satellites [START_REF] Srama | The Cassini Cosmic Dust Analyzer[END_REF].

Orbiter remote investigations:

Four instruments are mounted on the Cassini's Optical Remote Sensing Palette (RSP) (Fig. 1. The Radio Science Subsystem (RSS) searches for gravitational waves in the Universe, studies the atmosphère, rings and gravity fields of Satum and its moons by measuring telltale changes in radio waves sent from the spacecraft. It consists of a Kaband traveling wave tube amplifier, a translator, an exciter, a S-band transmitter and various microwave components. It détermines the température and composition profiles within Satum's and Titan's atmosphères as well as the températures and électron densities within Satum's and Titan's ionosphères (Kliore et al., 2004).

2. The Cassini RADAR (RADAR) mapper probes the surface of Titan using RADAR imaging and measures its topography as well. The RADAR is a multimode Ku-band instrument (13.8 GHz, X2.17 cm) having four operational modes: Synthetic Aperture RADAR (SAR) imaging, altimetry, scatterometry and radiometry. The RADAR also observes other targets in the Satumian System such as its icy moons and the giant planet (Elachi et al., 2004).

3. The Visible and Infrared Mapping Spectrometer (VIMS) searches for the Chemical compositions of the surfaces, atmosphères and rings of Satum and its moons by measuring colors of visible light and infrared energy emitted or reflected within the wavelength range 0.3 and 5.1 microns. Especially, VIMS probes the surface properties of Titan and maps them in a global scale. VIMS measurements are crucial for understanding Titan's geology and identifying any volcanic activity. It is consistent of two imagers, one in the visual part, the VIMS-VIS, and the other in infrared part of the spectmm, the VIMS-IR [START_REF] Brown | The Cassini Visual And Infrared Mapping Spectrometer (Vims) Investigation[END_REF].

4. The Composite Infrared Spectrometer (CIRS) (Flasar et al., 2004) measures infrared radiation emitted from the surfaces, atmosphères and rings of Satum and its moons to study their température and compositions. A detailed description of the instrument and its operational modes can be found in the following Chapter 2.

The Ultraviolet Imaging Spectrograph (UVIS) measures ultraviolet émission from

atmosphères and rings to study their structure, chemistry and composition. UVIS détermines the constituents' abundances in Titan's atmosphère, the aérosols in order to infer about the atmospheric circulation patterns. It also studies the UV émission from the upper atmosphère of Titan and its relation to the Satumian magnetosphere as well.

It hosts two moderate resolution telescope-spectrographs, providing images ranging from 56 to 118 nm (EUV) and from 110 to 190 nm (FUV) respectively [START_REF] Esposito | The Cassini Ultraviolet Imaging Spectrograph Investigation[END_REF]).

6. The Imaging Science Subsystem (ISS), a high-resolution two-dimensional optical device, takes pictures in visible, near-ultraviolet and near-infrared light of the objects within the Satumian System. It consists of two caméras, a narrow angle reflecting telescope and a wide-angle refractor. Both caméras operate at a spectral range from 1100 to 200 nm. ISS is also the optical navigation guide of the Cassini spacecraft [START_REF] Porco | Cassini Imaging Science: Instrument Characteristics And Anticipated Scientific Investigations At Satura[END_REF]).

The Huygens probe

The Huygens probe is the first man-made vehicle, which has performed a successfiil landing in such a distant place from the Earth. After a couple initial orbits of Satum, the orbiter has released the probe on a trajectory directed towards Titan. The Huygens mission can be divided into 3 phases: the Entry, the Descent and the Landing phase. It landed in Titan's équatorial région called Adiri at 10.3°S, 192.3°W. During its descent phase as well as after touchdown, the capsule revealed a solid surface under the thick haze. The Huygens descent lasted for 2 h 27 mins 50 s (Lebreton et al, 2009).

The orbiter then passed below the probe's horizon, breaking the established link between Huygens and the Earth [START_REF] Lebreton | An overview of the descent and landing of the Huygens probe on Titan[END_REF]. However, the Very Long Baseline Interferometry (VLBI) network, consisting of 17 large radio télescopes, recovered the Huygens signal [START_REF] Witasse | OverView of the coordinated ground-based observations of Titan during the Huygens mission[END_REF] and thus Huygens continued to send back to Earth more measurements. Eventually, before Huygens' five batteries ran out of energy, it had already submitted data of 2.5 h during its descent and 3 h and 14 min after its touchdown on the ground (Lebreton et al, 2009). The Huygens payload consists of six scientific instruments placed as shown in Fig. 1.7. This instrumentation has performed direct atmospheric and surface measurements of the environment of Titan during the descent and after the touchdown. The data collected is not only valuable for Titan's science, but also for reference for future missions. The six Huygens instruments are briefly described below.

1. The Gas Chromatograph and Mass Spectrometer (GCMS) measured the composition of the atmosphère of Titan, isotopic ratios and trace species, from the parachute release time to the surface. It consisted of ion sources, mass analyzer and ion detectors. The GCMS sampling System also analyzed samples sampled by the Aérosol

Collector and Pyrolyzer (ACP) in an ACP-devoted feed tube. It has three parallel gas chromatograph columns, a quadrupole mass filter and five électron impact sources [START_REF] Niemann | The Gas Chromatograph Mass Spectrometer for the Huygens Probe[END_REF]. The mass spectrometer operated in the range of 2 to 141

Dalton. It began its operation after the protection shield release at the altitude of 160 km during the descent phase (Niemann et al., 2005).

2. The Aérosol Collector and Pyrolyzer (ACP) has sampled aérosols during the descent in two stages: (a) from the top of the atmosphère to 40 km and from 23 km to 17 km altitude. Except for the sampling System it hosts an oven where pyrolysis was performed and a transfer System of the samples. Then, it prepared the material for further analysis and it sent the samples to the GCMS for deriving their Chemical composition [START_REF] Israël | Huygens Probe Aérosol Collector Pyrolyser Experiment[END_REF]2005).

3. The Descent Imager/Spectral Radiometer (DISR) is the optical remote sensing instrument in the range of 0.3 to 1.7 microns [START_REF] Tomasko | The Descent Imager/Spectral Radiometer (DISR) Experiment on the Huygens Entry Probe of Titan[END_REF]. The DISR acquired spectra and high-resolution images of Titan's atmosphère while also measuring the solar radiation in the atmosphère. Except for the imagers, a visible spectrometer, an IR spectrometer, a solar aureole caméra, violet photometers, and a sun sensor are parts of the instrument. DISR also had a 20 W lamp (surface science lamp), which switched on during the last stages of the probe's descent (at 700 m altitude) in order to enlighten the surface beneath it. It also revealed traces of hydrocarbon liquids on its surface through complex drainage Systems and sent back to the Earth the first images of Titan's surface [START_REF] Tomasko | Rain, winds and haze during the Huygens probe's descent to Titan's surface[END_REF].

4. The Huygens Atmosphère Structure Instrument (HASI) consists of a 3-axis accelerometer, a set of a coarse and a fine température sensors, a multi-pressure sensor, a microphone and an electric field sensor array [START_REF] Fulchignoni | The Characterisation of Titan's Atmospheric Physical Properties by the Huygens Atmospheric Structure Instrument (Hasi)[END_REF].

HASI studied the atmospheric structure of Titan, gave a detailed température vertical profile and recorded the surface pressure and température (1,467 ± 1 mbar and 93.65 ±0.25 K, respectively). HASI température and pressure sensors probed the atmosphère directly for the first time from an altitude of 1400 km in altitude down to the surface. Indeed, the P, T sensors were deployed after the shield release and sampled the local atmosphère. These parameters help to calibrate measurements from other instruments both on the probe and from the orbiter. The derived HASI température vertical profile is used for instance as a reference when adopting inverse methods to retrieve the température and abundance from CIRS data. Huygens studied the winds and the turbulence of the atmosphère, as well as its conductivity and searched for lightning (Fulchignoni et al., 2005).

5. The Doppler Wind Experiment (DWE) used one of the two redundant chains (Transmitter A) of the probe-orbiter radio link and it is mainly based on both probe and orbiter [START_REF] Bird | The Huygens Doppler Wind Experiment -Titan Winds Derived from Probe Radio Frequency Measurements[END_REF]. The DWE data was transmitted to Cassini through Channel A, which due to a command error, was lost. Fortunately, the Earth-based large radio telescope network recorded the signal directly (Bird et al., 2005). DWE determined the direction and the magnitude of Titan's zonal winds and confirmed the super-rotation of Titan's atmosphère (Bird et al., 2005;Lebreton et al., 2009).

6. The Surface Science Package (SSP) is a nine sensors suite which determined the properties of the lower atmosphère, the surface at the landing site and the subsurface.

The sensors are accelerometers, internai and external, penetrometers, sonarvelocimeter and density, permittivity, refractive, thermal properties index sensors as well as a tiltmeter [START_REF] Zamecki | Huygens' Surface Science Package[END_REF]. They showed a smooth surface but not fiat, resembling wet clay, packed snow or sand [START_REF] Zamecki | A soft solid surface on Titan as revealed by the Huygens Surface Science Package[END_REF].

Mission OverView

The Cassini-Huygens mission, launched in October 1997 from Cape Canaveral, used the gravity fields of the Earth, Venus, the Sun and Jupiter, to finally enter the Satumian System seven years later on July 1, 2004. Cassini carried the Huygens probe released in December 25, 2004 from the mother spacecraft towards Titan. After a successful enter, descent and landing (EDL) procedure, the probe touched down on the surface of Titan in January 14, 2005. The major phases of this seven-year space joumey are depicted in Fig. 1.8. Since its arrivai at Satura, and after the Satura Orbit Insertion (SOI), Cassini perforais continuous révolutions around the giant planet and frequently encounters its moons (flybys).

Table 1.1 below lists the phases of the Cassini/Huygens mission from its nominal stage to the second extension [START_REF] Nixon | Infrared limb sounding with Cassini CIRS: Optimal viewing strategy using horizon nodes[END_REF].

Table 1.1-The Cassini/Huygens mission phases [START_REF] Nixon | Infrared limb sounding with Cassini CIRS: Optimal viewing strategy using horizon nodes[END_REF]. Technologically, the Cassini orbiter has executed more than one hundred close flyby manoeuvers of planetary bodies, while orbiting Satum, more than the flybys, which hâve ever been performed in the entire planetary program. The successful entry, descent and landing (EDL) of the Huygens probe on the surface of Titan stands for another technologically remarkable point.

Scientifically, the Cassini-Huygens mission has significantly advanced our knowledge about Satum, its moons, its rings and its magnetosphere. Titan, one of the prime targets of the mission, has been thoroughly studied remotely by the advanced Cassini orbiter's instmmentation, during numerous flybys, and in situ by the Huygens probe.

The Cassini magnetometer measurements [START_REF] Backes | Titan's Magnetic Field Signature During the First Cassini Encounter[END_REF][START_REF] Neubauer | Titan's near magnetotail from magnetic field and électron plasma observations and modeling: Cassini flybys TA, TB, and T3[END_REF] confirmed Voyager 1 observations [START_REF] Ness | The induced magnetosphere of Titan[END_REF] that Titan has no détectable large-scale global magnetic field and showed that Titan is strongly influenced by the Satumian dynamic magnetodisk [START_REF] Arridge | Warping of Satum's magnetospheric and magnetotail current sheets[END_REF][START_REF] Bertucci | The Magnetic Memory of Titan's Ionized Atmosphère[END_REF][START_REF] Dandouras | Titan's exosphere and its interaction with Satum's magnetosphere[END_REF][START_REF] Simon | Titan's highly dynamic magnetic environment: A systematic survey of Cassini magnetometer observations from flybys TA-T62[END_REF]. Indeed, Titan orbits at an average of 20.2 Rs (1 Rs = Satumian radius = 60,330 km), while the Satumian magnétopause standoff distance follows a bimodal distribution with means at 21 Rs and 27 Rs, mainly controlled by the solar wind [START_REF] Achilleos | Large-scale dynamics of Satum's magnétopause: Observations by Cassini[END_REF]. This means that Titan is mostly located inside the Satumian magnetosphere, close to the Satumian magnétopause. Only during the T32 and T42 Cassini flybys, Titan was spotted in the shocked solar wind, inside the Kronian magnetosheath [START_REF] Bertucci | The Magnetic Memory of Titan's Ionized Atmosphère[END_REF][START_REF] Garnier | Titan's ionosphère in the magnetosheath: Cassini RPWS results during the T32 flyby[END_REF][START_REF] Rymer | Discrète classification and électron energy spectra of Titan's varied magnetospheric environment[END_REF][START_REF] Wei | Unusually strong magnetic fields in Titan's ionosphère: T42 case study[END_REF]. The Charge-Energy-Mass-Spectrometer (CHEMS/MIMI) measurements showed that although Titan can be considered as a source of the Satumian magnetospheric plasma, the relative low concentration of N-1" and N2+ in the outer magnetosphere suggests low nitrogen escape ratio of Titan [START_REF] Krimigis | Dynamics of Satum's Magnetosphere from MIMI During Cassini's Orbital Insertion[END_REF], contrary to previous believes [START_REF] Eviatar | Titan's gas and plasma torus[END_REF][START_REF] Richardson | Thermal plasma and neutral gas in Satum's magnetosphere[END_REF].

Titan's atmosphère interacts directly with the corotating Satumian magnetoplasma, generating a flow-induced magnetosphere. This interaction consists of coupling between charged particles escaping Titan and the extemal plasma. Therefore, the plasma variations with respect to the local time together with the changing solar wind conditions affect significantly the local plasma environment [START_REF] Wolf | Titan's Highly Variable Plasma Environment[END_REF]. Indeed, Cassini data showed horizontal variations, observed in the upper atmosphère, originating from the magnetic plasma variability, when Titan passes through different magnetospheric régions [START_REF] Bell | Simulating the time-dependent response of Titan's upper atmosphère to periods of magnetospheric forcing[END_REF][START_REF] Westlake | Titan's thermospheric response to various plasma environments[END_REF].

Cassini findings hâve put constraints on our understanding of Titan's atmospheric escape mechanisms and upgraded our knowledge of its exosphere and its interaction with the Satumian magnetosphere [START_REF] Garnier | The exosphere of Titan and its interaction with the kronian magnetosphere: MIMI observations and modeling[END_REF][START_REF] Dandouras | Titan's exosphere and its interaction with Satum's magnetosphere[END_REF]. The exosphere is the outermost région of a planetary atmosphère, where collisions between atmospheric particles cease to be important and the particles can escape and drift into the space [START_REF] Chamberlain | Planetary coronae and atmospheric évaporation[END_REF].

Since Titan lacks a protective intrinsic magnetic field, energetic particles originating from the Satumian magnetosphere directly bombard its upper neutral atmosphère, undergoing charge-exchange collisions with cold neutral atoms and producing energetic neutral atoms measurements [START_REF] Garnier | The exosphere of Titan and its interaction with the kronian magnetosphere: MIMI observations and modeling[END_REF], confirming pre-Cassini prédictions [START_REF] Dandouras | Production and imaging of energetic neutral atoms from Titan's exosphere: a 3-D model[END_REF]. The exospheric molécules included in this thermal model are CH4, N2, N, H2 and H, while the exobase, the lower boundary above the thermosphère, is located at 1425 km with a température of 149K [START_REF] Waite | Ion Neutral Mass Spectrometer Results from the First Flyby of Titan[END_REF][START_REF] De La Haye | Cassini Ion and Neutral Mass Spectrometer data in Titan's upper atmosphère and exosphere: Observation of a suprathermal corona[END_REF]. However, Titan's upper atmosphère is not in thermal equilibrium and INMS data led to the model of a non-thermal exosphere, using kappa distribution fonctions to fit the species distributions [START_REF] De La Haye | Cassini Ion and Neutral Mass Spectrometer data in Titan's upper atmosphère and exosphere: Observation of a suprathermal corona[END_REF]. ENA images variability is either related to the magnetospheric variability or to their distance from Titan. Strong asymmetries hâve been observed to the stable ENA halo around Titan mainly caused by the finite gyroradii effects of the parent ions [START_REF] Garnier | Statistical analysis of the energetic ion and ENA data for the Titan environment[END_REF]. INCA also detected Titan's exosphere outer extent at an altitude of 50,000 km [START_REF] Brandt | The distribution of Titan's high-altitude (out to 50,000 km) exosphere from energetic neutral atom (ENA) measurements by Cassini/INCA[END_REF] close to the Hill sphere radius, the limit of the gravitational influence.

Between the altitudes of 1500 km, close to the exobase, and 1000 km, ENAs can suffer multiple collisions, but they can still be detected. Below 1000 km, ENAs can be considered as absorbed [START_REF] Dandouras | Titan's exosphere and its interaction with Satum's magnetosphere[END_REF]. At the same altitude, the keV to tens of keV 0+ ions and energetic H+ ions (E > 50 keV) from the Satumian magnetosphere can penetrate across field lines and deposit their energy into Titan's atmosphère [START_REF] Cravens | Energetic ion précipitation at Titan[END_REF] resulting in heating and escape of suprathermal nitrogen and methane atoms and molécules from Titan's upper atmosphère [START_REF] Michael | Energy déposition of pickup ions and heating of Titan's atmosphère[END_REF].

The Cassini/INMS detected very heavy negatively charged particles (most likely aérosols) above the homopause level (800-850 km) [START_REF] Coates | Discovery of heavy négative ions in Titan's ionosphère[END_REF]Waite et al., 2007).

Hence, the process of aérosol formation appears to start at more than 1000 km above the surface through complex ion and neutral Chemical reactions in the atmosphère. These reactions are initiated by energetic photons and particles from the Satumian magnetosphere.

Extrapolation of Cassini INMS (with a mass limit up to 100 daltons) and Plasma Spectrometer (CAPS) data, suggests that up to several thousands daltons, high-molecularweight species may exist in the ionosphère, including polymers of high molecular weightup to and certainly beyond -C7 hydrocarbons. These results show that ionospheric chemistry plays an important rôle in the formation of complex hydrocarbons in Titan's environment (Waite et al., 2007). Fig. 1.9 illustrâtes the pattern of tholin formation in the upper atmosphère as resulted from INMS [START_REF] Waite | Cassini Ion and Neutral Mass Spectrometer: Enceladus Plume Composition and Structure[END_REF]. The INMS data analysis showed that Titan possesses the most compositionally complex ionosphère in the Solar System. Modeling of INMS data (Vuitton et al., 2007) detected polyynes (C4H2, C6H2, CgtU) and cyanopolyynes (HC3N, HC5N) with densities higher than photochemical models' prédictions (Yung et al., 1984;Toublanc et al., 1995; Wilson & Atreya, 2004). Moreover, the model also probably detected methylcyanopolyynes (CH3C3N, CH3C5N) and methylpolyynes (CH3C4H, CH3C6H) for the first time on Titan.

Finally, the model indicates the existence of ammonia (NH3), methanimine (CH2NH), other nitriles (C2H3CN, C2H5CN) and two unidentified N containing species (C5H5N, C6H7N).

INMS also detected the neutral mode of benzene (CôHô) in the ionosphère (Waite et al., 2007;[START_REF] Vuitton | Formation and distribution of benzene on Titan[END_REF]. These hydrocarbons and nitrile species form a complex layering of organic aérosols (tholins) that diffuse through the atmosphère and accumulate on the surface (Coustenis et al., 2007;Tomasko et al., 2008b) (Fig 1.9).

Titan atmospheric région from 500 to 950 km was observed by the Cassini Ultraviolet

Imaging Spectrometer (UVIS), which monitored the occultation of two stars by Titan during the second Titan flyby [START_REF] Shemansky | The Cassini UVIS Stellar Probe of the Titan Atmosphère[END_REF]. A mesopause was inferred at 615 km with a température minimum of 114 K. Methane, acetylene, ethylene, ethane, diacetylene, and hydrogen cyanide were identified at altitude ranges from 450 to 1600 km. The higher order hydrocarbons and hydrogen cyanide peak sharply in abundance and are undetectable below altitudes ranging from 600 to 750 km, leaving methane as the only identifiable carbonaceous molécule below 600 km in this experiment. Cassini instrumentation cannot probe at this altitude and this région is called agnostosphere [START_REF] Coustenis | Atmospheric/Exospheric Characteristics of Icy Satellites[END_REF] or ignorosphere [START_REF] Strobel | Atmospheric Structure and Composition[END_REF].

The Cassini Composite Infrared Spectrometer (CIRS) instrument provides a complété pole-to-pole global coverage of Titan's stratosphère since the beginning of the mission. It can operate in nadir and limb viewing modes and provide the most spatially resolved maps of constituent species to date. Nadir sequences provide global mapping of atmospheric température and composition, while the limb ones provide information on the vertical profiles. CIRS has detected organics, probing the atmosphère within the range of 150 and 450 km (Flasar et al., 2005;Coustenis et al., 2007;2010a;Vinatier et al., 2007a;2010b;Teanby et al., 2008;2009a) and confirmed the results of Vl/IRIS and ISO. CIRS has also confirmed the benzene détection at 674 cm'1 (Coustenis et al. 2007) by ISO observations (Coustenis et al. 2003). The retrieved mixing ratios show no longitudinal variation. Températures are mostly zonally symmetric with almost no variation with the longitude as well (Achterberg et al., 2008).

In general, ail the nitriles and several hydrocarbons such as C3H4, C4H2, CôHô and C2H6 show enhanced concentrations at northem latitudes in spring, implying the existence of a global-scale Handley cell. The trace gases abundance at 55°S has decreased during the mission, which is consistent with upwelling in the South (Teanby et al., 2009b). Cassini/CIRS also indicate that very strong zonal winds occur in the stratosphère of Titan in winter northem latitudes, while in the summer the stratospheric winds are weaker (Achterberg et al., 2008),

supporting the presence of a North Polar vortex (Flasar et al. 2005). When the Cassini mission has entered to its solstice phase (Table 1.1) Titan fulfïlled a complété révolution of the Sun from the VI encounter (in mid-2010) and seasonal studies can be conducted (Teanby et al., 2010).

The organics from Titan's atmosphère form the haze particles located at 80-100 km and finally, they accumulate on the surface. Cassini ISS images and RADAR echoes hâve recorded several earth-like features on Titan's surface such as mountains (Radebaugh et al., 2007;Lopes et al., 2010), ridges [START_REF] Soderblom | Topography and geomorphology of the Huygens landing site on Titan[END_REF], faults [START_REF] Radebaugh | Régional geomorphology and history of Titan's Xanadu province[END_REF], rectangular drainage patterns and cryovolcanic structures, which are controlled most likely, at least partially, by tectonic activity (Burr et al., 2009). Additionally, stable liquid lakes were spotted by both ISS and RADAR, located mainly at the polar régions (Fig. 1.10) (Mitri et al., 2007;Stofan et al., 2007;Hayes et al., 2008) and recently at the équatorial latitudes [START_REF] Griffith | Possible tropical lakes on Titan from observations of dark terrain[END_REF]. Seasonal changes hâve been recorded during the Cassini flybys when the shorelines of the Southern large lake Ontario retreated during the Cassini mission time period [START_REF] Turtle | Rapid and Extensive Surface Changes Near Titan's Equator: Evidence of April Showers[END_REF]). Titan's organic inventory has been estimated to exceed the terrestrial one by a factor of magnitude (Lorenz et al., 2008b).

Remote sensing data from Cassini indicate that Titan's interior is partially differentiated. Indeed, the variations of its degree 2 coefficient gravitational potential provided by the Cassini Radio Science Subsystem (RS S) support internai density variations [START_REF] Rappaport | Can Cassini detect a subsurface océan in Titan from gravity measurements[END_REF]. Additionally, the Cassini Synthetic Aperture RADAR data suggest non-synchronous spin rate (Lorenz et al., 2008c), and the Permittivity, Waves and Altimetry (PWA) sensor on the Huygens Atmosphère Structure Instrument (HASI) with the Cassini Radio and Plasma Wave Science (RPWS) recorded Schumann résonance [START_REF] Beghin | New insights on Titan's plasma-driven Schumann résonance inferred from Huygens and Cassini data[END_REF], giving evidence for internai océan. km beneath the surface. If the interior was not déformable (Fig. 1.11), such solid tides would be only at 1 m in height [START_REF] Iess | The Tides of Titan[END_REF].

The Huygens probe was released from the Cassini orbiter on 25 December 2004. On 14 January 2005, the probe reached Titan's upper atmosphère and after a descent procedure, which lasted for almost 2.5 hours, it landed. Cassini collected data from the Huygens probe for lh and 12 min after the landing before disappearing below the horizon [START_REF] Lebreton | An overview of the descent and landing of the Huygens probe on Titan[END_REF]. The probe was still in operation and its signal was monitored by a global network of 17 ground-based radio télescopes [START_REF] Witasse | OverView of the coordinated ground-based observations of Titan during the Huygens mission[END_REF]. For more than three hours after the touchdown, the probe operated normally, continuing its surface experiments (Lebreton et al., 2009).

The orbiter and the probe had two radio link channels, A and B, the transmitters of which were installed on board the probe, while the receivers on board the orbiter. Both transmitters and receivers were equipped with a température controlled crystal oscillators (TCXO) for providing sufficient frequency stability for telemetry. Channel A was additionally equipped with ultra-stable oscillators (USO) necessary for the Doppler Wind Experiment (DWE). The performance of USOs was satisfactory during the mission's joumey to the Satumian System and it was then decided to be used instead of TCXOs. However, the command to power the USO of Channel A on the receiver (aboard the orbiter) was omitted and the receiver did not hâve a reference oscillator to lock on the signal of the probe. DWE frequency measurements together with the non-redundant telemetry data from Channel A were lost. The signal of Channel A from the probe's USO was received from 15 Earth-based radio télescopes [START_REF] Lebreton | An overview of the descent and landing of the Huygens probe on Titan[END_REF].

The Huygens/HASI provided a continuous density, température and pressure vertical profile from the level of the exobase at 1380 km to the surface (Fulchignoni et al., 2005).

Probe's trajectory remained stable at the same latitude of 10.3°S during its descent and therefore the derived profiles can be considered as représentatives of the vertical structure of the atmosphère and as global average. From the entry level down to the heat shield release at 155 km, the probe measured the density indirectly from the accélération data and then, by assuming hydrostatic equilibrium, the relative pressure [START_REF] Harri | Vertical pressure profile of Titanobservations of the PPI/HASI instrument[END_REF]. The corresponding température is derived by assuming that the atmosphère is an idéal gas, which is valid above 40 km. Below the altitude of 155 km, HASI measured directly the pressure and the température. The HASI profile is illustrated in Fig. 1.12 below. Température (K)

Figure 1.12: HASI température profile (solid line). The dashed line shows the engineering model [START_REF] Yelle | The Yelle Titan Atmosphère Engineering Models[END_REF]. The mesopause is located at 490 km with a température of 152 K, the stratopause at 250 km (186 K) and the tropopause at 44 km (70.43 K) (adapted from Fulchignioni et al., 2005).

Inversion layers, gravity waves, gravitational tides or large wind shear, which créâtes Kelvin-Helmholtz instability, may hâve caused the recorded température variations in the upper atmosphère (Fig. 1.12). In this région HASI provided températures higher compared to the one of the engineering model. The température oscillations above the stratopause indicates that the atmosphère is not dominated by radiative processes and is strongly influenced by the wave activity (Fulchignoni et al., 2005). However, HASI and CIRS measurements do not agréé with the altitude of the stratopause. CIRS data set it higher at 312 km (183 K) [START_REF] Vinatier | The Titan l4N/15N and 12C/13C isotopic ratios in HCN from Cassini/CIRS[END_REF], while at this altitude HASI température is 185 K. Additionally, the calculation of the infrared radiance of HASI is inconsistent in relation to the observed CIRS radiance at methane 7.7 pm band (Lebreton et al., 2009). Finally, HASI measured the surface pressure and température at 1,467 ± 1 mbar and 93.65 ± 0.25 K, respectively (Fulchignoni et al., 2005).

The HASI's Permittivity Wave and Altimetry (PWA) sensors measured the electrical properties of Titan's atmosphère. PWA revealed a small conductivity and électron density from 80 to 140 km and a high-density électron layer with a peak at 63 km [START_REF] Grard | Electric properties and related physical characteristics of the atmosphère and surface of Titan[END_REF][START_REF] Hamelin | Electron conductivity and density profiles derived from the mutual impédance probe measurements performed during the descent of Huygens through the atmosphère of Titan[END_REF]. As discussed above, the recorded Schumann résonance by PWA supported the existence of internai liquid océan [START_REF] Beghin | New insights on Titan's plasma-driven Schumann résonance inferred from Huygens and Cassini data[END_REF].

The Huygens/DWE was planned to measure the wind profile during the probe's descent, but due to the problem with Channel A in the communication link between the probe and the orbiter, DWE recordings were lost [START_REF] Lebreton | An overview of the descent and landing of the Huygens probe on Titan[END_REF]. However, the weak signal of Huygens was monitored by several large radio télescopes, which measured the signal frequency during the whole descent and saved the experiment. The detected signal allowed an estimate of its frequency and yielded data for the probe's motion during the descent phase from the associated Doppler shift. With these measurements, the wind profile of Titan was retrieved. For the lowest part of the atmosphère, from 5 km and down to the surface, the winds are very weak with velocities less than 1 m/s. Above 10 km, winds are prograde (eastward, westerly) and above 32 km super-rotation occurs with a velocity of 11.74 m/s.

From 60 to 100 km, the wind velocity reached almost zéro. DWE confirmed the super rotation of Titan's atmosphère (Bird et al., 2005).

The Huygens/ACP provided the first in situ measurements of Titan's aérosols (Israël et al., 2005). The collected samples show homogenous composition and the presence of N2, CH4, NH3 and HCN after the pyrolysis procedure. It seems that the aérosols are made of refractory organic nucléus, covered with condensed volatile compounds and they release NH3

and HCN during the pyrolysis. The latter assumption needs to be confirmed by laboratory work. However, these preliminary findings support the tholin hypothesis (Sagan & Khare, 1979).

The Huygens/DISR took hundreds of images of the surface during the descent, which hâve revealed a terrestrial-like surface, with drainage networks and rounded pebbles, made probably by water ice. DISR also measured the haze density particles during the descent phase and indicated a thin haze layer at 21 km. Additionally, DIRS recorded the continuous réflectance spectra of the surface using a lamp for a few hundred meters before the touchdown. Water ice was identified, but no tholins [START_REF] Tomasko | Rain, winds and haze during the Huygens probe's descent to Titan's surface[END_REF].

About the composition of the atmosphère, Huygens/GCMS data analysis provided a mixing ratio of methane at 1.48%±0.09 and hydrogen at 1.010.16 x 10'3 (Niemann et al., 2010). The CIRS value of methane at 0.016±0.005 (Flasar et al., 2005) is consistent with the GCMS one. GCMS measured a 40% increase of the methane concentration after the landing, perhaps due to thermal conductivity (Lorenz et al., 2006a). GCMS also identified organics in the surface such as ethane and cyanogen, carbon dioxide and benzene (Niemann et al., 2005).

The Huygens/SSP detected a relatively smooth, but not fiat surface [START_REF] Zamecki | A soft solid surface on Titan as revealed by the Huygens Surface Science Package[END_REF]. According to SSP results, the fracture of the landing surface is like the terrestrial wet sand, wet clay or packed snow.

What précédés is a short summary of a great wealth of data that the Cassini-Huygens mission provided, considerably enhancing our knowledge of Titan. In the rest of the manuscript, I will corne back to some of these results in more detail as they are connected to my research work.

Although Titan's atmosphère is much colder than the Earth's, it présents many direct similarities with our planet at different levels that hâve been pointed out since the Voyager days. Both atmosphères are dominated by the same main constituent, dinitrogen, albeit Titan's atmosphère extends much higher than the Earth's. A similar vertical structure from the troposphère to the ionosphère is also présent, as well as a surface pressure of just 40% larger than on Earth (Fulchignoni et al, 2005). This is the only case of an extraterrestrial planetary atmospheric pressure with several similarities to that of the early atmosphère of the Earth.

Furthermore, a very exciting and complex organic chemistry takes place in Titan's atmosphère. Table 1.2 summarizes the basic parameters for Titan.

Table 1.2-Titan basic parameters (Source: NASA, ESA4 [START_REF] Zebker | Size and Shape of Satum's Moon Titan[END_REF] how the data processing was improved, when a new calibration algorithm was applied to them. The fundamentals of the theory of radiative transfer and the atmospheric model of Titan that I use for interpreting CIRS data are defined in Chapter 3. I also describe the upgrade of the radiative transfer code I hâve performed by adapting recent spectroscopic datasets, implementing new haze model and advancing some programmatic aspects from the previous versions.

The results of my research are presented in Chapters 4 and 5. Chapter 4 describes the méridional and temporal variations in the température and composition of Titan's stratosphère, as retrieved by CIRS data. The abundances' variations of the major trace gases of Titan's stratosphère are presented since the beginning of the mission up to early 2012, focusing on 5 latitudinal bins (50°N, 30°N, 0°, 30°S, 50°S). Constraints for photochemistry and dynamics are also given. Additionally, I présent my contribution to water vapor retrievals from CIRS/FP1 spectra and the comparison with other radiative transfer codes' outcomes.

Furthermore, I présent the obtained vertical distributions of benzene from the surface of Titan to the upper atmosphère by applying the 1-D photochemical model of Lawas et al. (2008a;[START_REF] Raulin | Planetary science: Organic lakes on Titan[END_REF]. By subtracting the observations and the model, I also show spectral régions of unidentified émissions and propose new molécules for future work.

In Chapter 5, I study the connection between Titan's lakes and its atmosphère and the lakes' contribution to Titan's methane cycle. I focus on Mayda Insula in Kraken Mare, the biggest lake of Titan, and study its surroundings using geomorphological techniques. I also describe the application of a new despeckle filtering technique for obtaining restored RADAR images.

In Chapter 6, I discuss the Astrobiological potential of Titan and other icy moons of the Solar System, since the recent results from the Cassini-Huygens missions show that the habitability zone should be extended. However, although the Cassini data helps us change our perspective of the Satumian System, several questions still remain unanswered and only a new mission will certainly shed light on them. In Chapter 7,1 describe my contribution to the design of experiments to be included in such missions. I propose two experiments to be included as payload in future missions to icy moons and especially Titan. One is focused on studying the interior of a lake on Titan by using micro-electro-mechanical devices (MEMS), while the other studies a seismograph. Finally, Cassini-Huygens results hâve raised the public interest for Satum and its moon. I report the outreach activities I hâve organized/participated of bringing Cassini-Huygens mission's accomplishments doser to the layman public in Chapter 8.

Chapter 2

Observations and CIRS data

2.1

The Cassini Composite Infrared Spectrometer (CIRS)

CIRS héritage and purpose

The Composite InfraRed Spectrometer (CIRS) (Flasar et al., 2004) on board the Cassini orbiter is an improved version of Voyager Infrared Interferometric Spectrometer (IRIS) [START_REF] Hanel | Infrared spectrometer for Voyager[END_REF]. CIRS high resolution (0.5 cm'1) measurements are ten times higher than the IRIS ones, with an extended spectral range from 7 to 1000 pm (10-1400 cm'1). The IRIS 4.3 mrad field of view (FOV) has been decreased to 0.27 mrad for the 600 to 1400 cm'1 région [START_REF] Kunde | Cassini infrared fourier spectroscopic investigation[END_REF].

CIRS probes both Satum and Titan's atmosphères from deep in their troposphère to high in their mesosphere. It mainly consists of two interferometers, which share the same scanning mechanism, one in the far infrared and the other in the mid infrared. By taking advantage of the dual interferometer concept, scientists achieve the maximum advantage of optics and detectors. For this reason, Cassini's CIRS was called "Composite".

CIRS sounds both in nadir and limb viewing mode. Due to its advanced spectral resolution compared to IRIS, it is able to identify new trace constituées and its isotopologues, the spectral signatures of which are very weak to be identified by previously used instrumentation. The CIRS high spectral resolution permits also to separate some of the spectral bands, which were blended in IRIS spectra, for species such as C3H4 and C4H2 at 630 cm'1.

CIRS is one of the four instruments mounted on the Cassini's Optical Remote Sensing Palette (RSP, Fig. 1.6) along with the Imaging Science Subsystem (ISS), the Ultraviolet Imaging Spectrograph (UVIS) and the Visual and Infrared Mapping Spectrometer (VIMS).

Description of the CIRS Instrument

The Composite InfraRed Spectrometer (CIRS) is a Fourier Transform Spectrometer (FTS) which senses the thermal radiation emitted from the objects in the Satumian System in the thermal range of 55 to 200K and from 10 to 1400 cm'1 with an apodized spectral resolution varying from 0.5 to 15.5 cm'1 (Flasar et al., 2004) (Fig 2.1). The conceptual layout of CIRS (Fig. 2.2) consists of a telescope, the relay optics, two interferometers, a reference interferometer, a Scan mechanism, the detectors, electronics assembly and a passive radiative cooler. The instrument's spécifications are listed in Table 2 Figure 2.1: CIRS mechanical configuration. Ail the parts of the instrument are depicted. The scale is printed below of the instrument in both inches and mm [START_REF] Kunde | Cassini infrared fourier spectroscopic investigation[END_REF]. Figure 2.2: The conceptual layout of CIRS [START_REF] Kunde | Cassini infrared fourier spectroscopic investigation[END_REF]. The single arrows show the light paths after been collected from the telescope, while the double arrows demonstrate the movable parts of the interferometers.

Table 2. 1 -Instrument spécifications [START_REF] Kunde | Cassini infrared fourier spectroscopic investigation[END_REF]Flasar et al., 2004).

Mass

kg

Telescope's primary mirror diameter X (mrad) and FP4 are the mid-IR focal planes (Flasar et al., 2004). The far-IR detector has larger FOV than the mid-IR ones.

CIRS interferometers hâve one focal plane array (FPA), the Focal Plane 1 (FP1), sensing the far-IR portion of the emitted radiation (16.7 to 1000 microns) and a pair of focal plane arrays, the Focal Planes 3 and 4 (FP3, FP4), sensing the mid-IR (7.1 to 16.7 microns).

In the primary proposai two focal planes the FP1 from 10 to 300 cm'1 and the FP2 ranging from 300 to 600 cm"1 covered the far-IR investigation. However, due to budget cuts, FP2 was not mounted but merged into one focal plane the FP1.

The FP1 is a polarizing interferometer, having a pair of thermopile detectors with a 3.9-mrad field of view (FOV) per pixel (Fig 2 .3). The mid-IR interferometer is a traditional

Michelson where the FP3 and FP4 hâve a photoconductive 10-element array of HgCdTe detectors and a photovoltaic 10-element array of HgCdTe detectors, respectively. Each pixel of the array has a 0.273-mrad FOV [START_REF] Kunde | Cassini infrared fourier spectroscopic investigation[END_REF]. The Table 2.2 below lists the spécifications of the focal planes. [START_REF] Kunde | Cassini infrared fourier spectroscopic investigation[END_REF]Flasar et al., 2004).

Focal Plane FP1 FP3 FP4

Spectral Range (cm'1) The CIRS telescope CIRS telescope is of Cassegrain type, made from béryllium with a very low scattering surface. This configuration combines two mirrors: a primary F/6 concave parabolic reflector at 50.8 cm and a secondary convex hyperbolic one at 7.6 cm. The gold-enhanced surface of the primary mirror provides low scattering. Both mirrors share one focus, while the other focal point of the hyperbolic mirror is located where the image is to be observed. A supporting tube, cylindrical in shape, extends from the primary mirror center to the secondary mirror.

The primary dish reflects the incoming IR light beams towards its one focus, while the secondary mirror converges these rays to the other focus. Then, the light beam falls on a field splitting mirror, the Entrance Aperture Plate (EAP) [START_REF] Kunde | Cassini infrared fourier spectroscopic investigation[END_REF]. The EAP divides the incoming light and directs it to the two interferometers via collimators and folding mirrors (see Fig. 2.2). Each collimator aligns the IR beam part to the fold mirror and then the light passes through a beamsplitter and onto retroreflectors [START_REF] Kunde | Cassini infrared fourier spectroscopic investigation[END_REF]Flasar et al., 2004).

The beamsplitters split the wave front of the incident ray in two separate beams, while the retroreflectors reflect the light back to its source with a minimum scattering. After being split, the recombined output beams go to the focal planes and are focused on the relative detectors. An electromechanical shutter is mounted in the mid-IR path, which can be used for internai calibration issues as it can be commanded to stop the incident beam. A solar blocking filter protects the FP1 thermopile detectors from any accidentai solar illumination, rejecting rays shorter than 16.7 microns. The optimum operation température of the instrument is at 170 K except for the mid-IR detectors, the température of which is at around 80 K.

By applying the Fourier transform formalization at the interferogram we can reconstruct the original spectrum. Larger sample numbers in the interferogram give higher resolution of the final product. In contrast with the grating spectrometers, the FTS has no loses of the incident light, which is called the multiplex advantage. To achieve higher resolution, the interferogram can be scanned more slowly and add a large number of interferograms [START_REF] Lewis | Physics and chemistry of the solar System[END_REF].

The far-IR interferometer (FP1)

For the far-IR interferometer, a polarizing Martin-Puplett type [START_REF] Martin | Polarised interferometric spectrometry for the millimétré and submillimetre spectrum[END_REF]) is mounted on CIRS in order to record the far-IR part of the incident light (Fig. 2.4). The linear polatization design hosts wire-grid polarizers to separate and recombine the incoming signal at FP1. CIRS wire-grid polarizers are nearly idéal for far-IR experiments. When the electromagnetic radiation meets the grid, the oscillating field parallel to its wires produces a wave originating from the électron movement along the wires, while the grid absorbs the rest of the light. Then, the radiation proceeds to the beamsplitter polarizer, located 45°in respect to the First polarizer, where it is split one more time into two orthogonal components. These beams fall in two roof top mirrors rotated at 90°and retum back to the beamsplitter.

CIRS FOCAL PLANE 1 ÀSSEMBLY

Figure 2.4: Cassini/CIRS FP1 assembly [START_REF] Kunde | Cassini infrared fourier spectroscopic investigation[END_REF]. The polarizing grid and the concentrators are shown as well as the focus mirror.

The beamsplitter recombines these components which now hâve a phase différence of 90°and eventually send the recombined beam, which has now elliptical polarization, to the focal plane. A polarizer/analyser (Fig. 2.4) is placed in the front of the focal plane's detectors and it converts the polarization modulation into intensity modulation. The polarizer, beamsplitter and analyser grids consist of photolithographic copper wires having each a diameter of 1 micron and a spacing of 2 microns, deposited on Mylar, a strong polyester film.

The modulated frequency is analogous to the wavenumber of the radiation, producing the résultant intensity of the combined beams, a conventional interferogram. The interferogram is a fonction of the path length différence between the two beams path which when they differ by intégral number of wavelengths they interfère constructively. In any other path différence, they interfère destructively and they do not appear in the observed interferogram.

The polarizer/analyser splits the polarized beam into the transmitted component and the reflected one and sends them to the thermopile pair detectors. A Compound Parabolic Concentrator (CPC) is installed in the front of each thermopile detector (Fig. 2.4). Each detector receives one of the split components and now the final signal is gained, with no unmodulated intensity fluctuations and provided redundancy.

The mid-IR interferometer (FP3 and FP4)

The mid-IR interferometer is a conventional Michelson design [START_REF] Michelson | On the Velocity of Light[END_REF], having KBr beamsplitter and cube-comer retroflectors. After the split of the incoming radiation into two beams of almost equal intensity, one of them proceeds to the stationary mirror, while the other directs to the moveable arm of the interferometer. After having been reflected and transmitted once, the beams are recombined at the beamsplitter that is also a compensator and travel to the detectors. The interferometers spécifications as well as the spécifications of the focal planes are showed in Table 2.3 below.

Table 2.3: Interferometers spécifications [START_REF] Kunde | Cassini infrared fourier spectroscopic investigation[END_REF]Flasar et al., 2004).

Spécifications

Mid-IR Far-IR Type Michelson Martin-Puplett Spectral Range (cm'1)

10-600 600-1400

Spectral Resolution (cm'1) 0.5-20 0.5-20 Intégration time5 (s) 5-52 5-52 The reference interferometer

The reference interferometer is the hardware component of CIRS, which provides an accurate calibration standard for the interferometer (Fig. 2.2). It is installed within the mid-IR interferometer of the instrument, uses a solid State 785 nm diode laser at 170 K providing a servo signal for velocity control. The laser is a source of monochromatic radiation and it is température depended. For redundancy the instrument is equipped with two laser diodes, selectable via command. The function of this interferometer is to generate a signal in order to sample the outputs from ail the focal plane detectors of the instrument. Sampling of the created interferograms occurs at zéro Crossing of equally spaced reference fringes. The same signal is additionally used for handling the velocity of the scan mechanism.

The term zéro path distance or zéro path différence (ZPD) describes the relative distance between the two moving mirrors of the instrument in order to hâve the two arms of the interferometer at equal optical path length. When this distance is zéro, ail the light rays interfère constructively, producing the white light spike in the interferogram. When the interférence is constructive for every wavenumber at the ZPD, the interferogram has the maximum relative intensity at one of the output ports and the minimum at the other. Thus the ZPD should be placed at every interferogram in order to ease the co-adding of interferograms and reduce the CIRS telemetry data rate. Such recording takes place due to a white light interferogram originating from the Light Emitting Diode (LED) mounted on the reference interferometer.

The scan mechanism

The scan mechanism is the same for the interferometers as well as for the reference interferometer and Controls the sampling sequence and the moving mirrors' adjustment (Fig.

2.2)

. It follows the pattern of Voyager/IRIS and hosts the carriage device and the motor 5 When using the Hamming apodization function. Unapodized widths hâve the corresponding scan times at about half as wide [START_REF] Nixon | Infrared limb sounding with Cassini CIRS: Optimal viewing strategy using horizon nodes[END_REF].

responsible for the moving interferometer mirrors (the moving roof top mirror of FP1, the moving retroreflector for FP3, FP4 and the reference interferometer).

The CIRS detectors

During its flybys, the Cassini orbiter approaches Titan close enough at an altitude of no less than 1000 km, so that the IR spectrometer can perform not only downward viewing, but also limb profiling. This type of sounding needs a small field of view and for this reason highefficiency photo-electronic detectors hâve been applied to the instrument [START_REF] Nixon | Remote sounding of the atmosphère of Titan[END_REF]. CIRS detectors can be distinguished to a pair of thermopiles on the far-IR part of the instrument and to mercury cadmium telluride (HgCdTe) photoconductive (PC) and photovoltaic (PV) detectors on the mid-IR part. The thermopile detectors operate at 170K, which is the CIRS functional température, while both the PC and PV operate at 80K. A passive radiative cooler is mounted at this part of the instrument, which reduces the initial température to 80 K.

In general, thermopile detectors are passive uncooled devices, which require no electrical bias. Such devices hâve negligible 1/f noise and they do not sense substrate température variations [START_REF] Foote | Uncooled thermopile infrared detector linear arrays with detectivity greater than 109 cmHzl/2/W[END_REF]. The thermal detectors of the far-IR sense the power of absorbed radiation. They receive one of the two components of the elliptical polarized beam from the polarizer/analyzer and improve the signal's gain by eliminating any unmodulated intensity fluctuations, providing redundancy. Each thermopile detector converts the incident infrared light into heat through an absorber, while a thermoelectric component créâtes an output voltage proportional to its internai température gradient.

The FP3 and FP4 detectors respond to the number of photons, which arrive per unit time. Both consist of HgCdTe linear arrays. A linear array of 1 x 10 photoconductive (PC) components is mounted in FP3, while a linear array of 1 x 10 photovoltaic (PV) components is mounted in FP4. Their back-to-back configuration is shown in Fig. 2.5. When photons hit the photoconductors, they generate free charge carriers and the résistance of the material is decreased. This results in a voltage change, which is the measured signal. On the other hand, when photons hit the photovoltaic detectors, électrons and holes are generated forming an electric field [START_REF] Hanel | Exploration of the Solar System by Infrared Remote Sensing[END_REF]. With this structure the instrument takes advantage of the low power dissipation which is less than 0.5 mW per element and a limited array detectivity by 1/f noise [START_REF] Martineau | HgCdTe detector technology and performance for the Composite Infrared Spectrometer (CIRS)ZCassini mission. Cassini/Huygens: A mission to the Saturnian Systems[END_REF].

i/2 in Figure 2.5: Configuration of photoconductive and photovoltaic arrays on FP3 and FP4 respectively (Martineau et al., 1996).

CIRS scientiflc objectives

CIRS measures the infrared émission from atmosphères, rings and surfaces, which are mainly used to map the température, the hazes and clouds and the Chemical structure of the probed atmosphères and surfaces (Coustenis & Taylor, 2008). CIRS objectives are the following (http://cirs.gsfc.nasa.gov/science.html):

For Titan:

• Thermal and compositional mapping of atmosphère The Cassini mission covers a large part of the Satumian year, which is circa 29.5 terrestrial ones. In mid-2010, a full Satumian year was completed since the Voyager 1 encounter. Moreover, ground and space borne observations hâve been published during the interval between the two in situ missions. Hence, CIRS infrared spectra, recovered during the Cassini mission period, détermine both the temporal and the spatial variations of the gas composition of Titan's atmosphère in a global scale: its température and its opacity sources such as the aérosol distribution, the haze and the clouds are investigated. In addition, CIRS has gathered information throughout the last 8 years (2004 to 2012) and will hopefully continue to operate until 2017 when its extended mission will end. Thus, a wealth of data is available which hâve already changed our perspective of the atmospheric composition, dynamics and évolution over time.

As discussed before, CIRS is a more advanced instrumentation comparing to its ancestor, the Voyager 1/IRIS, which first sounded in situ the atmosphère of Titan. Not only does CIRS exploit its higher spectral resolution, but it also maps an extended spectral régime, larger than the one IRIS did in longer intégration times. In fact, CIRS affords a programmable spectral resolution from 0.5 cm'1 to 20 cm'1 (Flasar et al., 2004). Its highly improved operability gives the opportunity to search for minor atmospheric species as well as for new isotopologues and calculate isotopic ratios. The déterminations of the Chemical composition of the atmosphère and the température are therefore done with much higher accuracy.

CIRS data analysis gives the opportunity to advance current modeling in atmospheric dynamics and general circulation and understand the complex photochemistry occurring on Titan's upper atmosphère. Moreover, the great success of its operation sets the requirements for future missions to the outer planets at high standards.

CIRS is not only dedicated to Titan's atmosphère but also to the other Satumian moons and of course to the study of Satum, its rings, their thermal properties, composition and surface. However, its achievements conceming these planetary bodies are beyond this study.

Without doubt, CIRS is one of the most important experiments of the Cassini-Huygens mission.

Types of CIRS observations

The Cassini orbiter executes flybys of Titan, while is touring Satum. Table 2.4 lists the Cassini flybys completed up to date (July 2012). The Composite Infrared Spectrometer on board Cassini opérâtes during the Titan flyby sequence beginning at about 40 hours before the closest approach. Cassini approaches Titan with a velocity relative to the satellite ranging from 5.5 to 6.1 km/s. Fig. 2.6 below demonstrates the observation sequence during each flyby and the type of data, which CIRS can obtain moving towards the Closest Approach (C/A). Ail the far infrared types of the CIRS observations are listed in Table 2.5, while the mid-IR recording types are listed in Table 2.6.

TO was the First Cassini flyby of Titan, which was performed just after the Satum Orbit Insertion (SOI) of the spacecraft. However, the distances were too high compared to the observations which followed; hence, TO data is excluded from our studies. A CIRS commanding error led to the corruption of most of the data obtained by TA flyby. Since Titan's atmosphère expériences température inversion, the derived thermal spectrum consists of the émission part and the absorption part, which both dépend on the température gradient of the probed atmospheric layers.

The FP3 pixel viewing modes are four: ODD, EVEN, CENTER and PAIRS. In the ODD mode only the 1, 3, 5, 7, 9 detectors are active, while in the EVEN mode the detectors 2, 4, 6, 8, 10 are only active. In the CENTER mode the 3,4,5,6,7 detectors are active, while in PAIRS mode the detectors 1&2, 3&4, 5&6, 7&8, 9&10 are combined and amplified.

As far as the FP4 pixel mode is concemed, we hâve the same modes: ODD, EVEN, CENTER and PAIRS. The active detectors are as in the FP3 modes, except that the detectors 4,5,6,7,8 are used in CENTER mode, to spatially match the detectors of FP3, which are numbered in reverse physical order. In the CO-ADD mode (COADD or NO-COADD) two successive scans are added together in the on-board electronics (buffers), to reduce the data rate by a factor of 2.

Normally, CIRS data is read out at a rate of 4 kbs and every 2 seconds 8 kbit datasets are transferred via the Bus Interface Unit (BIU) to the spacecraft Command Data Subsystem (CDS). CIRS has many ways of reducing the data rate, such as co-adding the spectra, dropping FP3 or FP4 readout (reduces by 45%), dropping FP3 and FP4 readout (reduces by 91%), or going into housekeeping-only mode.

The instrument electronics compress the data packets and transmit them to Earth by using the Deep Space Network (DSN). Then, the datasets are decompressed by software on the ground and go into a processing pipeline, which organizes the science data into binary tables creating the Vanilla database and interpolâtes housekeeping records. We will discuss more about the data calibration in the following section. The Vanilla software is a simple database query tool. However, the user can also access the binary data by using other programming retrieval tools.

2.3

Improving the CIRS data processing

Instrumental artifacts in CIRS data

The recordings of the Cassini Infrared Spectrometer onboard Cassini suffer from some quality problems having to do with quality, calibration, noise, etc. CIRS FP1, FP3 and FP4 interferograms are subject to the presence of single frequency (sine wave) ripples with varying period and amplitude and prominent sharp spikes or delta fonctions with periods of 0.125 s (8 Hz) and 2.0 s (0.5 Hz) respectively [START_REF] Nixon | Inflight Far-Infrared Performance of the CIRS Instrument on Cassini[END_REF]. The source of the sine wave ripples is still unknown [START_REF] Carlson | Removing Artifacts in the Calibration of Cassini CIRS Spectra ofSaturn and Titan[END_REF]. On the other hand, the source of the noise spikes is the electrical interférence from the CIRS electrical sub-systems and especially the Bus Interface Unit (BIU) [START_REF] Nixon | Inflight Far-Infrared Performance of the CIRS Instrument on Cassini[END_REF]. Figure 2.7 shows an example of such sine wave features in both FP3 and FP4 CIRS spectra. An example of prominent spikes is shown in Fig. zooms at the région of 580 to 660 cm'1 with a spike located at 574 cm'1, the center (b) focuses on the 700 to 800 cm'1 where the spike is at 764 cm'1 and the right panel (c) shows the 1000 -1100 cm'1 and the spike at 1085 cm'1.

Each Cassini science instrument is connected to the Cassini's data bus through BIUs, which allow for two-way communication between the instruments and spacecraft [START_REF] Henry | An Introduction to the Design of the Cassini Spacecraft[END_REF]. However, a bug exists when there is a heavy communication traffic, which affects the instruments' interférences. This bug has caused instrument resets and transmitted commands not to be received by the instrument [START_REF] Linick | Cassini DistributedInstrument Operations -What We've Learned Since Saturn Orbit Insertion[END_REF]. The CIRS BIU is queried 8 times per second by the Cassini Command Data Subsystem (CDS), on the spacecraft clock signal puises and it produces a small spike. Data is transferred from CIRS to the CDS every 2 seconds, which causes a larger spike and this enhances its intensity [START_REF] Nixon | Inflight Far-Infrared Performance of the CIRS Instrument on Cassini[END_REF]. The resuit in the spectrum is a sériés of spikes spaced by 12 cm"1 corresponding to the harmonies of 0.5 Hz. The 8 Hz spike at 191.38 cm"1 also appears in FP1 CIRS spectra as strong harmonies at 16 Hz (382.76 cm'1) and 24 Hz (574.14 cm'1). Unavoidably, both spikes impact atmospheric retrievals and can make interprétation problematic [START_REF] Carlson | Removing Artifacts in the Calibration of Cassini CIRS Spectra ofSaturn and Titan[END_REF].

Regular wavenumber calibration of the CIRS instrument

In-flight calibration of the absolute radiance expressed by the spectral responsivity has been performed on CIRS by using two separate techniques. The instrument is exposed to a blackbody source of accurately known température (shutter-closed) and to the deep space, which is assumed to be a perfect sink [START_REF] Hanel | Exploration of the Solar System by Infrared Remote Sensing[END_REF]Flasar et al., 2004;[START_REF] Nixon | Inflight Far-Infrared Performance of the CIRS Instrument on Cassini[END_REF]. The former calibration task results the local calibration database ("current" Pipeline),

while the latter the deep space calibration database (DS4000).

However, any new database édition may impact seriously the spectra. After having tested each correction algorithm, a new database inference is released to the CIRS data users with the new application. Hence, our purpose is not only to track any relative issues in the queries we make for our spectral analysis, but also to test each CIRS database using our radiative transfer code as well. In the next sections, I show the results of our testing procedure, how the new database inferences advance our analysis and measure the impact of the new release to radiative transfer calculation fits. I firstly retrieved spectra from The v3.2 current pipeline is the évolution of the v2.5 current pipeline, incorporating further refmement in the calibration algorithm. A spectral query with négative radiances is one of the occasional problems with the official CIRS current Pipeline. This problem has been experienced around from the beginning of the mission. The CIRS calibration team is working towards solving these issues and applies correction algorithms in both calibration sources in order to reduce spikes' effects on CIRS interferograms as well as vanish the presence of négative radiances.

The DS4000 is an independent database of CIRS data, which contains ail consistent current pipeline data, but calibrated with using 4000 deep space spectra taken under the same conditions as the data in as close proximity to the data as possible. The introduction of such a data processing aims to résolve négative radiances and reduce the instrument's noise.

The Grand Average version of CIRS data has been developed in order to vanish négative radiances. This database is calibrated using a unique average for the entire mission.

In principle, the comparison to the local (current) database will provide a much better signalto-noise ratio and give good relative amplitude, but with baseline bias. The DS4000 provide spectra with also better signal-to-noise ratio without bias.

Another issue that should be addressed is whether an eventual systematic drift of calibrations as fonction of time As far as the detector's response is concemed, a local calibration in time takes place.

That means that the calibration team uses observations of space and 170 K shutter taken around the same time as the science observations, to estimate the response. They continuously update the calibration during the mission, so that if the detector response ever dégradés, it will automatically be accounted for. For instance, if 1000 detector counts correspond to 1 mW of radiance in 2007, but 900 counts correspond to 1 mW in 2012, their calibration will always compute the right radiance level, by looking at the "known" targets (space and shutter).

Retrieval of CIRS spectra and averaging procedure

In order to evaluate the impact of each new CIRS database that became available after corrections by the CIRS team at Goddard Space Flight Center on our atmospheric retrievals, and at the same time validate the improvement in the new database as we were asked to do, I hâve performed some tests using our radiative transfer code and running simulations using our updated atmospheric model for a given sélection in the available developmental versions to détermine the improvement and uncover any remaining calibration and/or other issues. I then compared the results of this fit to the ones retrieved using the spectra selected from previous official éditions.

Because the purpose of this study is to obtain high précision inferences of the Chemical and thermal composition, we always query for spectral averages which include a large number of spectra with given restricting parameters. We thus increase the signal-to-noise ratio. We mainly focus on nadir spectral observations and therefore we restrict the émission angles to be less than 65°, in order to hâve reasonable airmasses in our calculations (airmass is the inverse cosine of the émission angle). These queries restricted the spacecraft range from 100 to 400,000 km and accepted ail longitudes.

The radiative transfer code is applied to match the CIRS spectra. The fitting procedure is a simulation of the observations and comparison with the selected real spectra. We search for the best agreement between data and model through an itérative process. Then, we extract the mixing ratios of the atmospheric compounds. The comparison of the fits retrieved from each database shows if each new developmental database provides improved, worse or the same data.

Evaluating new CIRS databases and improving the data processing

Hereafter, I show three different validation tasks, we were asked to perform in order to evaluate the improvement in the CIRS data processing in each new database inference, and to validate the new calibration database which would then be implemented on the public access machines.

On 1 October 2010, we were asked from Ron Carlson, Research Associate Professor at the Catholic University of America, to validate the spike suppression algorithm has been developed with Ever Guandique, from ADNET Systems Inc. This multi-parameter CIRS spike suppression algorithm is expected to reduce or vanish négative radiances from CIRS spectra, [START_REF] Carlson | Removing Artifacts in the Calibration of Cassini CIRS Spectra ofSaturn and Titan[END_REF]. This algorithm employs deep space and shutter-closed interferograms for calibration averaged over several years. It matches the positions and intensities of the spikes to the observed spikes in each interferogram and performs checks to optimize and characterize the quality of the fitting process. The main idea is to calibrating the coadded radiances is the query of a FP3 large average from March 01, 2009 to November 30, 20097 at high resolution (0.5 cm"1) within a latitudinal bin at . In this sélection, the presence of négative radiances is obvious. The resulted averaged spectrum is hereafter called Grand Average and is also shown in Fig. 2.9 (in blue).

The average spectrum of the current database version v2.5 has a continuum problem from 570 to 700 cm'1 with négative radiance we want to correct. One of the sequences of this query therein commanded co-added Titan interferograms to be recorded, along with a few coadded deep space and shutter-closed interferograms for calibration (Ever Guandique, pers.

comm.). The négative radiances occur only for the co-added spectra in this sequence. cm"1). The test-averaged spectrum (Grand Average) for the same conditions is plotted in blue (Ron Carlson, Ever Guandique pers. comm.).

Quite a few of the customary no-add Titan interferograms were recorded at a higher CIRS instrument température than the usual 76.3K, presumably because of the solar angle during the observations. These higher température Titan interferograms hâve to be calibrated separately using higher température deep-space and shutter-closed interferograms and then added to the 76.3K Titan calibrated spectra using a weighted average. To achieve a good signal-to-noise ratio for the higher température Titan spectrum, we may hâve to go backwards 57 in time to earlier CIRS reference laser mode 1 time periods to average enough higher température deep-space and shutter-closed interferograms. The algorithm used about 12,000 deep-space and 1,200 shutter-closed noadd 76.3 K interferograms for each of the ten FP3 detectors (11 -20) in order to obtain a superior signal-to-noise rate for the calibrated Titan spectra. This calibration should yield the correct continuum. The resulted averaged spectrum is shown in Fig. 2.10 (in red) below.

One problem is that the 8 Hz spikes are very prominent in the 25°S spectrum if we use about 10,000 deep-space interferograms for each detector (11 -20) On 6 October 2010, Marcia E. Segura, CIRS Operations Team leader in Goddard, asked us to validate the DS4000 database using our radiative transfer code. DS4000 spectra hâve been treated as in the current pipeline, but calibrated with 4000 deep space spectra. In this approach, typically issues such as négative radiances are resolved and the noise is reduced. The current v2.5 (and lately v3.2) is the primary source for CIRS data users. The version 3.1 incorporâtes further refinement in the calibration algorithm. For interferograms collected using CIRS Flight Software Version 6, non-co-added data will hâve the 1/2 Hz noise spikes suppressed, while no noise suppression will be done on co-added data.

The First effort towards a better calibration of CIRS data was the "custom calibration"

which was applied to support the Flight Software Version 6 (FSW6) by selecting 1,000 deep space spectra. The spectra were selected in order to match the conditions of the data to be calibrated and be recorded before the starting spacecraft event time (SCET). If such 1,000 spectra could not be found, then the sélection procedure moves to the 1,000 deep space spectra, which match the data conditions after the ending SCET.

The Global calibration DS4000 takes into account 4,000 deep space spectra, which match the conditions of the data to be calibrated. Contrary to the custom calibration, it uses the 4,000 deep space spectra from the beginning of the mission and moving through the deep space by always using the most recent 4,000, replacing the older blocks.

The task was to provide the results of a comparison between the current v2.5 version of CIRS DATA and the DS4000 and check if this database resolved the négative radiances without loosing any information. The 3.1 (and 3.2) version was not released in this date.

On 3 May 2012, Marcia E. Segura, asked for validating the new dataset, which was generated with a similar algorithm than the one used to generate the DS4000 database, but with a phase correction algorithm applied from Nicolas Gorius, science collaborator, member of the CIRS science team. This new algorithm is expected either to correct most of the négative radiance or too-positive radiance outriders or be identified easily but the user.

Outstanding outriders are being shifted back to level where they are expected to be.

My purpose is to provide to CIRS calibration team représentatives the requested validation for each database release.

A) Grand Average versus DS4000

Using our line-by-line radiative transfer code I inferred a model simulation that I compared to both sélections. When the best fits possible were found in the émission bands, I inferred the abundances of the major molécules. The fits are illustrated in Fig. 2.11 a, b and the mixing ratios of the major species in Table 2.7 below. For ail these spectra we hâve used the same température profile, which fits the v4 methane band in the associated FP4. .

, ^iuk

1 1 1 1 '
Liku.. jHfHk. Wavenumber (cm 1) In general, the derived abundances for the major molécules in the FP3 range hâve a différence on the order of 15-20%. The continuum is 1.5 to 2 times higher in Grand Average spectrum between 600-890 cm'1. Since the CIRSDATA v2.5 pipeline contains large calibration problems, the DS4000 is a large improvement with despiking. The Grand Average test database is still significantly different ffom DS4000 and yields higher abundances.

B) DS4000 versus v2.5

The characteristics of the sélections I hâve chosen to perform this task are listed in the Table 2.8 below. I compared two spectral sélections, one from the beginning of the mission and a recent one. According to the sélections criteria the nadir-averaged spectra from both December 2006 and January 2010 were queried at high resolution (0.5 cm'1) in FP3 and FP4. database as well as the v2.5 database at high resolution (0.5 cm'1) at 50°N. Ls represents the corresponding solar longitude of each query. The number of spectra, the émission angle of the averaged spectral sélection, the signalto-noise ratio as well as the airmass of the collection are also listed. 50°N DS4000 v2.5

CI RS

Query

Ls N u m b e roi s p e c t r a E m i s s i o n a n g l e S/N A i r m a s s N u m b e rol s p e c t r a E m i s s i o n a n g l e S/N A i r m a s s S / ND i f f e i D S 4 0 0 0 The v2.5 current pipeline présents an improvement by a factor of 30% on the S/N ratio in the FP3 query of January 2010. The DS4000 has higher S/N by a factor of 8% at the FP3 In order to depict the différences between the two database inferences conceming the despiking resuit, I hâve overplotted the FP3 (Figs. 2.12,2.13) and FP4 (Figs. 2.14,2.15) spectra from each database. Their différence within the spectral range in each query is drawn in the lower panel of Figs.

1 1 1 1 1- L - 1 • 1 1 1 1 r- 1 1 1 1 DS4000 -v2.5 J 1 i . i i i . i . i . i i i 1 1 . 1
Wavenumber (cm'1)

Figure 2.12: December 2006-FP3 average spectrum from the current database v2.5 (in red) and from the DS4000 (in blue). Their différence is plotted below (in black). Spikes of the DS4000 pipeline exist at 764, 956, 1033 and 1121 cm'1 wavenumber.

x 10

JANUARY 2010 FP3 at 0.5 cm'1 at 50°N o 3:

Wavenumber (cm 1)

Figure 2.13: January 2010-FP3 average spectrum from the current database v2.5 (in red) and from the DS4000 (in blue). Their différence is plotted below (in black). Spikes of the DS4000 pipeline exist at 651, 764, 795, 842, 956 and 1033 cm'1 wavenumber. Wavenumber (cm'1)

Figure 2.15: January 2010-FP4 average spectrum from the current database v2.5 (in red) and from the DS4000 (in blue). Their différence is plotted below (in black).

Both DS4000 FP3 averaged spectra show spikes at 764, 956 and 1033 cm1. In December 2006 query one more spike has been recorded 1171 cm'1, while in the January 2010 query spikes are also présent at 651, 795 and 842 cm'1. The différence in radiance of the major spikes at 764 and 956 cm'1 between DS4000 and v2.5 is listed in the Table 2.9 below.

Table 2.9: The major spikes in CIRS FP3 spectra. The percentage corresponds to the relative différence between the DS4000 spectra and the v2.5 around the wavenumbers of 764 and 956 cm'1.

Major spikes in spectra w/n (cm1) The spikes in DS4000 spectra are more prominent than in the v2. For the December 2006 query, the abundance of acetylene -C2H2 -(Tables 2.10 and2 .11) calculated from the v2.5 pipeline is significantly higher than the DS4000 one at about 26% and the HCN at about 10%. The rest molécules hâve similar mixing ratios retrieved from both database inferences, located within the error bars. and 13% within the spectral range of 575-590, 590-625 and 640-650 cm'1 respectively (Table 2.12). Instead, the continuum level in the January 2010 FP3 sélection (Fig. 2.21) shows no significant différences between the two inferences. The continuum level appears in the DS4000 database lower than for v2.5 and négative in some cases. The FP4 spectra of the 2006 December sélection (Fig. 2.22) shows significant différences between the two database inferences at about 50% from 1050 to 1120 and 1170 to 1200 (Table 2.13). Once again no significant différence is depicted in January 2010 FP4 data (Fig. 2.23). The v2.5 continuum is higher at about 50% than the DS4000 one. December 2006 -FP4 continuum level (cm1) higher values différence between DS4000 and v2.5 I applied our radiative transfer code in order to fit C2H2, C2H6, HCN signatures at the three spectral inferences. I use the corresponding température profile as retrieved by the V4 methane band at FP4 spectra (Achterberg et al., 2008;2011) to then infer the major molécules' abundances after a best-fit modeling procedure.

On September 24, 2010 the Cassini orbiter performed its 72th Titan flyby (T72). Table 2.14 below lists the characteristics of both FP3 and FP4 queries in each database. These queries hâve focused on équatorial latitudes. No spectra were found for FP1 for this sélection's parameters. The number of spectra as well as the émission angles are the same in ail database éditions, while the Signal-to-Noise ratio values differ but not significantly. September 2010 sélection at equator from the DS4000-PhC database still suffers ffom ripples on its edges in both FP3 and FP4. The négative différences in Fig. 2.24-lower panel-mean that the nominal DS4000 data hâve higher values in radiance than the DS4000-PhC ones in the FP3 spectral range. Instead, the différences plot in the lower panel of Fig. 2.25 shows that in FP4 the DS4000-PhC spectra hâve radiances higher than the corresponding nominal DS4000 ones. There is also a spike at 797 cm'1, which exists in ail database extracts. The P-wing of acetylene at 730 cm1, is better fitted when using the DS4000 databases The différences of both DS4000 databases are smoother at ethane band at 822 cm' compare to the v3.2 fit (Fig. 2.27 d, e, f). The derived mixing ratios as well as the previous v2.5 édition mixing ratios are listed in Table 2.15 below. No différences hâve been found for the abundances of the 3 major molécules of the FP3 spectral région. I concluded that the différence between the two DS4000 versions is less than 5% with respect to the phase corrected one, having somewhat lower values (Table 2.15).

The first Titan flyby (TO) by the Cassini orbiter occurred on July 3 2004 just after the Satum Orbit Insertion of the spacecraft. Table 2.16 below lists the characteristics of this query in each database. The geometry of these queries follows the restrictions we hâve adapted when studying the previous sélection. FP1 spectra hâve been retrieved. There is no différence in the FP1 spectra, as we can from the Table 2.15 and the Fig. 2.30.

Following the implémentation of the new phase algorithm in the DS4000 database, we fmd that the number of spectra that are now retum from this query in FP3 is significantly reduced by a factor of 10 and therefore we hâve lower Signal-to-Noise ratios in our spectral averages.

Figs. 2.28, 2.29 and 2.30 below depict the comparison between the two DS4000 versions in each focal plane. I hâve not plotted the FP1 DS4000 différences because both spectral averages were identical.

x 10"7 Fig. 2.28: FP3 averaged spectra from CIRSFRANCE. The upper panel shows the results of the database éditions (local calibration in blue, DS4000 in green and the new DS4000 with the phase correction algorithm in red). The lower panel shows the différence between the two DS4000 versions (black).

The new algorithm produces higher radiances at 570 cm'1 compared to the nominal version, and négative radiances at around 765 cm'1. Négative radiances are also found at around 960 The T23 and T24 flybys of Titan took place in January 2007. Table 2.17 below lists the characteristics of this query in each database. As previously, ail queries hâve been restricted to émission angles less than 65°, focused on équatorial latitudes, the spacecraft range was restricted frorn 100 to 400,000 km and ail longitudes were accepted. I now focus on the northem latitudes at 50°N. The number of spectra as well as the émission angles are the same in ail database éditions, while the Signal-to-Noise ratio values slightly differ. The continuum of this sélection in its v3.2 version suffers front négative radiances below 1000 cm'1 (Fig. 2.31) while, in the DS4000 database inference these spectral features do not exist. The spike at 573 cm'1 appears again in deep space calibration data. Négative radiance from 1000 to 1100 cm'1 still exists, but the problem with the bad pixel at 765 cm'1 has been fixed.

DS4000 and the DS4000-PhC are exactly the same and there is a question conceming the extent of the phase correction algorithm application in the DS4000.1 ran our radiative transfer code in order to fit the data of these three database inferences. The fits of ethane are plotted in The abundances from the best-fit process for the 5 major molécules of the FP3 spectral région from the fit of the three inferences (DS4000-PhC, DS4000 and v3.2) are listed in the Table 2.18 below. Ethane is better fitted with using the DS4000 inferences (Fig. 2.32). Ethane's abundance is slightly higher at about 5% using v3.2 to our calculations than using the rest databases, but located within the error bars. The rest molécules hâve similar mixing ratios using ail the databases. The DS4000 database optimizes the continuum (Fig. 2.34). The différences between v2.5 and DS4000 for this sélection are comparable to the January 2010 spectra (Table 2.10).

I hâve had the opportunity get access to the most recent CIRS databases, which are significantly improved, with respect to what was before and validate them. I hâve used five CIRSDATA éditions: the current pipeline v3.2 and the previous v2.5, which hâve been locally calibrated, and the DS4000 and the phase corrected DS4000, with deep space calibration. Additionally, I also validated the Grand Average database, which employs deep space and shutter-closed interferograms for calibration averaged over severalyears.

To begin with the lutter, the Grand Average calibration shows a better continuum level before 600 cm1, compared to the DS4000 and 1.5 to 2 times higher continuum between 600-890 cm'1. It also diminishes the négative radiances, successfully. However, the Grand Average still contains spikes. Instead, most of the spikes are removed in DS4000.

The derived abundances for the major molécules in the FP3 in Grand Average are higher on the order of 15-20% compared to the DS4000. Eventually, the Grand Average test database is still significantly differentfront DS4000 and yields higher abundances.

The DS4000 calibration of CIRSDATA is a large improvement with despiking.

Indeed, the DS4000 version diminishes in several cases the noise in the continuum around 570 and 1050 cm'1 and fixes the bad pixel problem at 765 cm'1. The abundances between DS4000 and the v3.2 are the almost the same in recent queries and their continuum levels hâve no différences. The retrieved abundances using DS4000 are more accurate compared to the v3.2 ones. The model shows that the shift I had to apply to my radiative transfer code calculations is smaller than the one in v3.2 and gives a better line shape.

However, when using old data queries, DS4000 lias higher abundances and higher continuum level compared to the v3.2 (and the v2.5). The spike at 573 cm'1 appears again in DS4000 calibration data, although it lias been vanished in current pipeline with local calibration in some cases. Moreover, DS4000 still suffers from ripples.

The new DS4000-PhC version of CIRSDATA, with respect to DS4000, reduces the négative radiances. In our radiative transfer simulations, the derived abundances hâve no significant variation comparing to the nominal DS4000 édition. However, ripples exist in several spectral sélections. There is a question if the new phase correction algorithm has been applied to the entire DS4000 database, since the queries in some sélections from the two DS4000 versions are identical.

/ recommend the phase-corrected DS4000 database in future studies of Titan CIRS data for ail users. If the newest database is not available, the DS4000 is still an improvement with respect to the v3.2.

Chapter 3

The Titan Atmospheric Model to retrieve the atmospheric Chemical composition. Ail the assumptions will be also defined.

3.1

The structure of Titan's atmosphère Titan is the only planetary body of the Solar System, except for the Earth, that possesses a thick, nitrogen-dominated atmosphère (which créâtes a pressure of about 1,5 bar at the surface). On Titan, methane is the second most abundant atmospheric species with a mixing ratio of 2.19 ±0.002 in the upper atmosphère [START_REF] Waite | Ion Neutral Mass Spectrometer Results from the First Flyby of Titan[END_REF], 1.6 ±0.5 in the stratosphère (Flasar et al., 2005), 1.48 ±0.09% in the lower stratosphère (139.8 -75.5 km) and 5.65±0.18 near the surface (6.7 km to surface) (Niemann et al., 2010). The interaction with the solar influx, the combination and dissociation of these two mother constituents initiâtes complex photochemical reactions in the upper atmosphère, which enrich the organic inventory of the satellite. Some of these organics condense and form a thick haze layer of aérosols covering the satellite globally or partially. The mixing ratios of these gases vary as a function of altitude, latitude and longitude and also time, due to atmospheric circulation. Titan is a cold environment, located at 10 AU and receiving 100 times less energy from the Sun, compared to the Earth.

Titan's température profile was first determined extensively by [START_REF] Lindal | The atmosphère of Titan: An analysis of the Voyager 1 radio occultation measurements[END_REF] from egress and ingress radio-occultation observations by Voyager 1. Its température structure was shown to resemble the terrestrial one, divided into régions with a clearly defined tropopause, stratopause and mesopause. As in the Earth, the température profile of Titan's atmosphère shows inversions, caused by local heating. The mid-atmospheric heating source in the Earth is the stratospheric ozone layer, while the haze layer in Titan plays such a rôle. Both of these heat the local atmosphère by absorbing solar energy and give local température maximum (Coustenis & Taylor, 2008).

The Huygens probe obtained, among other, Titan's vertical atmospheric thermal structure during its descent in January 2005, before landing at 10.3°S and 192.3°W (Lebreton et al., 2005). The Huygens Atmospheric Structure Instrument [START_REF] Fulchignoni | The Characterisation of Titan's Atmospheric Physical Properties by the Huygens Atmospheric Structure Instrument (Hasi)[END_REF] measured the décélération of the probe ffom 1400 to 155 km and determined the atmospheric mass density. By assuming that the atmosphère is in hydrostatic equilibrium, the pressure profile can be inferred. Then, by considering the atmosphère at these levels as a perfect gas and assuming a value for the atmospheric mean molecular weight, the température as a fonction of altitude is derived.

At the altitude of 155 km, Huygens released its protection shield and HASI measured directly the atmospheric pressure and température by using its sensors from this level down to the surface by better than 1 K accuracy (Fulchignoni et al., 2005). The density and température profile derived from HASI is depicted in Figure 3.1 from the surface up to 1380 km.

HASI showed that Titan has a well-defined tropopause at about 44 km with 70.43 K.

The HASI température measurements for this région are in agreement (1-2 K) with the Voyager measurements [START_REF] Lindal | The atmosphère of Titan: An analysis of the Voyager 1 radio occultation measurements[END_REF]. The stratopause is located at the altitude of 265 km with 188 K, 13 K warmer than the Voyager IRIS results. The mesopause has been identified at 494 km with 153 K. In this région, the températures are 5-10 K higher than the model prédictions [START_REF] Yelle | The Yelle Titan Atmosphère Engineering Models[END_REF]. For altitudes over 500 km and up to 1020 km, which is the thermosphère, HASI shows températures variations of around 170 K (varying within ±10-15 K), caused by inversion layers, dynamic phenomena like gravity waves and graviational tides (Fulchignoni et al., 2005). In this part of the atmosphère, the HASI density measurements are higher by a factor of 2 compared to the re-interpretation of Voyager solar occultation measurements [START_REF] Vervack | New perspectives on Titan's upper atmosphère from a reanalysis of the Voyager 1 UVS solar occultations[END_REF]. (1997) as well as the measurements of the Voyager radio occultation experimented revisited by [START_REF] Lellouch | Titan's atmosphère and hypothesized océan: A reanalysis of the Voyager 1 radio-occultation and IRIS 7.7-pm data[END_REF]. No considérable différence exists from the surface up to an altitude (1989) Voyager radio occultation data are also shown for comparison (adapted from Strobel (2010). CIRS retrieved température profile from the T3 Cassini Titan flyby is plotted (red solid curve) (Flasar et al., 2005).

INMS density results are also shown (blue solid line) [START_REF] Muller-Wodarg | Horizontal structures and dynamics of Titan's thermosphère[END_REF]).

Titan's atmosphère is more extensive than the atmosphère of the Earth and its exosphere begins much higher. Indeed, the homopause level is located at 1,195 ±65 km according to Cassini/INMS mass spectrometer results from TA flyby (80-90 km on Earth), while the exobase is found to be at 1429±5 km [START_REF] Waite | Ion Neutral Mass Spectrometer Results from the First Flyby of Titan[END_REF] (400-500 km on Earth).

3.1.1

The Troposphère

The troposphère, the atmospheric part which lays above the surface, has a négative température vertical gradient. HASI measured the surface température at 93.65 ±0.25 K (Fulchignoni et al., 2005). Only 10% of the incoming solar radiation reaches Titan's surface.

The other solar photons are scattered by haze and cloud layers. Part of this scattered radiation is reflected back to space, while some of it is absorbed by the aérosols and the atmospheric gases (Coustenis & Taylor, 2008). The surface absorbs the incident radiation and heats the lower atmosphère, producing atmospheric convection depending on the local température gradient.

The infrared émission of the surface heats the lower tropospheric layers, which become less dense than the colder layers above. The density différence produces vertical movement of the atmospheric layers and the hotter layers rise, while the colder fall. The convection does not occur when the température gradient is shallow and it is only possible when the gradient value reaches the adiabatic lapse rate. For Titan, this rate is at 1.4 K/km.

Such convective instabilities occur in the lower part of optically thick atmosphères. In the upper part of the atmosphère, the density becomes quite low due to radiative cooling caused by infrared émission to the space and reduces the température. The température gradient becomes sub-adiabatic and the convection stops. Titan's troposphère reaches a minimum of 70.43 K at 44 km with a pressure of 115 ± 1 hPa (Fulchignoni et al., 2005). In this level, a température inversion occurs and the stratosphère begins.

3.1.2

The Stratosphère

The lower part of Titan's stratosphère is quasi-isothermal near the altitude of 50 km. No absorption occurs in the lower stratosphère, above the tropopause, and the température lapse rate is almost zéro, which indicates that the température is constant with height. In the mid and upper Titan's stratosphère, similar to Earth's, the température increases with the altitude.

From 50 km to 300 km, the température increases due to the heating caused by the absorption of UV and visible solar radiation.

In Titan's stratosphère, the température has a high static stability with a lapse rate at about lK/km. The convection stops and the atmosphère is stratified in layers, which are mixed. On Earth, the ozone layer is the main absorber and the heating source of the stratosphère. On Titan, aérosols and methane are the main radiation absorbers. The aérosols, the methane, the acetylene and the ethane re-emit isotropically in infrared wavelengths and this mechanism cools the stratosphère, balancing the solar heating. The température reaches a maximum value of 186 K at the altitude of 250 km, as derived from HASI measurements (Fulchignoni et al., 2005), which defines the location of the stratopause.

Cassini/CIRS also retrieved the atmospheric température from remote sensing probing.

Its results are not consistent with HASI measurements about the location of the stratopause which CIRS places at about 312 with 183 K (Flasar et al., 2005;Vinatier et al., 2007a). This différence is illustrated in Figure 3.1. This discrepancy and how it affects radiation transfer calculations is discussed in Section 3.2 below.

3.1.3

The Mesosphere Above the stratopause, the température decreases again with the altitude and it reaches a minimum of 152 K at 490 km, the mesopause (Fulchignoni et al., 2005). In this région, HASI measured perturbations due to dynamic phenomena and inversion layers. Between the tropopause and the mesopause, the atmosphère remains fairly well mixed by turbulence produced by a variety of instabilities in wave motions and the mean flow (Coustenis & Taylor, 2008).

The Thermosphère

Above the mesopause, the atmospheric température rises up to the exobase due to the absorption of solar EUV from nitrogen, methane and acetylene [START_REF] Lellouch | Titan's thermosphère profile[END_REF]. In these altitudes, the pressure values are a few microbars and gas densities are quite low. The solar EUV photons as well as energetic particles from the Satumian magnetosphere, pass through this layer, causing ionization, dissociation of the molécules and releasing kinetic energy. Such mechanisms heat the gas and the température gradient increases rapidly with height.

The thermosphère is the most extensive part of the atmosphère, in which the energy is transported by thermal conduction. Diffusion is the dominant process and the atmosphère starts to separate into its lighter and heavier components. CIRS températures profiles are retrieved from an inversion method, which is described in Appendix A. In brief, the température profile is retrieved by using the methane V4 band at 1304 cm'1 as a thermometer in the stratosphère. Methane's mole fraction versus altitude is well known from Huygens/GC-MS in situ measurements from the surface up to about 140 km (Fig. 3.2) and we adopt it to retrieve the température (Niemann et al., 2010). By introducing the FP4 CIRS spectra into an inverse algorithm (Achterberg et al., 2008;2011), we retrieve the température profile related to each CIRS sélection of interest. This algorithm solves the radiative équation, which is an inverse ill-posed problem, by considering a reference température profile for optimizing the results as the initial guess. Two are the main profiles used as the original guess in the retrievals, the one derived from Huygens/HASI (Fulchignoni et al., 2005) and the second retrieved from Cassini/CIRS T3 Titan flyby (Flasar et al., 2005) (Fig. 3.1). The HASI profile is a resuit of in situ measurements during the Huygens probe descent. From 1400 km down to 155 km, the température was retrieved indirectly from the accélération measurements and by assuming hydrostatic equilibrium and an idéal gas State for this atmospheric région.

Below 155 km and after the thermal shield release, HASI measured directly the température and the pressure by using sensors exposed to the environment. CIRS T3 profile at 15°S was retrieved from a combination of CIRS nadir and limb-viewing spectra. Voyager radio-occultation data [START_REF] Lindal | The atmosphère of Titan: An analysis of the Voyager 1 radio occultation measurements[END_REF] as well as radiative mesospheric models [START_REF] Yelle | Non-LTE models of Titan's upper atmosphère[END_REF] were used as the initial guess in order to interpolate the values between nadir and limb spectra (5-60 mbar) (Flasar et al., 2005). Since HASI data were obtained by in situ measurements, its results should be considered as reliable for the équatorial latitudes.

The stratopause température and level discrepancy between CIRS and HASI observations

The CIRS température profile retrieval at 15°S from T3 flyby is close to the HASI landing date and descent path but it does not agréé in the upper stratosphère with the HASI results (Fig. 3. 1). HASI shows a hotter stratosphère than CIRS, with the stratopause located at about 250 km at 186 K (Fig. 3.3). The CIRS profile locates the stratopause higher at 316 K with 183 K (Flasar et al., 2005).

The cause of this discrepancy remains unsolved. I spent, however, some time at the beginning of my PhD interacting with both with the CIRS and the HASI team (Prof. M.

Fulchignoni and Dr. F. Ferri) and with the help of Dr. D. Strobel, trying to shed some light to this problem. We also discussed the Huygens probe descent parameters with ESA représentatives, among which J-P. Lebreton and P. Couzin from Thaïes Alenia Space, ESA's prime contractor for building the Huygens probe.

The HASI température at this level being 185 K, leads to a computation of the inffared radiance with the HASI derived atmosphère inconsistent with the observed CIRS radiance of the 7.7 fim methane band (Lebreton et al., 2009;[START_REF] Strobel | Atmospheric Structure and Composition[END_REF]. The methane V4 band probes the 150-350 km région, around the stratopause and the HASI profile is too warm. We also hâve ground-based observations in the submillimeter range, from the Institut de Radio Astronomie Millimétrique (IRAM)/Plateau de Bures telescope that can help in this problem. IRAM heterodyne measurements inferred a température profile sounding altitude range from 80 to 800 km at very high spectral resolution. Carbon monoxide line wings probe the Titan atmosphère from 80 to 180 km [START_REF] Hidayat | Millimeter and Submillimeter Heterodyne Observations of Titan: The Vertical Profile of Carbon Monoxide in Its Stratosphère[END_REF] and at these levels the HASI profile is similar with the IRAM one, while the CIRS one is too cold (Fig. 3.4). Hydrogen cyanide lines probe higher altitudes (300-500 km in the center and 100-200 km in the wings) [START_REF] Hidayat | Millimeter and Submillimeter Heterodyne Observations of Titan: Retrieval of the Vertical Profile of HCN and the 12C/13C Ratio[END_REF] and the HASI profile is too cold. In conclusion, without having solved the problem, we can say that the HASI near-equator data directly sounding the lower atmosphère are the more reliable, while at higher altitudes, where HASI only indirectly infers the température, CIRS inferences may be the correct ones.

HASI

Impact of the original assumption in the température calculations

In order to deal with the question of how the radiative transfer calculations are affected by the choice of the initial profile in température retrievals, I hâve performed radiative transfer calculations on the same CIRS data FP3 and FP4 queries. I hâve selected 4 queries from CIRSDATA, two at équatorial latitudinal bins and two at 50°N. Since HASI data were obtained at 10°S, the application of CIRS T3-based température profile at équatorial sélections can give an estimation of the température profile influence to the resulted abundances. On the other hand, I ran the same test in northem sélections where the stratospheric températures are indeed lower compared to the équatorial ones. Table 3.1 below lists the characteristics of the sélections used in these tests. The retrieved température profiles for each sélection, using both HASI and CIRS-T3 profiles as the initial guess, are illustrated in Figure 3 At mid-latitudes ail molecular abundances are within 10%, except for HCN, witli the CIRSbased profile generally giving higher abundances. This is expected, because the HASI profile was inferred above the Huygens landing site, which was at 10°S. In the northern latitudes, the divergence depending on the profile is more pronounced, with most of the molecular abundances still within 10% whatever the température profile. Only the émission bands of C2H2 and HCN are more affected by the température variation as expected, still within the 20% level, since the HASI original assumption gives rise to higher abundances.

3.3

The opacity sources on Titan Voyager 1 observations showed that the main infrared opacity sources of the Titan's atmosphère are the main molécules' absorption, N2, CH4 and H2. Usually symmetrical atoms like N2 do not interact with the infrared radiation field due to the fact that they lack dipole moment. However, at longer wavelengths collisions between these molécules produce a temporal dipole moment and therefore opacity (Coustenis & Taylor, 2008). The trace gaseous species also contribute to the infrared opacity of the atmosphère (Fig. 3.6). the stratosphère caused by the aérosols (Coustenis et al., 2007). The température is adjusted from an initial guess profile and the surface température is considered to be at 93.7 K according to Huygens/HASI in situ measurement after the probe's touchdown (Fulchignoni et al., 2005) at 10°S. CIRS brightness measurements at 19 pm showed that the surface températures in North pôle is 3 K lower and at the South pôle 2 K lower, compared to the équatorial one [START_REF] Jennings | Titan's Surface Brightness Températures[END_REF].

The aérosols are the end product of the complex photochemistry that occurs on Titan's upper atmosphère and they strongly affect the thermal balance of the entire satellite. They are the main heating source of the stratosphère, since they absorb the 40-60% of the sunlight, depending on the altitude, as measured by the Huygens/DISR (Tomasko et al., 2008a).

Vl/IRIS measurements provided the aérosol distribution between the tropopause at 40 km and the upper stratosphère at 300 km [START_REF] Samuelson | Thermal infrared properties of Titan's stratospheric aérosol[END_REF]. They showed that the stratospheric opacities increase monotonically with the wavenumber from 250 to 600 cm'1.

According to experimental results, aérosols are dark in the visible and bright in the infrared [START_REF] Ramirez | Complex Refractive Index of Titan's Aérosol Analogues in the 200-900 nm Domain[END_REF]. DISR measured directly the haze particles' vertical distribution from the altitude of 150 km down to the surface [START_REF] Tomasko | Rain, winds and haze during the Huygens probe's descent to Titan's surface[END_REF].

As Cassini approaches Titan during each flyby, its radial component of the velocity is négative, it goes to zéro at closest approach, and then it becomes positive. For a nadir viewing geometry this should, in theory, produce a Doppler shift on the observed CIRS spectra. It is probably a very small effect, but the question is if we can give an évaluation of it.

After personal communication with Dr. C. Nixon from Goddard/NASA, we evaluated the Doppler effect and came to the conclusion that is probably too small to significantly affect CIRS results. A rough estimation of the size follows below. Cassini approaches Titan at typically 20,000 km/h, or 5.5 km/s. The Doppler shift is defïned as:

AL Àv ~L=~5 .5 3jc105 2*1 0"5
Thus, the shift is about 2 parts in 100,000, applied to frequency/wavenumber. For example, at 1000 cm'1 the shift should be about 0.02 cm'1, but at 100 cm'1 would be 0.002 cm'1, while the maximum would be about 0.03 cm'1 at 1500 cm'1. However, larger wavenumber shifts hâve been noticed, likely due to other effects, they can be noticeable (~0.1 cm'1) and should be corrected for empirically.

Radiative Transfer on Titan

The theory of radiative transfer is a quantitative assessment of the relation between the interactions of the matter with radiation (absorption, scattering and émission of radiant energy on a microscopie scale) and the resulting radiation field [START_REF] Hanel | Exploration of the Solar System by Infrared Remote Sensing[END_REF]. Atmospheric émission, which at planetary températures occurs mainly in the thermal-infrared part of the spectrum, dépends on the composition and thermal structure of the atmosphère. Absorption and scattering mainly occur at ultraviolet, visible and near-infrared wavelengths and dépend on the prevailing molecular opacity and on aérosol and cloud properties.

3.4.1

The radiative transfer process

The émission of an infrared source is described both by Planck's law and Kirchhoff s law.

The former gives the radiant energy emitted by a black body at the same température and wavelength as the source, while the latter détermines that the absorption of a body equals to its emissivity. The radiative transfer équation (RTE) describes how the optical properties of the medium are related to the spectrum we measure.

When we hâve a beam which passes through a medium, a portion is absorbed by the material due its physical properties and according to Kirchhoff s law the body should émit the same amount of heat as if it was a black body in the same température described by Planck's law. In physical terms, while a radiation with the intensity I{(X) passes through a path within a medium of a distance Ax, the medium absorbs a fraction of kAxI(/li) /and emits a fraction of thermal radiation of KAxB/fi [START_REF] Twomey | Introduction to the Mathematics of Inversion in Remote Sensing: Courier Dover Publications[END_REF] where ^is the cosine of the émission angle measured from the vertical, k is the absorption coefficient of the medium and B(v,T)

represents the Planck radiance (B(v,T) = 2hv:

1 2 "hvlkT c e -1
) at the frequency v and the température T. Hence, the fraction of the total intensity which propagates through the medium can be written as:

A/<» = absorbed + 1emitted = Xl(fj) + ~KAxB = er/"(-/(rt + £) dr However, since ---= exl,x -+ I(fi)er/,x -the RTE can be written as:

dr dr (x 4^))_ iei,B dr fj,
Then, by integrating both parts of the latter relation in respect to the optical depth, we dérivé:

\d(er/,iI(iu)) \ i f _V yrlldT = f B-eT/tidr o dr o P1
The intensity of the radiation is the initial I0 when r = 0 at the level of emergence and I=I(i0) when r = z0 at the level of émission and by applying these boundary values in the above intégral équation we dérivé [START_REF] Hanel | Exploration of the Solar System by Infrared Remote Sensing[END_REF]:

T°. dr = le''"1 + f (1) o B
Equation ( 1) describes the intensity of a propagating radiation through layered gaseous media such as planetary atmosphères. The first term of the second part of (1) corresponds to the surface contribution, while the second term depicts the atmospheric contribution.

The emerging radiation from the surface I0 is the fraction of the Planck function émission at the same température. According to the Beer-Lambert law, the transmittance of the atmosphère Tr can replace the exponential expression in (1). Therefore, following the above considération, the spectral radiance measured remotely by the CIRS instrument is [START_REF] Hanel | Exploration of the Solar System by Infrared Remote Sensing[END_REF]:

= evB(v,T,)Tr<ji,v,z,) + / B(v,T(z)) dT^'v'z) dz (2) , dz
where /i=cosine of the émission angle and v=frequency

Here, the termev is the emissivity8 of the surface at the température Ta. The altitude z can be derived by the équation z = -ln/?, where the atmospheric pressure is denoted by the symbol p. With the index "s" the surface level values are indicated. The ztop refers to the highest altitude from the surface where the atmospheric gases sense the recorded radiance.

The transmittance Tr can also be written in the following form:

Tr(n,v,z) = f <j>(yy)e-'lv'i)ll'dv' (3) <5v

The terni <p(v,v') is the instrument spectral response fonction9 extending over the spectral interval <5v with central number v. By définition, the transmittance is the rate of the radiation, which passes a gaseous sample versus the incident one. The exponential expression in équation ( 3) cornes from the Beer-Lambert law where r is the optical depth.

On Titan, since the stratospheric haze absorbs most of the surface radiance, the first term of the équation ( 2) can be neglected. This assumption simplifies the équation (2) in the form [START_REF] Hanel | Exploration of the Solar System by Infrared Remote Sensing[END_REF]:

/(z,M) = JB<v,T(z))âTr(^V'Z> dz (4)
{ dz

It should be noted, however, that both scattering and surface reflection hâve been neglected in (2) and ( 4). Indeed, in the thermal part of the spectrum, where absorption and émission by atmospheric gases dominate and solar radiation is negligible, scattering processes can often be neglected and the solution becomes considerably simpler [START_REF] Hanel | Exploration of the Solar System by Infrared Remote Sensing[END_REF].

The solution of the radiative transfer équation

In physical terms, the intensity of the radiation î(p, v) of the équation ( 4) is known (the recorded data), while what causes it (the quantity in the intégral) is unknown. Hence, the solution of the équation (4) provides the transmission of the absorbing gasses of the atmosphère related to the recorded radiation and, eventually, the composition of the probed atmospheric région as well as the thermal structure.

The problem is how to retrieve the atmospheric parameters from a set of measurements. The infrared data dépends on the spectral range, the spectral resolution, the 8 The emissivity is the ratio of the radiated energy by the surface to relative black body energy radiated (for the black body £v=l) 9 The spectral response function détermines the ability of an instrument to record radiation in a spécifie wavelength comparing to the one in another wavenumber where the response is unity. Practically, it évaluâtes the operational spectral range.

observational geometry and the signal-to-noise (S/N) ratio [START_REF] Hanel | Exploration of the Solar System by Infrared Remote Sensing[END_REF]. Several models hâve been proposed to solve équation (4). These solutions can be divided into two kinds:

a) To those which solve directly the équation ( 4) by producing a model which is then compared to the data and b) To those which solve indirectly the relation ( 4) by importing the recorded data to an inversion algorithm.

The forward approach is the direct intégration over the frequencies of équation ( 4) and the use of the fine parameters of the absorbing bands [START_REF] Scott | A direct method of computation of the transmission function of an inhomogeneous gaseous medium -I: Description of the method[END_REF] to produce a synthetic model of the atmosphère. By assuming the abundances of the contributing species and following an itérative best-fit process, the model spectrum is fitted to the real data. Then, the atmospheric composition can be retrieved.

The inversion approach is an alternative tool that allows for the calculation of the values of atmospheric parameters from the recorded radiation intensities by solving the inverse problem. Following a reversai sequence, we begin from the results of a physical process and seek for the embedded parameters which hâve caused the recorded results [START_REF] Tarantola | Inverse Problem Theory and Methods for Model Parameter Estimation: Society for Industrial Mathematics[END_REF].

Most of the models proposed so far hâve been used for radiative transfer calculations in the terrestrial atmosphère [START_REF] Snell | FASCODE for the environment (FASE)[END_REF][START_REF] Berk | MODTRAN Cloud and Multiple Scattering Upgrades with Application to AVIRIS[END_REF][START_REF] Stamnes | Numerically stable algorithm for discreteordinate-method radiative transfer in multiple scattering and emitting layered media[END_REF].

Apparently, they contain assumptions and parameters which satisfy the conditions on the Earth but they need modifications before applied in other planetary atmosphères (Irwin et al., 2008).

Cassini/CIRS as well as its predecessor Voyager 1/IRIS hâve probed the infrared part of the spectrum of Titan's atmosphère and sent back to the Earth their recordings. Here, I describe the forward model I use to solve the radiative transfer équation (4) for interpreting CIRS FP1 and FP3 spectra.

We fïrstly consider an atmosphère in hydrostatic equilibrium with a vertical température structure and without any scattering effects. We model the atmosphère by dividing it into 80 variable individual pressure levels. Each layer is determined by the pressure values in its boundaries Pj.j and while the température in each one varies from (Tmin)î and (Tmax)i. Considering a spectral région limited by frequencies va and vt, we apply a fast line-by-line code introduced by [START_REF] Scott | A direct method of computation of the transmission function of an inhomogeneous gaseous medium -I: Description of the method[END_REF] for ail the layers and for ail the températures of each layer from (Tmin)j to (Tmax)i [START_REF] Scott | A fast line-by-line method for atmospheric absorption computations -The Automatized Atmospheric Absorption Atlas[END_REF]. The températures are imported by the température profile retrieved from the FP4 CIRS recordings using a constrained inverse algorithm (Achterberg et al., 2008;2011).

For each individual layer we assume local thermodynamic equilibrium (LTE) which means that the source fonction of each layer is given by the Planck fonction, stable in the layer between its top (ztop) and bottom (zs), at the corresponding level température. According to Houghton, the LTE condition does not exist at high altitudes due to the différence between the Planck and the source fonctions in these layers [START_REF] Houghton | The physics of atmosphères[END_REF]. [START_REF] Coustenis | Titan's atmosphère from voyager infrared observations : I. The gas composition ofTitan's équatorial région[END_REF] calculated this altitude at 8x1O'3 mbar or 550 km and set this value as the upper limit of the solution of équation ( 4).

The radiative transfer code intégrâtes directly in frequency and uses the fine structure parameters of the absorbing gases [START_REF] Scott | A direct method of computation of the transmission function of an inhomogeneous gaseous medium -I: Description of the method[END_REF]. These spectroscopic parameters incorporated into the code hâve been provided from laboratory results through spectroscopic databases.

Since we hâve an atmosphère in non-isothermal State, the température contrast between different pressure levels of the atmosphère causes spectral variation. Then, the emerging radiation of équation ( 4) for each layer can be written as:

Ztop I(z,n)= fB(z)W(z,fi)dz *s
The fonction W (z,/u) is the weighting fonction of an opacity source. By définition, W(z,^) is the partial dérivative of the transmittance in the direction of the (the cosine of the émission ôT (u,v,z) angle), that is W (z,[i) = -----. Since the Planck fonction varies quite rapidly with the dz height, the peak of the contribution fonction C(z,/lï) :

C(z,/u) = B(z)W (z,v)
is more appropriate to describe the level of the maximum contribution of this source to the emerging radiation [START_REF] Hanel | Exploration of the Solar System by Infrared Remote Sensing[END_REF]. The monochromatic radiance is then calculated by incorporating the contribution functions of the opacity sources to équation (4). The outcome of this procedure is a model spectrum which is then convolved with the Fourier Transform of the Hamming function used for smoothing the edges of the interferograms.

In our calculations we take into account the absorption of 39 molécules and their isotopologues. For each species, we hâve used the most recent spectroscopic parameters as they are published in online spectroscopic databases HITRAN (Rothman et al., 2009) and GEISA10 (Jacquinet- Husson et al., 2011).

For deriving the abundances of Titan's trace stratospheric constituents we use constant-with-height distributions as well as vertically-varying distribution profiles. In the former case, the abundances are assumed to be constant-with-height above the condensation level and set to follow the corresponding saturation law below (Coustenis et al., 2007). The 2-D General Circulation Model (GCM) provides the latter profiles (Rannou et al., 2005).

However, the minor trace gases of the stratosphère hâve not strong enough spectral lines in their émission bands to apply the relative vertical distribution profiles. The vertical mixing ratios are applied only for acetylene, hydrogen cyanide and ethane, since they probe multiple altitudes via optical thin and thick lines [START_REF] Coustenis | Titan trace gaseous composition from CIRS at the end of the Cassini-Huygens prime mission[END_REF].

The fît of the continuum of the recorded spectrum is acquired by using a simple twocloud System (Borysow & Tang, 1993;[START_REF] Coustenis | Titan's atmosphère from Voyager infrared observations. IV: Latitudinal variations of température and composition[END_REF][START_REF] Samuelson | Gaseous abundances and methane supersaturation in Titan's troposphère[END_REF] fitting the continuum at different latitudes, starting out with CIRS limb data and DISR measurements at the équatorial latitudes (Vinatier et al., 2010a;[START_REF] Vinatier | Optical constants of Titan's stratospheric aérosols in the 70-1500 cm'1 spectral range constrained by Cassini/CIRS observations[END_REF]. We then adjust the haze opacity parameters in order to achieve the best fit of the continuum.

The Atmospheric Radiative Transfer code for Titan (ARTT)

The Atmospheric Radiative Transfer code for Titan (ARTT) is a forward model which solves the radiative transfer équation ( 4) for the atmosphère of Titan. It is a fast line-by-line and layer-by-layer method which sums ail the contributions of the absorbing gases in an inhomogeneous atmosphère [START_REF] Scott | A direct method of computation of the transmission function of an inhomogeneous gaseous medium -I: Description of the method[END_REF][START_REF] Scott | A fast line-by-line method for atmospheric absorption computations -The Automatized Atmospheric Absorption Atlas[END_REF].

The method incorporâtes both the atmospheric inhomogeneities and the instrument function and it directly intégrâtes the radiance over a band of spécifie frequencies in order to compute the radiance in a single layer. It is called line-by-line due to the calculation of each spectral line of the atmospheric gases in the layers of interest. For this purpose, it adapts the spectroscopic files of the molécules of interest as inputs in the code which are provided from laboratory studies. It relies on (a) a précisé température profile and (b) on the correct représentation of ail opacity sources (Coustenis et al., 1989a).

ARTT is based on the précisé knowledge of the thermal structure or of the opacity sources of the atmosphère. Then, it produces a model spectrum which is compared to the real data to infer one or the other of the parameters. An itérative process during which we adjust the abundances of the molécules of interest, allows us to obtain best fîts of the data. We firstly dérivé the transmittances monochromatically for the whole spectral région for a determined set of atmospheric conditions and then we reconstruct them for any atmospheric conditions to generate a mean weighted spectrum.

The scheme of the fitting procedure is the following. We firstly incorporate a précisé model of thermal structure of the probed atmospheric région to the synthetic spectra.

We then also assume either constant-with-height abundances or vertical distributions of the absorbing gases. By an itérative best-fit process, we fit the synthetic spectra to the real data at FP1 (10-600 cm"1) or at FP3 (600-1100cm'1).

ARTT takes the gravity and the radius of Titan as constant values considering the satellite as a perfect sphere. Moreover, it takes into account the sécant of the observations, which can be modified by the user. The other inputs of the software are the pressure levels and the température vertical profile which can be adjusted to fit the spectra. The code then simulâtes the vertical pressure and température distribution for a spécifie sélection. This is necessary in order to model the Planck fonction for the recorded infrared radiation.

Afterwards, we induce the mixing ratios of the major molécules (nitrogen, methane, hydrogen and argon) of the atmosphère.

The next step is to incorporate the spectroscopic parameters of the trace gases. ARTT is a fast line-by-line code and therefore it reads ail the line lists from each file and calculâtes the relative contribution. Each molecular mixing ratio can be constant-with-height or vary vertically with height. Ail the available spectroscopic line lists hâve been incorporated in the software during the installation procedure. Finally, the continuum is simulated through the presence of a two-cloud model which reproduces the opacity in between lines, and is due to the absorption from aérosols or particle condensâtes. The haze distribution we apply to the data is adapted from [START_REF] Vinatier | Optical constants of Titan's stratospheric aérosols in the 70-1500 cm'1 spectral range constrained by Cassini/CIRS observations[END_REF] spectral discription near the equator that we adjust to match our spectra at different latitudes.

After having extracted the model, the code compares it with the recorded spectra. This comparison is the outcome of the whole procedure and leads to the détermination of the thermal and Chemical structure of Titan. The comparison between the synthetic and the observed spectra is dual. In general, we try to accomplish the best possible match of the Qbranch (the center) of ail the molécules. That can be done visually through an itérative process. However, for large and extended émission bands (as for CH4, C2H2 and C2H6 for instance) we may want the best fit throughout the band. In that case, the best fit can be deducted from the différence between the model and the data by minimizing the residuals.

When the différence is minimal, we assume a best fit. We hâve applied a spécifie code using régression analysis to optimize these différences. The software upgrade has optimized the computational time which is literally depended on the hardware. Generally, each radiative 2) Incorrect scale of propane V26 line strengths at 748 cm'1 included in GEISA 2003, the band of which should be multiplied by a factor of 0.420 (Nixon et al., 2009).

3) Isotopologue investigation in Titan atmosphère requires new molécules to be added to the code. One of the aspects we had to tackle was the low flexibility of the code in implementing new spectroscopic data and making any sort of changes in the input parameters.

"Gestion et

First of ail, there was a lack of a readme file with guidelines on how to use the source files in order to create the exécutable files, which we would use in order to create our graphs. In addition, there was no manual on the actions, which had to be taken in order to find how to enrich the molecular database of the program. By studying the workflow of the program and by adding breakpoints into the source code of the program, we managed to identify the actions needed in order to complété the upgrade. Moreover, the use of the program had to be more user-friendly and for this reason we added output messages to the software and replaced or removed any messages left of the debugging process.

The code also had to meet the demands of the increased number of data imported, which increased the computational time. For this reason, we had to remove any delays on the workflow by using less memory resources and by removing some redundant conditional statements and itérative processes. So, the new version has better stability and an increased performance compared to the previous version since the number of the calculated molécules has almost been doubled but the calculations time has remained the same.

The new édition of the radiative transfer software is completely compatible with the files used by the previous one and therefore the migration from the one version to the other is bug-ffee.

Description of the new aérosol model

The ARTT code has also been upgraded in the sense of new aérosol distribution to fit the continuum. As noted previously, the continuum was fitted by adapting a simple two-cloud System, one in the troposphère and the other in the stratosphère (Borysow & Tang, 1993;[START_REF] Coustenis | Titan's atmosphère from Voyager infrared observations. IV: Latitudinal variations of température and composition[END_REF][START_REF] Samuelson | Gaseous abundances and methane supersaturation in Titan's troposphère[END_REF]. As opacity sources were considered the collisions-induced absorption (CIA) between the major molécules (N2, CH4 and H2) and its combinations. In the new ARTT model we hâve adapted the recent aérosol distribution provided by [START_REF] Vinatier | Optical constants of Titan's stratospheric aérosols in the 70-1500 cm'1 spectral range constrained by Cassini/CIRS observations[END_REF], which put constrains at the refractive indices of Titan's aérosols in the mid infrared part of the spectrum. 0.95 mbar (193 km) and Huygens/DISR in situ measurements at 10°S to model the extinction cross-section in mid-infrared (600-1500 cm"1). The extinction cross-section is derived by the ratio of the aérosol volume extinction coefficient to the aérosol number density. The former was retrieved by the CIRS limb data [START_REF] Vinatier | Optical constants of Titan's stratospheric aérosols in the 70-1500 cm'1 spectral range constrained by Cassini/CIRS observations[END_REF], while the latter by DISR measurements. According to (Tomasko et al., 2008b) the aérosol number density of 5 particles per cm"3 at 80 km decreases exponentially with a scale height of 65 km from the altitude of 80 km up to 150 km (Fig. 3.8).

Aérosol Number Density vs. Altitude (2008b). The value of 5 particles per ce at 80 km is the maximum of the haze aérosol concentration. The prédictions of a fractal algorithm, which was used to describe the haze particles, agréé with the measurements. This means that above 80 km some hydrocarbons condense onto fractal aggregate particles, which may fill the voids in particles and therefore cause higher single scattering albedos (Tomasko & West, 2009). [START_REF] Vinatier | Optical constants of Titan's stratospheric aérosols in the 70-1500 cm'1 spectral range constrained by Cassini/CIRS observations[END_REF] assumed the same scale height for altitudes above 150 km. Although no relative measurements exist for higher than 150 km, this assumption agréés with Vinatier et al. (2010a) results. Anderson & Samuelson (2011) hâve derived the aérosol volume extinction coefficient in far-infrared (70-600 cm'1) at 15°S. The calculated extinction crosssection by Vinatier et al., combined with the latter results of is illustrated in Fig. 3.9. Both parts of the complex refractive index are interlinked by a spécial form of Hilbert transforms, which are termed Kramers-Kronig relation (KKR) [START_REF] Lucarini | Kramers-Kronig Relations In Optical Materials Research[END_REF]. By applying the KKR, through an itérative process, the real and the imaginary parts of the refractive index as a function of the wavenumber is calculated numerically (color lines in Fig. 3.9).

Figure 3. 9 -Aérosol extinction cross-section in the range of 30 to 1500 cm-1 (adapted from Vinatier et al.

(2012)). The observed extinction cross-section (black) is plotted with 1-a error bars. The calculation was performed by assuming a constant real part of the refractive index equal to 1.69, while the imaginary part of the refractive index ni varying within the range of 0.01 and 0.2. The color lines are calculations assuming a fractal cluster of 3000 monomers with individual radii of 0.05 pm and fractal dimension of 2 as predicted by [START_REF] Cabane | Formation and growth of photochemical aérosols in Titan's atmosphère[END_REF] and observed by DISR (Tomasko et al., 2008b). The ni used for each model calculation is shown in the legend. The imaginary part of the refractive index of Titan's aérosols varies between 0.018 and 0.2.

The extinction cross-section is the imaginary part of the complex refractive index. In an opaque medium, the real part the complex refractive index of describes the refraction, while its imaginary part the absorption loss. The results of [START_REF] Vinatier | Optical constants of Titan's stratospheric aérosols in the 70-1500 cm'1 spectral range constrained by Cassini/CIRS observations[END_REF] calculations of the complex refractive index are plotted in Fig. 3.10. Within the région spectral région of 600 to 1500 cm'1, the imaginary part of the refractive index présents maxima at 630, 700, 750, 900, 1300 and 1450 cm'1. These results differ from [START_REF] Khare | Optical constants of organic tholins produced in a simulated Titanian atmosphère -From soft X-ray to microwave frequencies[END_REF], being less absorbent than tholins in the thermal infrared (Fig 3.10). Similar results hâve been reported in the literature [START_REF] De Kok | Characteristics of Titan's stratospheric aérosols and condensate clouds from Cassini CIRS far-infrared spectra[END_REF][START_REF] Rannou | Titan haze distribution and optical properties retrieved from recent observations[END_REF]Vinatier et al., 2010a;Anderson & Samuelson, 2011). By integrating the volume extinction coefficient over the altitude for two given altitudes, we dérivé the optical depth of the haze aérosols. I hâve applied these values to our radiative transfer code to fit the continuum. However, although this distribution can be applied in ail latitudes, there are différences between the north to the south hemisphere and these values should be adjusted in order to obtain the best fit. In Chapter 4, I show that this description if adjusted in latitude can provide results as to the haze distribution on Titan's disk.

The updated spectroscopic parameterization

The spectroscopic parameters of the main molécules (CH4, CH3D, CO, CO2, HC3N, C2N2, C2H2, C2H4, C4H2, C3H4, H20, C6H6, C2HD, HCN, HC15N, C2H6, C3H8 and CH3CN) and the e) For the acetonitrile (CH3CN), GEISA 2009 provides data ranging from 890 to 1650 cm'1. These lines are imported in the code. However, for the région of 318 to 1135 cm'1 the PNNL data from the previous ARTT version was kept.

As a second step, we hâve added new isotopologues in the new release of ARTT. The identification of isotopologues in Titan's atmosphère is essential for radiative transfer studies.

Isotopologues are molecular entities which differ in their isotopic composition. The isotopic compounds hâve at least one atom which consists of a different number of neutrons from the one of the main molécule, it exhibits different vibration States and can therefore be observed in infrared spectra. The knowledge of isotopic ratio sets constrains on the origin and évolution of a planetary atmosphère, understands its dynamics and describes contemporary Chemical and photochemical paths.

Moreover, contributions from isotopologues may influence the observed spectra and could possibly explain any misfits between the data and the model [START_REF] Coustenis | Titan trace gaseous composition from CIRS at the end of the Cassini-Huygens prime mission[END_REF].

The minor isotopologues exhibit extremely low abundances in Titan's atmosphère. The identification of isotopologues in Titan's infrared spectrum can be obtained only where the adéquate spectroscopic data files are available. GEISA 2009 distribution has been updated to include such line lists. In HITRAN, the minor isotopologues are embedded with the main molécule.

Eighteen new minor isotopologues hâve been added to the ARTT code, one for ethane (13C2H6), two for hydrogen cyanide (DCN and H13C14N), three for cyanoacetylene (HI3CCCN, HC13CCN and HCC13CN), five for water (HI6OD, H17OH, Hl7OD, H,8OH and Hi8OD) and seven for carbon dioxide (l3C02, l2C180160, 12C170I60, 13C170I60, 13CI80I60, l2C180l70 and l2C1802). The spectroscopic parameters of ail these molécules hâve been adopted from GEISA 2009, except for HI3C14N and cyanoacetylene. The HI3CI4N line lists hâve been adopted by HITRAN 2008, while cyanoacetylene minor isotopologues hâve been picked up from the CIRS team website. In fact, they are Jolly et al. (2007) data, modified by C. Nixon in GEISA format.

Some of these minor isotopologues hâve been detected in Titan's atmosphère. The fïrst identification of 13C-ethane (13C2H6) V12 band was derived from CIRS spectra (Nixon et al., 2008a). H13C14N has been detected in Titan from ground -based observations of IRAM/JCMT data in May-June 1995 [START_REF] Hidayat | Millimeter and Submillimeter Heterodyne Observations of Titan: Retrieval of the Vertical Profile of HCN and the 12C/13C Ratio[END_REF] and from Submillimeter Array data in February 2004 (Gurwell, 2004). The same molécule has been inverted from CIRS limb data at 83°N [START_REF] Vinatier | The Titan l4N/15N and 12C/13C isotopic ratios in HCN from Cassini/CIRS[END_REF].

The existence of three cyanoacetylene minor isotopologues has been studied by [START_REF] Jennings | Isotopic Ratios in Titan's Atmosphère from Cassini CIRS Limb Sounding: HC3N in the North[END_REF] in CIRS spectra. The H13CCCN has identified for the first time at 658.7

cm'1 but it has been impossible to detect the other two, HC13CCN and HCC13CN at 663 cm'1 because they are blended with the stronger V5 band of the main isotopologues at 663.3 cm'1 [START_REF] Jennings | Isotopic Ratios in Titan's Atmosphère from Cassini CIRS Limb Sounding: HC3N in the North[END_REF].

As far as the carbon dioxide minor isotopologues are concemed, [START_REF] Nixon | Isotopic Ratios in Titan's Atmosphère from Cassini CIRS Limb Sounding: C02 at Low and Midlatitudes[END_REF] reported a 6.5o détection of 13CC>2 and a possible 3.5o détection of 13C180160 at 648.75 and 662.5 respectively using CIRS data. The détection of the latter is noted as possible due to need for higher resolution spectroscopic data [START_REF] Nixon | Isotopic Ratios in Titan's Atmosphère from Cassini CIRS Limb Sounding: C02 at Low and Midlatitudes[END_REF].

It should be noted that although the 13C-isotopologue of diacetylene has been detected in Titan's atmosphère at 627.9 cm'1 [START_REF] Jolly | The v8 Bending Mode of Diacetylene: From Laboratory Spectroscopy to the Détection of 13C Isotopologues in Titan's Atmosphère[END_REF], it has not been included in the current ARTT upgrade. GEISA 2009 lacks the line lists of this isotopologues. From the new molécules in ARTT, DCN, water minor isotopologues, 12C180160, 12C170160, 13C170160, 12C180170 and 12C18C>2have not been detected yet in Titan's infrared spectrum. The Table 3.4

below lists these molécules and their sources.

In the following Figs. 3.11 -3.20,1 hâve plotted the output of the modeled émission features of ail the molécules included in the new upgraded radiative transfer software. Figure 3. 12 -Output of the model simulation of the FP3 but with zoom in the spectral région from 642 to 665.9 cm'1. The 13-C isotopologues of carbon dioxide are located in this spectral région as well as the ones of cyanoacetylene. 13C180160 has émission band at 643 cm'1, 13C170160 at 646 cm'1, 13C1602 at 648 cm'1, 12C1802 at 657 cm'1, 12C180170 at 660 cm'1, 12C180160 at 662.5 cm'1 and 12C170160 at 664.9 cm'1. 12C180160, 12C170160, 13C170160, 12C180170 and 12C1802 hâve not been detected yet in Titan's infrared spectrum. Only H13C12C2N has been detected at 658 cm'1 [START_REF] Jennings | Isotopic Ratios in Titan's Atmosphère from Cassini CIRS Limb Sounding: HC3N in the North[END_REF]. The rest two are blended with the Q-branch of HC3N at 663 cm'1 and higher resolution than CIRS is needed to separate them. ARTT can identify them as we can see in Fig. 3.13, but only with using much higher resolution for its calculations. cyanoacetylene, which are blended with HC3N at 663 cm'1. These calculations hâve been made by using a resolution of 0.001 cm'1, two orders of magnitude higher than the CIRS one (0.5 cm'1). The signatures of these isotopologues are indeed blended, but they can be identified by the other bands. The HC3N has not been plotted.

HCC13CN has émission bands at 650, 652. exists at 674 cm'1. In this région, there is also the signature of C2HD at 674 cm'1 and 13-C HCN isotopologues at 706 cm'1. This sélection contains 160 spectra recorded from the T68 Cassini flyby. We had firstly applied the relative température profile as retrieved from the corresponding FP4 data and then, through an itérative process, we had been running the code until we reached the best fit.

This procedure had been done for both ARTT versions. Fig. From the différence plot (Fig. 3.21-b) we can see that the old version has higher émission from 670 to 870 cm'1 (négative différences). On the other hand, from 940 to 1000 cm'1, the old model has lower émission (positive différences). Indeed, from 770 to 680 cm'1 both models seem identical, but from 660 to 700 cm'1 the old version has a slightly higher émission as well as from 700 to 720 cm'1 and from 720 to 750 cm'1. The émission of the old model becomes higher in the région of 740 to 800 cm'1 and continues similarly from 760 to 870cm'1.

Eventually, from 880 to 1100 cm'1 both fits seem identical. Judging from the différences between each model and the data, the new version fîts better the data at the range between 670 to 870 cm'1, while in the rest wavenumbers the resuit is the same. In this point, we hâve to note that we hâve used the same continuum fit with the same adjustments. We will discuss the reason of the différences between ARTT 1.0.0 and ARTT 0.3.8 below. The major molécules in this spectral région (CIRS/FP3) are: HCN, C2H2, C3H8, C2H6, C2H4 and CO2. We examine these molécules one by one. Table 3.5 summarizes the testresults of the two ARTT versions.

HCN:

Hydrogen cyanide is one of the dominant trace gaseous species in the spectral range of 650 to 780 cm'1. However, the line lists for this molécule are from the same database distribution The next plot (Fig. 3.24) shows the updates of C2H2 spectroscopic dataset in GEISA 2009.

ISpectroscopic data recently measured and added in GEISA-09

. 'v Work in progress for line intensifies measurements The C2H2 upgrade in GEISA 2009 contained new lines, but not in the région of FP3, therefore both ARTT versions produce the same signature in FP3 spectral simulation. Indeed, the différence between the two models in Fig. 3.19 is close to zéro for acetylene. Any déviance from the zéro level can be explained by the contribution of the HCN émission.

C3H8:

Both models hâve high négative différence within the région between 700 and 800 cm'1 (Fig. band are incorrectly scaled (Nixon et al., 2009). When comparing the abundances retrieved by both code versions, the relation between the 0.3.8/1.0.0 is 0.420 as described above. The adoption of the new spectroscopic file for propane in ARTT 1.0.0 fixes this malfunction.

C2H6:

Both models hâve notable négative différence within the région between 800 and 880 cm'1 . The lines lists used in ARTT 1.0.0 hâve been separated into two files. One was adopted from Vander-Auwera et al. ( 2007) at 1350-1496 cm'1 (FP4-region) and the other from GEISA 2009 at 613-843 cm'1 (FP3-region). For the FP3 part, the différences are notable.

The same fit is achieved with a value 7.5% higher in 1.0.0 than in 0. The différences in the région of C2H4 (900-1000 cm'1) are indeed significant. Contrary to previous results, the ARTT new édition generates simulations with higher émission than the older version. The previous ARTT version used spectroscopic lines from [START_REF] Blass | Absolute intensities in the v7 band of ethylene: tunable laser measurements used to calibrate FTS broadband spectra[END_REF] and Rotger (2006pers. comm.-Rotger et al., 2008) In CIRS FP3 spectra, the émission of carbon dioxide at 667 cm'1 has no significant différence when using the previous and the new . However, when fitting the same spectral sélection by using both codes, the différence in abundance is at about 2% higher values for 1.0.0. This can be explained by the increase of the number of the transition lines of C02 in the new GEISA édition (Fig. 3.25). The following plot (Fig. between the two versions.

The work described hereabove on the comparisons among the code versions, serves the purpose of helping adjusting new inferences to past ones, without having to redo ail the calculations and also understanding any discrepancies found with previous results, while identifying significant and real variations.

Conclusions

The new radiative transfer code upgrade functions in an enhanced and optimized way and allows for an improved fit to the data, thanks to the updated haze and spectroscopic parameters. Furthermore, the previous version generated the synthetic spectra by taking into account 21 molécules. In the upgraded édition, ARTT can search for 18 additional molécules and isotopologues providedfrom the latest available spectroscopic databases with more précisé molecular parameters. The molecular inventory of the code has been almost

doubled. This allows us first of ail improve the fit with molécules such as those that hâve been recently detected by Nixon et al., (2008a;[START_REF] Raulin | Planetary science: Organic lakes on Titan[END_REF] and [START_REF] Jennings | Isotopic Ratios in Titan's Atmosphère from Cassini CIRS Limb Sounding: HC3N in the North[END_REF]. Secondly, the research for the presence of more complex and new weak species, as will be discussed in Chapter 4. Moreover, by using large spectral averages from CIRS data, the identification of more minor isotopologues is easier.

Some problems of the previous version hâve been fixed, such as the propane V26 line strengths wftich are corrected. Additionally, the continuum is better fitted due to the more accurate aérosol distribution, which has been adopted. The computational time stays within acceptable levels considering the amount of calculations, which the software performs during each simulation.

Chapter 4

Titan's stratosphère by Cassini/CIRS

The Cassini-Huygens mission in the Kronian System has revealed a complex world in its largest satellite, Titan. The stratospheric organic chemistry, resulting from several processes like photodissociation and recombination starting with the mother molécules N2 and CH4, comprises a large number of hydrocarbons and nitriles. In addition, a few oxygen compounds were detected, possibly of extemal origin. The research so far has concluded that the relations between Titan's atmospheric veil with the surface, and probably the interior, are very close.

Hence, the study of the atmospheric composition, and especially its trace gaseous molécules and their isotopologues, from the surface to the upper atmosphère, is crucial for obtaining a complété picture of the moon.

After giving the basis for my research (processes, method), I présent the results of yearly CIRS spectral averages analyses during the Cassini nominal mission, which were published in Icarus [START_REF] Coustenis | Titan trace gaseous composition from CIRS at the end of the Cassini-Huygens prime mission[END_REF]. This study provided us with the latitudinal distribution of température and composition from South to the North in Titan's stratosphère and allowed me to train in CIRS spectral processing and analysis.

In this Chapter, I also présent more time-resolved results of the Cassini/CIRS spectra analysis which I hâve performed, in order to investigate the température structure and the trace gaseous composition of Titan's stratosphère from the beginning of the mission up to early 2012. This work has been recently published in the Astrophysical Journal (Bampasidis et al., 2012a). I report in this Chapter the inferred variations in température and Chemical composition in the stratosphère during the Cassini mission, which include Titan's Northern Spring Equinox (NSE) in mid-2009.

During these studies, I also worked on the mixing ratios of benzene and HC3N, weaker molécules, hence more difficult to détermine and not yet published in their entirety. For benzene, we hâve obtained vertical distributions from the surface of Titan to the upper atmosphère by applying the 1-D photochemical model of Lavvas et al. (2008a;b). This work was presented in Faraday Discussions 147: Chemistry of the Planets, which was held in Saint Jacut de la Mer (France) from 14 to 16 June 2010 [START_REF] Coustenis | Titan trace gaseous composition from CIRS at the end of the Cassini-Huygens prime mission[END_REF].

Moreover, we hâve applied our own radaitive transfer code to analyze CIRS data in the focal plane 1 (FP1) in order to search for water vapor, following a first détection by ISO (Coustenis et al. 1998). Our own analysis verified the calculations performed with another radiative transfer code used by our collaborators at the GSCF and in Oxford. The results of this study hâve been recently published in Icarus (Cottini et al., 2012a).

Recently, we hâve also focused on changes observed in the émission bands of several components at higher northem and southem latitudes than the ones we usually explore and the évolution in time that we find indicate a seasonal cycle which affects some species (like HC3N and the haze products) more than others. This work was recently accepted for publication in the Astrophysical Journal Letters (Jennings et al., 2012, in press).

I complété this Chapter by describing some perspectives for the future where we look at the possible détection of some weak species or new molécules in large averages and also in the évolution with time of some of the most abundant species in Titan's atmosphère since the Voyager times.

Titan's stratosphère before the Cassini-Huygens mission

Cassini entered into Satum insertion orbit (SOI) on 2 July 2004 and since then it observes Titan closely through several programmed flybys. The Composite Infrared Spectrometer (CIRS) (Flasar et al., 2004) opérâtes during most of them, delivering numerous infrared spectra with spectral resolutions varying between 0.5 cm'1, 2.5 cm'1, and 15.5 cm'1. A detailed description of the instrument was given in Chapter 2 along with a description of my personal contribution to the improvement of the CIRS data processing. In order to analyze CIRS spectra, we also had to upgrade the existing radiative transfer code to produce a new powerful software version that incorporâtes ail the recent spectroscopic data (Rothman et al., 2009;Jacquinet-Husson et al., 2011) and a new aérosol model (Flasar et al., 2004;[START_REF] Vinatier | Optical constants of Titan's stratospheric aérosols in the 70-1500 cm'1 spectral range constrained by Cassini/CIRS observations[END_REF]. The procedure and the results of this upgrade were described in Chapter 3. Before now describing the data I hâve used in my work here, I will briefly discuss the pre-Cassini knowledge of Titan's stratospheric thermal structure and composition as relevant to my research.

Voyager 1 & 2 encounters

Infrared spectroscopy has provided the characterization of Titan's température and Chemical structure. The first in situ measurements in the thermal radiation of Titan were obtained during the Voyager 1 encounter, in November 1980, right after Titan's northem spring equinox. The Infrared Spectrometer Radiometer (IRIS) on board the spacecraft [START_REF] Hanel | Infrared spectrometer for Voyager[END_REF] performed disk-resolved observations of Titan's stratosphère at different latitudes and longitudes with a spectral resolution of 4.3 cm'1 in the 200 and 1500 cm'1 région.

The results of the VI/IRIS spectral analysis provided the température of the stratosphère as well as its composition. Molecular nitrogen was determined to be the dominant species, followed by methane in a few percent mole fraction [START_REF] Broadfoot | Extrême ultraviolet observations from Voyager 1 encounter with Satum[END_REF][START_REF] Hanel | Infrared observations of the Satumian System from Voyager 1[END_REF][START_REF] Tyler | Radio science investigations of the Satum System with Voyager 1 -Preliminary results[END_REF]. Voyager also detected a rich inventory of trace gases in the atmosphère consisting of hydrocarbons and nitriles. Indeed, acetylene, ethylene, ethane and hydrogen cyanide were identified through their émission spectral signatures [START_REF] Hanel | Infrared observations of the Satumian System from Voyager 1[END_REF] . Moreover, methyl acetylene and propane were tentatively identified [START_REF] Maguire | C3FL and C3H4 in Titan's atmosphère[END_REF] and confirmed since then, while diacetylene, cyanoacetylene and cyanogen were also found in the stratosphère [START_REF] Kunde | C4H2, HC3N and C2N2 in Titan's atmosphère[END_REF]. The existence of mono-deuterated methane was reported a little later [START_REF] Kim | The abundance of CH3D in the atmosphère of Titan, derived from 8-to 14micron thermal émission[END_REF], while the first oxygen-bearing compound (CO2) was also identified in the Vl/IRIS spectra a year later [START_REF] Samuelson | CO2 on Titan[END_REF].

Equatorial vertical atmospheric refractivity profiles were yield by Voyager 1 radio occultation measurements (Radio Science System-RSS) from the surface up to 200 km. These measurements were used to retrieve vertical thermal profiles [START_REF] Lindal | The atmosphère of Titan: An analysis of the Voyager 1 radio occultation measurements[END_REF]. These température profiles depended on the atmospheric composition and a pure nitrogen atmosphère was assumed at the time. [START_REF] Lellouch | Titan's atmosphère and hypothesized océan: A reanalysis of the Voyager 1 radio-occultation and IRIS 7.7-pm data[END_REF] reanalyzed the Vl/RSS measurements and updated the thermal profiles by using a more précisé atmospheric composition, including methane and various Argon molecular fractions.

A thorough analysis of IRIS spectra provided the vertical and spatial variations of the température and composition from 53°S to 70°N in 1980. The température profiles used in these studies were retrieved from the best fit of the v4 methane émission band at different latitudes (Coustenis et al., 1989a;[START_REF] Coustenis | Titan's atmosphère from Voyager infrared observations. II -The CH3D abundance and D/H ratio from the 900-1200/cm spectral région[END_REF]1991;[START_REF] Coustenis | Titan's atmosphère from Voyager infrared observations. IV: Latitudinal variations of température and composition[END_REF]. The warmest région was found to be close to the equator (5-7°S). A température decrease of about 17 K at 0.4 mbar (225 km) was observed between the 5°S and 70°N, while a shorter decrease of about 3 K was observed at the south hemisphere from 5°S to 53°S [START_REF] Flasar | Titan's stratospheric températures -A case for dynamical inertia[END_REF][START_REF] Coustenis | Titan's atmosphère from Voyager infrared observations. IV: Latitudinal variations of température and composition[END_REF].

As far as the trace gaseous concentrations are concemed, no longitudinal variation has been found by Vl/IRIS data in Titan's stratosphère between 4 and 9 mbar (Coustenis et al., 1989a). HCN increased significantly from pôle to pôle, while carbon dioxide mole fraction remained constant in respect to the latitude. The abundances of the other trace constituées presented variations from the équatorial latitudes and northwards. Ethane, acetylene and propane showed a moderate increase by a factor of two. Methyl acetylene, ethylene and diacetylene showed significant enrichments at latitudes northem than 50°N. The other nitriles, cyanogen and cyanoacetylene hâve no spectral signatures from 50°N and southwards (Fig 1 .3

in Chapter 1). The results of this research are listed in Table 1 below. The ethane mole fraction remained constant compared to the VI era as well as those of CO2, C2H4, C3H8, C3H4 and C4H2 within the error bars. On the other hand, acetylene and HCN exhibited an increase of 30% in the ISO spectra. Cyanoacetylene was observed in diskaveraged spectra for the first time at 553 cm'1. One of the most important discoveries of the ISO Short Wavelength Spectrometer (SWS) was the firm détection of two water vapor signatures near 40 micron (Coustenis et al., 1998) and of benzene in Titan's atmosphère at 674 cm'1 with a maximum émission coming from atmospheric levels between 0.2 and 20 mbar (Coustenis et al., 2003).

Ground-based observations

Ground-based observations of Titan offer significant information on the stratospheric composition as well as on the vertical distribution of molecular abundances is several cases.

After CO2, the other oxygen-bearing species, carbon monoxide (CO), was discovered on Titan by Lutz et al. (1983) with near infrared observations at the Mayall telescope of Kitt Peak National Observatory (KPNO) -USA. Acetonitrile (CH3CN) was firstly detected by [START_REF] Bezard | Détection ofAcetonitrile on Titan[END_REF] from spectra acquired at the Institut de Radioastronomie Milimetrique -IRAM -(Spain) radiotélescope. Using the same facility, [START_REF] Marten | New Millimeter Heterodyne Observations of Titan: Vertical Distributions of Nitriles HCN, HC3N, CH3CN, and the Isotopic Ratio 15N/14N in Its Atmosphère[END_REF] provided the first vertical profiles of Titan's stratospheric nitriles (HCN, HC3N and CH3CN), as also reported by [START_REF] Paubert | The millimeter spectrum of Titan -Detectability of HCN, HC3N, and CH3CN and the CO abundance[END_REF]. Moreover, the isotopic ratios of 12C/13C from HCN [START_REF] Hidayat | Millimeter and Submillimeter Heterodyne Observations of Titan: Retrieval of the Vertical Profile of HCN and the 12C/13C Ratio[END_REF], l4N/15N from HCN [START_REF] Marten | New Millimeter Heterodyne Observations of Titan: Vertical Distributions of Nitriles HCN, HC3N, CH3CN, and the Isotopic Ratio 15N/14N in Its Atmosphère[END_REF] and l60/l80 [START_REF] Gurwell | Carbon, Nitrogen And Oxygen Isotopic Ratios In Titan's Stratosphère[END_REF] were measured, but the latter hâve not been fully published yet. Millimeter observations at IRAM also delivered the vertical profile of HCN [START_REF] Paubert | First radiodetection of HCN on Titan[END_REF][START_REF] Tanguy | Stratospheric profile ofHCN on Titan from millimeter observations[END_REF].

NASA Infrared Telescope Facility (IRTF) ethane measurements [START_REF] Kostiuk | Ethane abundance on Titan[END_REF]2005;[START_REF] Livengood | NOTE: Improved Détermination of Ethane (C2H6) Abundance in Titan's Stratosphère[END_REF] at 12 pm (émission coming from 7 to 0.1 mbar) gave values comparable to the Vl/IRIS results. The first spectrally-resolved détection of the propane émission band on Titan's stratosphère was made at NASA/IRTF [START_REF] Roe | Propane on Titan[END_REF]. A south spring polar accumulation of ethylene was also shown just before solstice (Roe et al., 2004b) at Keck Observatory spectra using the Long Wavelength Spectrometer (LWS). These spectra were taken from the south polar région of Titan, which had not been observed by Voyager 1.

High-resolution spectroscopy at 3 pm provided measurements of HCN and C2H2 high in Titan's stratosphère and mesosphere and CH3D in Titan's troposphère [START_REF] Geballe | High-Resolution 3 Micron Spectroscopy of Molécules in the Mesosphere and Troposphère of Titan[END_REF][START_REF] Kim | Clouds, haze, and CH4, CH3D, HCN, and C2H2 in the atmosphère of Titan probed via 3 pm spectroscopy[END_REF]. In the Table 4.4 below, I hâve gathered these and more the ground-based observations of Titan's stratospheric trace gases published so far to the best of our knowledge.

Recently, the Heterodyne Instrument of the Far-Infrared (HIFI) on board ES A's Herschel Space Observatory (de Graauw et al., 2010;[START_REF] Pilbratt | Herschel Space Observatory. An ES A facility for far-infrared and submillimetre astronomy[END_REF] reported the first détection of hydrogen isocyanide (HNC) in Titan's atmosphère at 18.14 cm'1 [START_REF] Moreno | First détection of hydrogen isocyanide (HNC) in Titan's atmosphère[END_REF], during observations of Titan on June 14 and December 31, 2010 and provided abundances of several species [START_REF] Courtin | First results of Herschel-SPIRE observations of Titan[END_REF][START_REF] Moreno | The abundance, vertical distribution and origin of H20 in Titan's atmosphère: Herschel observations and photochemical modelling[END_REF]. [START_REF] Waite | Ion Neutral Mass Spectrometer Results from the First Flyby of Titan[END_REF], 1.6 ±0.5% in the stratosphère (Flasar et al., 2005), 1.48 ±0.09% in the lower stratosphère and 5.65 ±0.18 near the surface (Niemann et al., 2010)]. The existence of a large number of hydrocarbons and nitriles in Titan's atmosphère, as were firstly shown in Voyager 1/IRIS observations (Hanel et al., 1981;Coustenis et al., 1989a), originates through methane and nitrogen photolysis essentially and recombination with nitrogen, as predicted by [START_REF] Strobel | The Photochemistry of Hydrocarbons in the Atmosphère of Titan[END_REF].

Figure 4. 1 -The methane cycle on Titan (Atreya et al., 2006).

The products of the photochemistry, which occurs above the level of 700 km in the atmosphère, are short-lived radicals of CHn (n=l,2,3). The self-reaction between CH3 radicals forms ethane, which is the most abundant methane photolysis product. Ethane is a stable molécule and hence a key molécule in the methane cycle from the time it is formed until it précipitâtes on the surface [START_REF] Choukroun | Is Titan's shape caused by its meteorology and carbon cycle[END_REF].

The products of methane photolysis are a high variety of complex hydrocarbons and once methane is dissociated, the Chemical procedure is irréversible. These organics form the haze layer, which obscures the lower atmosphère and the surface at visible wavelengths.

Substantial quantities of these aérosols accumulate on the surface.

Part of the organic surface deposits form liquid lake-like features, discovered by Cassini, which are extremely significant in photochemical modeling and habitability. When aérosols reach the équatorial surface, the environment is dry enough for them to remain in solid State and dunes are formed. However, the methane-based hydrologie cycle has not been

well constrained yet by the Cassini-Huygens mission.

Photochemical reactions and hydrogen escape dissociate methane in the upper atmosphère [START_REF] Strobel | The Photochemistry of Hydrocarbons in the Atmosphère of Titan[END_REF]Yung et al., 1984;Wilson & Atreya, 2004;Lawas et al., 2008b) and if no source replenishes it, it will be totally depleted from the satellite within 10-100 million years [START_REF] Strobel | The Photochemistry of Hydrocarbons in the Atmosphère of Titan[END_REF]Atreya et al., 2006). Like the hydrological cycle on the Earth, this procedure is expected to be irréversible. The study of possible methane sources in Titan's environment is crucial to understand the methane cycle. The sources that could replenish the dissociated atmospheric methane in Titan's atmosphère are still unknown, but they are possibly located on the surface and in the interior (see following Chapter). Atreya et al. (2006) and other investigators hâve suggested methane outgassing from the interior and/or hydrothermal sources to satisfy the mission methane.

In my PhD studies, I hâve looked at Titan's environment from different points of view, from the atmosphère to the surface conditions and even to some of the possibly associated internai processes, as will be discussed in the next chapter. The methane cycle is the leading theme among ail these processes and helps us understand the various physical and geological processes, while by studying the processes themselves we better understand this alcanological cycle, whose importance is major in the work presented here.

Photochemistry

Gaseous species are idéal agents for tracing the photochemical and dynamical processes occurring on Titan's atmosphère. The combination of Voyager, of other space-bome and of Earth-based observations of Titan provides constraints to photochemical models.

Complex photochemistry in the upper atmosphère

The Chemical and physical pathways involved in the Chemical évolution of the atmosphère of Titan hâve been described by several photochemical models (Yung et al., 1984;Toublanc et al., 1995;Wilson & Atreya, 2004;Lawas et al., 2008a;[START_REF] Raulin | Planetary science: Organic lakes on Titan[END_REF]. These models estimate the resulting vertical concentration profiles of the different molécules starting from dissociation of N2 and CH4 and then fulfilling a volatile Chemical cycle. Solar EUV and UV photons are the main sources of photo-dissociation [START_REF] Agren | On the ionospheric structure of Titan[END_REF][START_REF] Galand | Ionization sources in Titan's deep ionosphère[END_REF]. Other photo dissociation sources are galactic cosmic rays and energetic particles from the Satumian magnetosphere (Fig. 4.2) [START_REF] Krasnopolsky | A photochemical model of Titan's atmosphère and ionosphère[END_REF]. The photochemical processes provide the pathways for the formation of ethane, acetylene and hydrogen cyanide in the higher atmospheric levels. These molécules play a key rôle: after their formation, they diffuse down to the lower atmosphère where they produce higher hydrocarbons and nitriles and perhaps aromatic compounds. Besides, methane dissociation probably also occurs in the lower stratosphère through photocatalytic processes involving acetylene and polyynes. Polyynes are the group of organics in which single and triple bonds altemate. The simplest member of this group is acetylene.

Contrary to the very short lifetime of C2H2 on Earth (Rudolph et al., 1984), acetylene expériences a seasonal variation during a Titan year (29. Wilson & Atreya (2004). The stable species and the principal reactions are plotted in bold.

Figure 4. 5 -Plot of the nitrile chemistry in Titan according to Wilson & Atreya (2004). The stable species and the principal reactions are plotted in bold. Boxes represent the ions.

As the haze particles fall through the atmosphère and grow, they become détectable with imaging Systems such as the Cassini Imagine Science Subsystem (ISS) at about 500 km altitude and are ubiquitous throughout the stratosphère (Porco et al., 2005). Since they are strong absorbers of solar UV and visible radiation, they heat the stratosphère and drive wind Systems in the middle atmosphère, like the ozone layer in the Earth's middle atmosphère. variability of Titan's atmosphère and its photochemical haze (Lavvas et al., 2008a;[START_REF] Raulin | Planetary science: Organic lakes on Titan[END_REF] during his PhD (A. Coustenis, co-advisor) which he updated and upscaled in the past years.

The model consists of three parts: the radiative transfer calculations, the chemistry sub-model and the microphysics sub-model (Fig. 4.6). The model firstly incorporâtes detailed radiation transfer calculations which describe the short and long wave fluxes. These calculations provide the vertical radiation field and the thermal structure which are used for simulating photochemical processes in the atmosphère.

The photolysis of Titan's major compounds (N2, CH4) produces the haze particles in its atmosphère. The Chemical sub-model solves the time-depended continuity équation for the profiles of the contributing molécules' abundances. The haze pathways lead to haze monomer production at each altitude.

The growth of the haze particles from an initial monomer is described by the microphysical branch of the model. Haze particles are assumed to be spherical with no fractal structure and the haze particle production rate is quite different from the Gaussian. The model generates the vertical haze distribution from the vertical production rates. Then, the calculated aérosol and gas opacities are included in the radiation transfer calculations in order to find their influence in the température profile and géométrie albedo. The model thermal structure dépends on the haze vertical distribution. This itérative process occurs until the equilibrium is reached.

By applying this model to Titan's atmospheric envelope we dérivé a better understanding of the unknown connection mechanism between the production of neutral gases and their transformation to haze particles. The model reproduces well the observations (Lawas et al., 2008b) as we hâve also reported in [START_REF] Coustenis | Titan trace gaseous composition from CIRS at the end of the Cassini-Huygens prime mission[END_REF].

4.4

Composite Infrared Spectrometer (CIRS) data sélection My work in studying Titan's stratosphère is twofold. I first studied the composition and température as retrieved from CIRS spectra averaged yearly during the nominal mission (2004)(2005)(2006)(2007)(2008) and then looked at the évolution of the thermal structure and Chemical composition from the beginning of the mission in 2004 up to early 2012 with more timeresolved sélections. Large averages hâve been also used for the study of water vapor and the weak species and of course to search for new molécules as I describe hereafter.

For ail selected CIRS datasets, I hâve restricted the émission angles to be less than 65°. I sum ail longitudes (because it was demonstrated that no longitudinal variations are expected (Coustenis et al., 2007;Teanby et al., 2008) and retrieve spectra only for different latitudinal bins.

As discussed in Coustenis et al. (2007), we need to correct the data for any latitude smearing effect observed for high émission angles and/or data taken at latitudes higher than 50°N (Fig. 4.7). The sounded latitudes can be about 5°lower when the line-of-sight intercepts the surface at 55°N or 55°S. Thus, the actual latitude of the stratospheric altitudes observed being rather 50°N and 50°S, respectively. The latitudes listed in data Tables below (4.4-4.14) are corrected for this effect. The contribution functions peak at about 150 km (provided by Dr. N. Teanby). This plot shows that the actual location at which CIRS data are acquired dépends highly on the émission angle of the observation. However, in this work we hâve taken into account this so-called smearing effect (Coustenis et al. 2007) and corrected the line of sight before the retrievals. (2004)(2005)(2006)(2007)(2008) In our paper [START_REF] Coustenis | Titan trace gaseous composition from CIRS at the end of the Cassini-Huygens prime mission[END_REF], we gathered nadir CIRS spectra from the first 44

Cassini flybys which hâve been performed within the first 4 years of the Cassini-Huygens mission. It was the first time that Titan was explored continually and that study is complementary to the previous one described in Coustenis et al. (2007). The spectra analyzed in the 2007 study were obtained within the first two years of the mission (July 2004-January 2006, flybys TB-T10). A full description of the instrument and the processing and calibration procedure can be found in Chapter 2 of this manuscript.

As in Coustenis et al. (2007), in the 2010 study we averaged CIRS FP3 and FP4 spectra from the CIRSDATA database in latitudinal bins of 10°ranging from 85°S to 80°N in both mid and high resolutions (2.54 and 0.53 cm"1 respectively). We hâve excluded the T0

and TA flybys from our queries. Data from T0, the first Cassini flyby of Titan, was omitted since they were obtained at distances too high to be compared to the observations which followed. Most of the data obtained from TA flyby were corrupted due to a CIRS command error. Tables 4.5 and 4.6 list the high and medium resolution spectra that I assembled and used in Coustenis et al. (2010b -c.f. Table 1). In the more recent study, in 2012, we looked at more time-resolved data now that a large number of spectra hâve become available. We hâve thus used ail available and exploitable data at high spectral resolution (0.5 cm'1) and I performed averages of the spectra acquired during one or several flybys as necessary to attain a high signal-to-noise ratio for my calculations. To do this, I hâve queried for spectra averaged over a given date (flyby, or over a month) and attributable to a given latitudinal bin between (50°S and 50°N) containing a large number of spectra in general, with some exceptions (Bampasidis et al., 2012a).

The TB-T9 flybys (December 2004-December 2005) Tables 4.7 and 4.9 contain nadir averages at high resolution (0.5 cm'1) for the northem hemisphere (FP3 and FP4 spectra respectively), while Tables 4.8 and 4.10 contain the datasets for the southem hemisphere (FP3 and FP4 spectra respectively). The sélections we made in order to enhance the signal-to-noise ratio cover different latitudes on Titan. Each query was restricted to émission angles lower than 65°. The average signal-to-noise ratio is given in the Tables, as well as the corresponding Cassini Titan fly-by(s) and the relative solar longitude (Ls). Foliowing the same sélection criteria I hâve also analyzed 6 datasets of medium resolution spectra (2.54 cm'1) taken at FP3 and FP4 spectral ranges for January 2008 (flyby T40), April 2009 (flybys T52-T53) and May 2010 (T68). Table 4.11 and 4.12 list these midresolution data characteristics acquired at FP3 and FP4 respectively. In order to detect and retrieve the abundances for water vapor, I hâve used our radiative transfer code (ARTT) to simulate large CIRS averages in the FP1. These calculations offer a chance to compare the results of our radiative transfer code with the retrievals of the NEMESIS code -Non-linear Optimal Estimator for Multivariate Spectral Analysis (Irwin et al., 2008).

The water vapor détection in Titan's atmosphère can be achieved only through large spectral averages, because, as demonstrated by the ISO first détection in 1997 data (Coustenis et al., 1998) the émission signatures are rather weak. In our paper Cottini et al. (2012a), we queried for 3 FP1 sélections within the season of Titan's northem winter which is from December 2004 up to December 2008. These FP1 datasets are listed in Table 4.13 below. The method for the analysis of the spectra I hâve analyzed the Cassini/CIRS spectra using a method described in previous articles (Coustenis et al. 2007;2010). In brief, we model Titan's thermal infrared spectrum using a line-by-line monochromatic radiative transfer code upgraded ffom the one we used in previous Titan atmospheric structure retrievals (ARTT, for Atmospheric Radiative Transfer for Titan, c.f. Bampasidis et al. 2012 and references within).

I computed synthetic nadir spectra in each of the FP1 (10-600 cm'1), FP3 (600-1100 cm'1) and FP4 (1100-1400 cm'1) spectral ranges, which I then compare to the observations in order to retrieve the vertical température profile (from the V4 methane band émission in the FP4, at 7.7 micron) and the composition of Titan's atmosphère (ffom the gaseous molecular signatures in FP1 and FP3). The dominant émission bands in the CIRS/FP4 spectra (1100-1500 cm'1) are the methane V4 band at 1304 cm'1 and the CH3D V6 band at 1156 cm'1. Several molécules hâve émission bands in this région, but their signature has not been verified yet, like the propane band V24 at 1472 cm'1 (Nixon et al., 2009). The 13-C bearing isotope of CH3D has also a weak émission at 1148 cm'1. As shown in the contribution function of CH3D (Fig. 4.8), the atmospheric région which is probed in this band at 50°S is between 0.03 and 10 mbar with a peak at 0.24 mbar. Similarly, the Q-branch of CH4 peaks at 0.29 mbar with a FWHM between 0.04 and 1 mbar. As far as the methane wings are concemed, the R-and P-branches probe the atmosphère between 0.01-7 mbar and 0.1 -5 mbar respectively.

The best fit for the sélection of May 2009 at 50°S in the FP4 région is illustrated in Wcm'"sr' /cm' . When other sources of uncertainties (in the methane abundance, the thermal structure, the calibration and the continuum fit) are included, the total error is about 2x10'9

Wcm^sr'Vcm'1. The misfit between the model and data from 1180 to 1250 cm'1 has not been yet understood; perhaps is caused by uncertainties in methane line lists [START_REF] Coustenis | Titan trace gaseous composition from CIRS at the end of the Cassini-Huygens prime mission[END_REF].

"I-1-1-|-1-1-r
Data-Model The way the radiative transfer calculations work is that either the opacity is known and the thermal composition inferred or the température profile is injected in the code and the mixing ratios inferred. To get the température profiles from the V4 methane band, we use the CH4 vertical mixing ratio profile in the stratosphère as measured by the Huygens probe (Niemann et al., 2010), which yields 1.48% above the cold trap, compatible with the CIRS inferences from FP1 (Flasar et al., 2005). We also take 0.11% for H2 [START_REF] Samuelson | CO2 on Titan[END_REF] and 98.41% for N2. We then use the inversed température profiles to infer the abundances of the emitting gases.

In the ARTT code, I hâve included more constituents relevant to Titan's chemistry following a procedure which is described in detail in Chapter 3. The contributing molécules to the synthetic model spectra are therefore 26 in total, which is ail the so-far detected hydrocarbons, nitriles and oxygen compounds in addition to the main molécules N2, H2, CH4

and Argon.

Furthermore, the updated list of ARTT molécules includes the 13CH4, 13CH3D, 12C13CH2 and H15CN isotopologues that are essential in fitting the émission bands of methane, acetylene and hydrogen cyanide respectively and allow us to infer the 12C/13C and 14N/15N ratios. Finally, new oxygen and hydrocarbon isotopologues that were recently detected [START_REF] Jennings | Isotopic Ratios in Titan's Atmosphère from Cassini CIRS Limb Sounding: HC3N in the North[END_REF]Nixon et al., 2008a) are also incorporated.

As said in Chapter 3, the need for a massive upgrade was obvious since new spectroscopic parameters for the molécules found in the FP3 and FP4 spectral ranges became available from recent laboratory studies, as for instance in the case of ethane and propane spectroscopic parameters, both of which were provided from the individual line lists of Vander Auwera et al. ( 2007) and Flaud et al. (Flaud et al., 2010) respectively. These new line lists are from the GEISA 2009 (Jacquinet- Husson et al., 2011) and the HITRAN 2008 (Rothman et al., 2009) databases, in which ail these updates hâve been included.

More in detail, the analysis we perform comprises the following steps: we begin by inferring the température profiles by the inversion of the observed ru methane émission band at 1305 cm'1 in CIRS/FP4 individually for each sélection (Tables 4. 5-6, 4.9-10 and 4.12-14).

These thermal profiles are provided by R. Achterberg of the GSFC following the method described extensively in Achterberg et al. (2008;2011). Once the vertical thermal structure is known, the contribution fonction of each molécule incorporated in the model can be calculated (thus providing the altitude/pressure range from where the maximum of the émission originates) and more importantly its abundance can be retrieved when the best fit is achieved.

The a priori reference température profile for these calculations was the 15°S profile from Flasar et al. (2005). As I hâve described in Chapter 3, Section 3.2,1 hâve also run tests using the HASI measured température structure in the stratosphère and troposphère (Fulchignoni et al., 2005). At the équatorial latitudes, the impact of these thermal profiles in the abundances retrievals is at about 10% less for ail molecular abundances, except for HCN (20%). This small impact is expected because the HASI profile was inferred above the Huygens landing site (10°S). In the northem latitudes, the divergence is more pronounced (within 20% of increase in abundance) for the molécules with the stronger émission bands (C2H2, C2H6 and HCN). Most of the molecular abundances vary by less than 10% whatever the initial température profile (see Tables 3.2 and 3.3).

Then, I adapt each individual température profile to solve the radiative transfer équation in the FP3 (and FP1 for the water vapor) part of the CIRS spectrum. Through an itérative best-fit process, I dérivé the trace gases abundances of the various components in each sélection of data (Tables 4.5- 6,4.7-8, 4.11 and 4.13-14).

For ail of the molécules analyzed here in surface-intercepting (nadir) geometry conditions, I use only constant-with-height abundance profiles above the condensation level relevant to stratospheric levels of 0.1-20 mbar essentially (roughly 80-280 km). Only C2H2, HCN and C2H6 can provide some vertical information and there I hâve tested vertical profiles against the observations. The others do not hâve strong enough émission bands to provide the retrieval of vertically-dependent information.

Moreover, in order to fit the continuum observed in the spectra, I hâve adapted the haze description with wavenumber recently reported in [START_REF] Vinatier | Optical constants of Titan's stratospheric aérosols in the 70-1500 cm'1 spectral range constrained by Cassini/CIRS observations[END_REF] adjusted to match the level of radiance observed between the molecular bands in each sélection. This distribution provides the aérosol reffactive index from CIRS spectra in the far and midinfrared régions near 15°S and 20°S.

4.6

Titan's stratospheric composition during the nominal Cassini-

Huygens mission

The study of CIRS data obtained during the Cassini-Huygens mission has significantly advanced our perspective of Titan's atmospheric Chemical structure. In the paper of Coustenis et al. (2007), the authors hâve analyzed the CIRS spectra retrieved ffom flybys TB to T10 [START_REF] Coustenis | Titan's atmosphère from Voyager infrared observations. IV: Latitudinal variations of température and composition[END_REF], but the enhancement was greater (Fig. 4.11).

Figure 4. 11 -Latitudinal distribution of major organic trace gases from Voyager [START_REF] Coustenis | Titan's atmosphère from Voyager infrared observations. IV: Latitudinal variations of température and composition[END_REF] and Cassini (Coustenis et al., 2007), adapted from [START_REF] Coustenis | Titan trace gaseous composition from CIRS at the end of the Cassini-Huygens prime mission[END_REF].

However, due to the Cassini position during the first 10 flyby of Titan, the most northem latitude observed was the 60°N one. Since there is a significant enhancement of the trace gases mixing ratios polewards to the north, further analysis of polar data is needed.

Thanks to the Cassini mission, a wealth of spectra became available in order to study Titan at ail latitudes.

In the paper of [START_REF] Coustenis | Titan trace gaseous composition from CIRS at the end of the Cassini-Huygens prime mission[END_REF], which I hâve co-authored, we hâve used large spectral averages from flybys TB to T44 that cover the whole nominal mission of the Cassini orbiter. These averages provide more detailed abundances retrievals as well as improved méridional variations as described hereafter.

4.6.1

Titan's stratospheric chemistry retrieval

In [START_REF] Coustenis | Titan trace gaseous composition from CIRS at the end of the Cassini-Huygens prime mission[END_REF], we hâve divided Titan's globe into latitudinal bins of 10°from 90°S to 90°N with émission angles less than 50°and searched to détermine spatial variations in température and composition from 2004 to 2008, more refined with respect to previous studies. No longitudinal variations existing in the Chemical composition of the stratosphère (Coustenis et al. 2007) and the longitudinal variations in température not being significant (Achterberg et al., 2008), we summed up ail longitudes.

This work was a great introduction for me in CIRS data analysis by using ARTT (Atmospheric Radiative Transfer code for Titan) and the ensuing interprétation. It was also important that through that project I got familiar with the radiative transfer code modifications and began to interact with the CIRS data acquisition and calibration team in Goddard.

The modeling procedure is described in Section 4.4 and the data used are listed in Tables 4.5 and 4.6. Due to the détection of C2HD (Coustenis et al., 2008) and 13CH3D [START_REF] Bezard | Détection of 13CH3D on Titan[END_REF], these molécules hâve been included in the 2010 version of ARTT.

The vertical distributions of the molécules are assumed to be constant with altitude with the exception of acetylene, for which we tried some vertical profiles. We hâve also used distributions predicted from General Circulation Models (Rannou et al., 2005). Contrary to the previously described analysis of the temporal and spatial variations of the Titan's stratosphère, the GCM distributions for C2H2 work well [START_REF] Coustenis | Titan trace gaseous composition from CIRS at the end of the Cassini-Huygens prime mission[END_REF]). We confirm the méridional variations reported in Coustenis et al. (2007) with a significant enhancement at northem latitudes with new smaller uncertainties. Figure 4.12 illustrâtes the latitudinal variations during Cassini prime mission.

Différences in the mixing ratios hâve been found for ethane and for other species (C2H4 and C3H8) compared to the previous work (Coustenis et al. 2007). The decrease in the abundance near the North pôle of C2H4 and C3H8 is not confirmed. Instead, the propane's mixing ratio increases slightly, while the values for ethylene are higher.

Ethane, acetylene, propane and carbon dioxide show abundance increase from South to North by a factor ranging from 1.5 to 2. The mixing ratios of benzene and cyanoacetylene increase even more significantly to the North by factors of 30 and 150

respectively. The other molécules hâve enriched values by factors within 3 and 8 [START_REF] Coustenis | Titan trace gaseous composition from CIRS at the end of the Cassini-Huygens prime mission[END_REF]. In the 2007 study, the enrichments found in the North were not at the levels of the Voyager ones. The ones in the 2010 study are lower with respect to the ones observed at the time of the Voyager encounter.

The results of our work are compatible with the ones of Teanby et al. (Teanby et al., 2009a) and with the isotopic ratio retrievals of [START_REF] Bezard | Détection of 13CH3D on Titan[END_REF], Nixon et al. (2008a;[START_REF] Raulin | Planetary science: Organic lakes on Titan[END_REF] and [START_REF] Vinatier | The Titan l4N/15N and 12C/13C isotopic ratios in HCN from Cassini/CIRS[END_REF]. Additionally, our findings are in good agreement with photochemical models (Lavvas et al. 2008a;[START_REF] Raulin | Planetary science: Organic lakes on Titan[END_REF] and GCM models (Crespin et al., 2008).

The main reason of the observed enhancement from the Southern latitude to the northern ones seems to be the stratosphère which is in the shadow of Titan 's pôle during winter. During the polar night, the mesosphere is illuminated by the sunlight and Saturnian magnetospheric energetic électrons. Thus, HCN and other nitriles are formed in the upper atmospheric layers, they diffuse downwards and accumilate close to the polar région [START_REF] Yung | An update of nitrile photochemistry on Titan[END_REF]Coustenis étal., 1989a).

Temporal and spatial variations during the Cassini-Huygens mission

Latitudinal variations in Titan's stratospheric thermal and Chemical structure hâve been previously reported from Cassini CIRS limb and nadir data analyses (Flasar et al., 2005;Coustenis et al., 2007;2008;2010b;de Kok et al., 2007a;Teanby et al., 2007;2008;2009b;Vinatier et al., 2007a;2010b;[START_REF] Nixon | Isotopic Ratios in Titan's Atmosphère from Cassini CIRS Limb Sounding: C02 at Low and Midlatitudes[END_REF].

In this section, I investigate the possible thermal and Chemical évolution within 2006 and 2012, complementing and refining previous studies (Teanby et al. 2008;2010;Achterberg et al. 2011) by adding temporal resolution. I focus on Titan's stratosphère and lower mesosphere between -120 and 500 km in altitude and latitudes from 50°S to 50°N.

Thermal évolution of Titan's stratosphère during the Cassini mission

The température profiles used in this study were retrieved by Dr. R. Achterberg by fitting of the iu methane band at 1305 cm"1 for the 2006-2012 CIRS data acquisition period (see Fig.

above

). These retrievals were obtained by a constrained algorithm which is an inverse approach of the radiative transfer équation solution (Achterberg et al. 2008;2011). A more general view of solving inversely the radiative transfer équation is presented in Appendix B of this manuscript.

The a priori reference profile was the one published in Flasar et al. (2005) and In the northem hemisphere, the température decreases in the mesosphere, which is more sensitive to the seasonal insolation variations. This decrease starts from the stratopause level above 0.1 mbar, where the température significantly decreases from 2007 to 2010. Since the beginning of the Cassini mission, the warm northem lower mesosphere has cooled by about 14 K at 50°N (Figure 4.13a) and by about 8 K at 30°N (Figure 4.13b) at around 250 to 300 km in altitude, which was also suggested by Achterberg et al. (2011).

In the mid-stratosphere (around 200 km), the northem polar région initially began The latter phenomenon is confirmée! from CIRS limb data analysis (Figure 4.13f). We retrieved the température profiles at 50°S, 1°N and 51°N from CIRS limb observations acquired in June 2010 with the FP4 at 15.5 cm'1 spectral resolution. The température uncertainty is about 1 to 2 K. Figure 4.13f shows the température profiles retrieved from June 2010 compared with température profiles retrieved using CIRS limb spectra acquired in the 2006-2007 period with a 0.5 cm'1 spectral resolution [START_REF] Vinatier | Analysis of Cassini/CIRS limb spectra of Titan acquired during the nominal mission. I. Hydrocarbons, nitriles and CO2 vertical mixing ratio profiles[END_REF]). In the lower mesosphere, between 250-300 km, the température is lower by 10 K at 50°S and 3 K at 0°N in 2010 than in the 2006-2007 period, while at 50°N the 2010 température is about 3 K higher than the 2006 value.

In the stratosphère, below ~250 km, from 2006 to 2010, the température at 50°N increased by about 10 K, shows no significant variation at the equator and has increased at 50°S more than 10 K. The différence between the 2007 and 2010 température profiles from limb data furthermore indicates not only a significant cooling (by more than 20 K in the 300-400 km range), but also a trend for the température to become more vertically uniform and the pronounced mesosphere to be dampened, as we find from the nadir data.

These changes at northem latitudes suggest a weakening of the descending branch of the middle atmosphère méridional circulation (Achterberg et al. 2011), which implies less adiabatic heating and hence lower températures. Moreover, the changes in the circulation affect the distribution of aérosols [START_REF] Vinatier | Optical constants of Titan's stratospheric aérosols in the 70-1500 cm'1 spectral range constrained by Cassini/CIRS observations[END_REF], which in tum affect the thermal structure.

Chemical composition variation during the Cassini-Huygens mission

After fixing the température profile, I inject it in our radiative transfer code, in order to dérivé the mixing ratios of the contributing molécules from their émission bands in the CIRS FP3 and FP4 spectra through an itérative procedure. When the best fit between the model simulation and the observation is accomplished, I infer the corresponding abundances (see also [START_REF] Bampasidis | Thermal and Chemical structure variations in Titan's during the Cassini mission[END_REF][START_REF] Coustenis | Titan trace gaseous composition from CIRS at the end of the Cassini-Huygens prime mission[END_REF]). For each sélection, I match the continuum level between the émission signatures by adapting the new model of haze distribution of [START_REF] Vinatier | Optical constants of Titan's stratospheric aérosols in the 70-1500 cm'1 spectral range constrained by Cassini/CIRS observations[END_REF].

The émission bands of Titan's hydrocarbon trace gases, nitriles and CO2 are contained in the 600-1100 cm'1 région of the spectrum. These molécules (C2H2, C2H4, C2H6, C3H4, C3H8, C4H2, HCN, CO2, HC3N) hâve been previously detected in Voyager/IRIS spectra, ground based data and/or ISO/SWS recordings. The presence of ail of the above-mentioned molécules has been verified in Cassini/CIRS FP3 spectra. The list of the molécules detected in Titan's spectra, is presented in Table 4.14 below. Table 4.15 summarizes the candidate molécules that are expected to be detected in CIRS FP3 and FP4 spectra and which are discussed further in this Chapter 3, Section 3.4.6, when we describe our search for weak molécules. Ail these molécules are included in the model. The continuous probing of Titan's stratosphère by CIRS provides information for the évolution of the Chemical composition. Towards this end, I hâve calculated the abundances of the trace gases from the beginning of the mission up to early 2012.

In order to define the stratospheric régions probed by the émission observed in the different gaseous bands, we calculate the contribution functions C(z,n) by using the formula The dominant bands in the FP3 région are those of acetylene, hydrogen cyanide and ethane at 729.5 and 712.5 and 822 cm'1 respectively. For achieving the best fit, after fitting the continuum, we match these bands first for each sélection. The following show the resuit of the best-fit process for some of the aforementioned sélections throughout the FP3 région. errors due to instrumental noise itself at the location of the émission peak, the uncertainty in the placement of the continuum above which the émission rises and the uncertainty on the thermal profile and its influence on the line formation région (Coustenis et al., 1995;2010b).

(see Chapter 3, Section 3.4.2): C(z,n) = B(z)W (z,n) With W(z,
Other systematic errors are from calibration uncertainties, température profile inconsistencies in the datasets and haze description uncertainties.

The Noise Equivalent Spectral Radiance (NESR) is the signal when the instrumental signal-to-noise ratio (S/N) is unity [START_REF] Hanel | Exploration of the Solar System by Infrared Remote Sensing[END_REF]. NESR varies with the wavenumber as it shown in The ratio of NESR to the square root of the total number of averaged spectra is the l-o standard déviation. The estimation of uncertainties has included the température profile errors as described above, calibration, and line positions.

Except for the early mission dates (within 2004-2005) and sometimes in the northem sélections (e.g. February 2008), the noise level is low in our usually large averages. The errors from the instrument noise are small in température retrievals, about 0.1 K (Achterberg et al., 2008;[START_REF] Coustenis | Titan trace gaseous composition from CIRS at the end of the Cassini-Huygens prime mission[END_REF]. We only consider the relative uncertainties that mainly originated in the uncertainties of methane abundance and température (ail within 15%) and neglect the systematic errors which are the same for ail sélections.

Propane (C^Hg) and ethylene (C2FI4) show ratios that hâve error bars of about ±35% and they are not as reliable as the other inferences for the temporal or spatial variations. These molécules are affected by abundance retrieval difficulties and some insuffîcient laboratory spectroscopic data. The strong P-branch of the C2H2 band influences the retrieval of the C3H8 abundance from its émission at 748 cm'1. Around 765 and 950 cm'1, the spectrum suffers from some electrical interférence artifacts ("spikes", see Chapter 2 for further information).

The retrieved abundances at Southern latitudes (50°S and 30°S), as well as near the equator, remain rather constant-in-time during the Cassini mission within the error bars. C2H4 and C3H8 are identified as complicated molécules to process, in particular at the mid and Southern latitudes. In general, for the other molécules, the équatorial data yield higher abundances than the 30 or 50°S inferences by about 20%, almost at the same level for ail molécules. The mixing ratios at 30°N are 20% higher than the équatorial values except for the acetylene which is almost 10% higher.

When moving from the South to the North, I hâve found a definitive trend for increased gaseous content in the stratosphère. Additionally, the enhanced abundances hâve been observed in most cases at latitudes higher than 30°N during the duration of the northem latitudes.

In the latitudinal bin of 50°N, ail abundances except for CO2, are thus significantly enhanced with respect to 50°S and the equator. These results are consistent with our previous work with this model (Coustenis et al. 2007;2010b) and with Voyager 1 results (Coustenis et al. 1991;Coustenis & Bézard 1995). The 50°N values are higher than the 30°N by 20% and by 40% with respect to the equator.

These enhancements appear to be greatest for the shortest-lived Chemicals which are explained by a combination of chemistry and dynamics. The vertical distributions are steepest for the shortest-lifetime species in a purely Chemical model, while the presence of a circulation cell in the real atmosphère, with subsidence in the polar régions (at the north pôle during the period of the observations), causes the lower stratosphère to be greatly enriched in these species at the winter pôle as compared to the equator. The study of the temporal évolution of this enhancement in the higher northem régions within the Cassini mission duration, as well as changes observed with time at southem polar régions are described in more detail by Teanby et al. (2008;2010 and2012, in press) as well as in a work to which we contributed and will be published by Vinatier et al (2012, in préparation). 4.24). C^Hs seems to hâve increased abundances around NSE, but the data analysis results for this molécule at lower latitudes is rather uncertain.

We hâve detected a maximum around the time of the Northern Spring Equinox in mixing ratios with increases in abundances by about 30-40% for C3H4 and C2H6, 60-70% for C2H2 and C4H2 and about a factor of 2 for HCN and C2H4 (albeit with higher uncertainties) relative to the adjacent time periods in 2008 and 2010. Ethane shows variations with time only at this higher latitude.

The observed increase in abundance for some molécules is followed by a strong decrease which reduces considerably the observed enhancement by 2010. It is a quick decrease, within 1-2 terrestrial years and brings the abundances back to their levels prior to the ascent.

There is only one CIRS FP3 sélection at high resolution available in nadir viewing geometry after the NSE and one in lower resolution. The latter confirms this resuit, but further vérification is required. Additionally, both high and low-resolution nadir data taken in the 2010-2012 period seem to support the observed increase as well as the follow-up decrease (see open circles in Fig. 4.24).

We hâve also derived the 12C/13C ratio in CH4 and C2H2 to be 95 ±15. Similarly, the 14N/15N ratio in HCN has been inferred to be 50 ±10 for the latitudes analyzed in this section as an average. The l2C/13C ratio is consistent with results reported by [START_REF] Vinatier | The Titan l4N/15N and 12C/13C isotopic ratios in HCN from Cassini/CIRS[END_REF] and Nixon et al. (Nixon et al., 2008a;[START_REF] Raulin | Planetary science: Organic lakes on Titan[END_REF] and with the terrestrial inorganic standard value (88.9, Fegley 1995). The 14N/15N ratio is also consistent with the values given by [START_REF] Vinatier | The Titan l4N/15N and 12C/13C isotopic ratios in HCN from Cassini/CIRS[END_REF] and about 4.8 times lower than the terrestrial value (272, Fegley 1995).

Much stronger Chemical and température variations hâve been reported at higher northem or Southern latitudes (poleward of 70°N or 70°S) within the recent years (Vinatier et al., 2012b;Teanby et al., 2012, in press). 4.7.3 Comparison between nadir and limb data at the NSE I hâve tried to verify the NSE enhancement by performing limb data analysis in collaboration with Dr. S. Vinatier. However, there are no available high-resolutio limb spectra of this period at 50°N, only in July 2009 we hâve spectra taken at 15 cm'1 resolution during the T59 Cassini flyby. The comparison with this data delivered a discrepancy. No enhancement is observed in T59 limb spectra. We hâve checked if this discrepancy cornes ffom the température profile used in both retrievals (nadir and limb).

The limb température profile is indeed hotter by about 10 K at 1 mbar compared to the nadir one, while their différence is less significant in higher altitudes (Fig. 4.25). I hâve to note here, that the limb température profiles hâve been retrieved ffom limb spectra from S.

Vinatier, while R. Achterberg provided the nadir température profiles. Since limb spectra do not give information below 2 mbar, the initial guess used for retrieving both profiles does not affect the results. We hâve run some tests in order to understand the discrepancy in the thermal profiles.

The nadir température profile could not fit the limb FP4 data except for the altitude of 213 km or 0.3 mbar (see Fig. 4.26). did not fit well the wings of CH4 band at 1304 cm1. The température profile is hotter by 10 K than the nadir température profile and a colder profile is needed for the probed altitudes.

x 10'

Figure 4. 27 -Fit of the v4 methane band at 1304 cm'1 by using ARTT and the limb température retrieval. This température profile is too hot in the probed atmospheric level for fitting the nadir data, especially in the wings.

The contribution functions of the methane Q-branch and its wings when we used the limb température profile are plotted below (Fig. 4.28). The contribution functions confirm that the methane central branch as well as its wings peak at the pressure levels where nadir and limb thermal profiles hâve a 10 K différence and are therefore not compatible. The HCN and C2H6 limb vertical distributions fit the nadir data. The C2H2 limb vertical distribution is lower than the one we hâve calculated by using constant-to-height mixing ratio and it needs adjustment. The retrieved value for C2H2 is 6.25x10'6. The vertical distributions from S. Vinatier fit better the spectra compared with the GCM files that hâve a steeper slope.

In discussions that ensued, R. Achterberg pointed out that these observational sélections are not comparable because they are not really looking at the same latitude as the limb observations in the v4 methane band are also looking at about 10 degrees of arc or more in front of the tangent points. Indeed, the 2009 sélection is at a fairly high émission angle.

Assuming that the latitudes used for averaging are the surface latitudes, the latitudes in the upper stratosphère can be different by several degrees, and at 50°N, the latitudinal température gradients are fairly large. Therefore we renounced in trying to fit the limb and the nadir sélections taken at the same time but looking at different real estate and hâve focused on the nadir spectra.

Interprétation of the results

In this section, I hâve retrieved the abundances of trace gaseous compounds in Titan's stratosphère from 50°S to 50°N (Fig. 4.24). CO2 présents no latitudinal variations anywhere because of its long photochemical lifetime. For the other molécules, mid and Southern latitudes show no significant temporal variations during the Cassini mission.

However, I confirm that a rapid change in the atmosphère took place at 50°N during the Cassini mission (Teanby et al., 2010, Bampasidis et al., 2012a). A compositional enhancement is reported with an indication for a maximum at the time of the Titan northern spring equinox (NSE -August 2009), followed by a sharp decrease of the gaseous Chemical content within the next terrestrial year that remains to be confirmed with further data. These results are compatible with the findings of Teanby et al. (2010) in that I find HCN and C2H2 to display a rapid increase in northern latitudes up to mid-2009, while the abundances at équatorial and Southern latitudes remain stable.

Short-term variations observed during the Cassini mission can arise from changes

in the circulation around the equinox The collapse of the detached aérosol layer (West et al., 2011) suggests that the dynamics during this period go through a rapid transition wltich should also affect the gas distribution. The rapid decrease after NSE for wltich the most straightforward explanation is that the vortex has shrunk somewhat, would be consistent with the weakening thermal gradient and the profile becoming more constant with height, losing it warm mesosphere that we find here. Also compatible is the weakening of the winds as reported by Achterberg et al (2008;2011) and Teanby et al. (Teanby et al., 2009b). The finding also ties into the location of the maximum température gradient, which appears to be moving northward over the winter/spring season (Teanby et al., 2010, see Fig. 3, T panel). If 50°N is emerging from the vortex core, it would cause a large réduction in the abundances, hence explaining our observations. Thus, réduction at 50°N could be due to a weakening vortex with reduced latéral mixing across the vortex boundary (Teanby et al., 2010).

Another cause could be the spatial (due to Titan's inclination) variations in the energy input to Titan 's atmosphère as a driver for changes in the advection patterns, which in turn provide a stronger variability in the latitudinal abundances of photochemical species. Cassini entered into Saturn Orbit Insertion (SOI) at Ls=293°. The NSE is at Ls=0°a nd the average insolation at the top of the atmosphère has been increased (Fig. 4.33).

Daily average insolation (W/m2): x = 0

Water vapor in CIRS spectra

The first détection of water vapor in Titan's atmosphère was made in disk-average observations taken with the Short Wavelength Spectrometer of the Infrared Space Observatory (ISO/SWS) yielding a mixing ratio of 4xlO'10 at about 180 km (Coustenis et al., 1998). Due to their low signal-to-noise ratio, the first CIRS spectra could not provide the détection of water vapor and only an upper limit of 9xlO"10 was inferred [START_REF] De Kok | Characteristics of Titan's stratospheric aérosols and condensate clouds from Cassini CIRS far-infrared spectra[END_REF]. After the end of the Cassini prime mission, more spectra are available and the signatures of the weak molécules can be detected in large spectral averages where the signalto-noise ratio is high. In the paper of Cottini et al. (2012a), which I hâve co-authored, we confirm the détection of stratospheric water vapor on Titan using CIRS spectra. The analysis of these observations was performed by using two independent radiative transfer codes: the NEMESIS code used by the Oxford team and by GSFC colleagues -Non-linear Optimal Estimator for Multivariate Spectral Analysis (Irwin et al., 2008) and our newly updated ARTT -Atmospheric Radiative transfer for Titan (Coustenis et al., 2007;2010b;Bampasidis et al., 2012a). NEMESIS is a retrieval code that performs a combination of correlated-k forward model computation (Lacis & Oinas, 1991) and the retrieval scheme based on the method of optimal estimation (Rodgers, 2000). I was asked by Dr. V. Cottini (GSFC/NASA) to work on spécifie CIRS queries by using ARTT in order to validate the NEMESIS results. 5, 170.4, 177.56, 202.77, 208.50, 227.8, and 253.93 We First fit the continuum at the FP1 range thus providing an ab initio simulation of the aérosol contribution in there (as indicated and more properly done by de Kok et al. 2007a) which however serves our purpose sufficiently well since we focus on the lines. We then make different tests using various values for the water abundance, starting with of lxlO'10. By fitting the different lines, with particular accent to the line at 150 cm'1 which shows the strongest émission, and then getting weighted averages of the inferred best-fit values, we obtain a final best-estimated water vapor mixing ratio. (Down) Plot of the différence between model and observations. The vertical black lines depict the positions of water lines in the spectrum.

The same constant value fïts also the water lines for the 45°S-10°S sélection (Fig. 4.38), except perhaps the 250 cm'1 line. (Down) Plot of the différence between model and observations. The vertical black lines depict the positions of water lines in the spectrum. Except for 250 cm'1, ail the water lines are fitted.

We then use a slightly higher value, at 2xlO'10, and the line at 250 cm'1 is fitted but the fit is not well for the other lines (see Fig. 4.39). (lower panel) Plot of the différence between model and observations. The vertical black lines depict the positions of water lines in the spectrum. Only the line at 250 cm'1 is well fitted.

For the latitudinal bin of 80°S-45°S, the simulation resuit is the same. The water vapor constant-with-height abundance of lxlO"10 fits ail the lines except for the one at 250 cm"1 (Fig. (Down) Plot of the différence between model and observations. The vertical black lines depict the positions of water lines in the spectrum. Only the line at 250 cm'1 is well fitted.

We find that the 254 cm'1 water line is best fitted with 2±0.8xl0"10. For the other lines the H2O abundance is rather l±0.4xl0'10. The Full Width at Half Maximum (FWHM) of the contribution function for 254 cm'1 water line is slightly higher than for the others, so that seems to then bring us into compatibility with ISO, where a higher value is observed at higher altitudes for the 228 and 254 cm'1 lines (4±2)xlO'10. The ISO 4.0xl0'10 constant value profile had a contribution function peaking at 180 km and not near 100 km (Coustenis et al., 1998).

The error bar we find is ±2x10'10, and this includes ail the uncertainties due to calibration, methane mixing ratio, Satum background, etc. The water vapor value retrieved from Herschel Space Observatory data is in good agreement with the ISO results (Moreno et al., 2012, in press).

From the analysis of nadir spectra with the NEMESIS radiative transfer code, we hâve derived a mixing ratio of (1.4±0.5)xl010 at 93+37-io km (Table 1 of Cottini et al., 2012). In the same paper, using limb observations, we obtained mixing ratios of (1.3±0.4)xl0'10 at 115+47_22 km and (4.5 ± 1.5)xlO'10 at 230+45^0 km of altitude, confirming that the water abundance has a positive vertical gradient as predicted by photochemical models (Cottini et al. 2012).

The important resuit from these simulations I hâve performed on the CIRS available FP1 water lines is that the CIRS resuit is compatible with the value found by ISO and that the photochemical models for water vapor should be revised. The ratio of CO2/H2O on both Satum and Titan is unexpectedly large and Enceladus variable activity has been suggested to be its source (Moreno et al., 2012, in press).

4.9

Benzene vertical distribution profiles

In this section, I présent the results of analysis of the averaged Cassini/CIRS (Composite Infrared Spectrometer) spectra at both medium (2.5 cm'1) and high (0.5 cm'1) resolutions during the first six-and-a-half years of the mission focused on the benzene mixing ratio.

Benzene in Titan's environment

The presence of benzene (CôHô) in Titan's atmosphère is extremely interesting, as it is the first and the simplest polycyclic aromatic hydrocarbon (PAH) detected. In particular, following its first détection at 674 cm'1 based on ISO/SWS data (Coustenis et al., 2003), benzene was also detected in the thermosphère (950-1150 km) by modeling Cassini/INMS data (Waite et al., 2007) and firmly in the stratosphère (100-200 km) in ail latitudes [START_REF] Flasar | Titan's atmospheric températures, winds, and composition[END_REF]Coustenis et al., 2007;2010b). Moreover, benzene has been tentatively identified on Titan's surface by Huygens/ CGMS measurements (Niemann et al., 2005).

The unsaturated molécule of benzene has a six-carbon ring with three double bonds altemating with single ones. The length of its bonds are equal in ail carbon pairs with a value of 139 pm, while the électron density in the regular hexagon of benzene is identical. It is well represented as a résonance hybrid of two équivalent forms with six pi électrons, and thus its conjugated molécule is more stable than typical alkenes. Comparing to the whole organic inventory, aromatic substances perform different Chemical behavior [START_REF] Mcmurry | Organic Chemistry[END_REF].

Benzene is a key compound for the synthesis of polycyclic aromatic hydrocarbons (PAHs)

and may play an important rôle in the formation of the hydrocarbon haze of Titan. The planar molécule of benzene could help in understanding of the photochemical mechanisms occurred in Titan. When two benzene rings fuse together the molécule of naphthalene is created.

Naphthalene has not yet been identified in Titan.

The presence of the cyclic conjugated molécule of benzene in Titan's atmosphère demonstrates multiple pathways to the possible synthesis of biological building blocks.

Additionally, the existence of the liquids on the surface of the satellite (Stofan et al., 2007) in combination with the low température of 93 K (Fulchignoni et al., 2005) could host the proper environment for this biosynthesis. Recent laboratory experiments hâve showed that aromatic compounds could plausibly be produced on icy surfaces (Menor-Salvân et al., 2008).

Considering ail the above described characteristics, the prebiotic potential of Titan is significant and a large effort in astrobiological research has recently been focusing on Titan's environment (Raulin et al., 2008). The firm détection of hydrocarbons and especially acetylene and benzene has really enhanced the inventory of photochemical organic products.

Moreover, the abundance of benzene in the thermosphère shows a value four times higher than in the stratosphère, as it discover by the Cassini/INMS (Waite et al., 2007).

Combined with acetylene, benzene could produce polycyclic aromatic hydrocarbons (PAHs) or/and simpler compounds, the polyphenyls. The polycyclic aromatic hydrocarbons (PAHs) seem to be the most interesting product of the photochemical chain. However, since Titan lacks oxygen and sufficient température as Early Earth had, probably different pathways from Earth's hâve been followed and polyphenyls may possibly be created (Delitsky & McKay, 2010). The signature of benzene (CôHô) at 674 cm'1 is very weak in the CIRS spectra and only by large averages and/or at medium resolution reliable abundances can be retrieved. For this reason, I hâve gathered the same Cassini/CIRS averaged spectra from Tb to T44 flybys as presented in [START_REF] Coustenis | Titan trace gaseous composition from CIRS at the end of the Cassini-Huygens prime mission[END_REF] -Table 1). The new sélections hâve been improved by the fact that new calibration inferences were incorporated in the CIRS database. These data cover the beginning of the mission on July 2004, until the T44 flyby in late May 2008. In this work, as no longitude variation has been observed in Titan's atmosphère, we hâve scanned the globe in latitude bins of about 10 , starting from the equator and ending at the pôles. We hâve analyzed data of both medium and high spectral resolution (2.5 cm'1 and 0.5 cm'1 respectively).

We refer here only to nadir observations with émission angles restricted to less than 60°. The corresponding température profiles are retrieved by R. Achterberg following the method described in Achterberg et al (2008;2011). The results of this research were presented at the Faraday Discussions 147: Chemistry of the Planets, which was held in Saint Jacut de la Mer (France) from 14 to 16 June 2010.

CIRS detected benzene above 40°S and its profile is significantly affected by the downwelling part of the circulation at this région. Benzene has been detected in a large abundance in Titan's upper atmosphère [START_REF] Waite | Ion Neutral Mass Spectrometer Results from the First Flyby of Titan[END_REF] and its formation has been correlated with ion-neutral Chemical reactions taking place in the thermosphère [START_REF] Vuitton | Formation and distribution of benzene on Titan[END_REF]. The contribution of neutral chemistry in CôHô production is small and focused in the stratosphère. Thus, both production mechanisms, ion and neutral, must be included in the From the model calculations, the benzene mole mixing ratio profiles exhibit an almost constant-with-height value from 600 km up to 1000 km in the upper atmosphère. In the stratosphère, CIRS measurements (green boxes) are well satisfied by the model except in the higher Northern latitudes. However, the model overestimates the benzene abundance for latitudes greater than 40°(Fig. 4.43a-f). with green boxes (Bampasidis et al., 2010).

From the équatorial latitudes to the South pôle, the model of C^Hô mixing ratio reproduces very well the CIRS inferences (Fig. in Figs. 4.43a,c,d,e,f,the INMS results for the benzene retrieved abundance at 1100 km (Waite et al., 2007;[START_REF] Vuitton | Formation and distribution of benzene on Titan[END_REF]Cui et al, 2009) -in blue boxes -are also presented. The simulation also satisfies the INMS value for the latitudes of 23°N and 33°N (Figs. 4.43ef). For higher latitudes from 38°N to 53°N (Figs. 4.43a,c,d) In the paper of Jennings et al. (2012, in press), in which I participate, we associate the émission at 220 cm1 seen for the first time to appear in Southern CIRS/FP1 spectra, with stratosphcric condensâtes enhancement rising in the south. Since the South pôle of Titan moves deep into the winter season, the lack of solar EUV and UV provide downwelling accumulation of nitriles in the stratosphère and form condensâtes. However, the increase of the 220 cm'1 émission was much faster compared to the decrease in the north. Additionally, in the north, HC3N is présent in CIRS September 2011 spectra but not in June 2011 ones, suggesting that the mixing ratio of the molécule initiated its enhancement a year before the émission at 200 cm"1.

The decrease of cyanoacetylene in the northern spectra indicated that changes in the condensâtes and gases are observed indicative of seasonal changes (Vinatier et al., 2012a;2010b;Jennings et al, 2012, in press). The change in season is also evidenced in the south by the appearance of the 200 cm1 émission band (Fig. 4.47) described by an exponential law within a lifetime of 4.4 years (Jennings et ai, 2012, in press). Seasonal Phase (degrees) The explanation is that during the Southern summer, the sunlight influences significantly the trace gaseous composition in the stratosphère. Due to the low supply rate from the upper atmospheric layers, the condensâtes precipitate on the surface. However, as Titan 's

Southern hemisphere moves into the winter, haze particles start to form and therefore the relative émission increases in the spectra.

Jl North

Search for weaker molécules in CIRS spectra

Large CIRS nadir FP3 averages are necessitate to observing the signature of the weak trace gases and the isotopologues in Titan's atmosphère, since the use of such averages reduces the signal-to-noise ratio. Fortunately, CIRS has provided a huge number of datasets, which can be averaged in order to retrieve large averaged spectra. This work was presented at the EGU General Assembly 2012 in Vienna [START_REF] Bampasidis | Evolution of minor trace gases and isotopic ratios in Titan's stratosphère using CIRS/Cassini spectra[END_REF].

I hâve queried the 2009 CIRS annual FP3 spectra and assembled them in an average (Titan flybys T50 to T65) at 50°N, consisting of 1306 spectra at high resolution. Large averages for searching for weaker molécules are summarized in Table 4.16 below. signature of its isotopologues are also shown. For FI|JCN, only a 3-cr upper limit of 1.7x10"8 was derived. Top (a): Averaged spectra of T50 to T64 flybys at high resolution (0.5 cm'1) (in blue) at 50°N, best-fit model with main gases and isotopologues (in red) and best-fit model without H13CN isotopologue at 706 cm'1 (in green).

Bottom (b): Différence between model and data for both cases: with ail gases and isotopologues (in green) and without H1 'CN 706 cm'1 (in blue).

In Figure 4.48-a, I hâve plotted the model atmosphère with and without the H13CN contribution. The residual between data and model in each fit is also depicted in Figure 4.48b. For H13CN at 706.2 cm1, only a 3-a upper limit of 1.7xl0'8 has been derived, which is comparable with previous results in CIRS limb data [START_REF] Vinatier | The Titan l4N/15N and 12C/13C isotopic ratios in HCN from Cassini/CIRS[END_REF]. I hâve also

• 12 13
derived the "C/ C ratio of this average at 64.7±14 which is consistent with the relative ratio of 68 M6_i2 calculated in limb spectra at 83°N [START_REF] Vinatier | The Titan l4N/15N and 12C/13C isotopic ratios in HCN from Cassini/CIRS[END_REF], the estimation of 80.2±2 [START_REF] Nixon | Isotopic Ratios in Titan's Atmosphère from Cassini CIRS Limb Sounding: C02 at Low and Midlatitudes[END_REF] and the Huygens GCMS in situ measurements (82.3±1) (Niemann et al., 2005). This ratio is similar to the terrestrial one (at 89). The same procedure has been conducted for ,3C,80160 at 662.5 cm '1 (Fig. 4.49). We hâve derived only a 3-o upper limit of 8xlO'10 and the 160/180 ratio at 386±130 which is comparable with previous estimations [START_REF] Nixon | Isotopic Ratios in Titan's Atmosphère from Cassini CIRS Limb Sounding: C02 at Low and Midlatitudes[END_REF]. However, this resuit is significantly lower than the terrestrial ratio (circa 500) [START_REF] Bampasidis | Evolution of minor trace gases and isotopic ratios in Titan's stratosphère using CIRS/Cassini spectra[END_REF].

This test shows that our upgraded radiative code with the new molécules incorporated can help detect the signatures of new isotopologues or even entirely new species by fitting better the data where already identified molécules appear. We aim to apply this search for new molécules in the future with large averages seeking to identify C2H5 or other such predicted molécules, see Table 4.15.

Cassini will continue orbiting Satum up to the next summer solstice in 2017. Thus, a considérable amount of CIRS spectra will be available for studying Titan's stratosphère. The follow-up of this study is to investigate temporal variations on Titan due to seasonal and dynamical effects. The reported peak and sharp decrease in the abundances at the northem spring equinox at 50°N needs confirmation from northem CIRS spectra when are available.

Moreover, it will be very interesting to include in our photochemical analysis of benzene vertical distribution Cassini/INMS measurements in a global sense.

Seven years after Cassini's arrivai to the Satumian System, the combination of The cause of these variations should be also investigated. A further discussion on the influence of the solar cycle on Titan's stratospheric interannual variations in chemistry and température should be conducted. In fine, this allows us to set constraints on seasonal, photochemical and general circulation models for Titan and to better describe the physical processes in this complex environment.

Chapter 5

Numerical tools for analyzing the Titan lakes environment

In the previous Chapter, I hâve studied Titan's atmosphère focusing on the temporal and spatial variations of trace gases as derived from the Cassini Composite Infrared Spectrometer (CIRS) data during the Cassini-Huygens mission, from 2004 up to early 2012. Cassini's remote sensing measurements as well as Huygens' in situ observations hâve also revealed a variety of features on Titan's surface uniquely similar to the terrestrial ones. Comparative

Planetology studies can provide essential information about the physical processes that form or formed these surface features.

Although Titan's atmosphère is the basic subject of my research, this active geological environment and the processes that occur or occurred on Titan's surface as well as the exchanges with its atmosphère and interior are also among my scientific interests. This is due to the fact that, besides my Physics background, I also possess a degree in Geology. The subject of my thesis included the study of "Titan's environment" and this of course bears relevance to the interactions between the atmosphère and the surface or even the interior and my contribution in that respect to Titan science is what is described in this chapter. I focused mainly on the manifestations of the methane cycle in its connecting the atmosphère, the surface and the interior.

To that effect, the Cassini/RADAR discovery of numerous liquid organic depositslakes -on the surface of Titan bearing witness to the condensation of methane and its products and their deposits on the surface. Without doubt, lakes are the idéal medium for linking the surface processes, the atmosphère and the interior on Titan. By studying the lakes, we can infer valuable information about the complex processes occurring within Titan's environment and understand better the pattern of the methane cycle on a global scale.

My contribution and involvement in the surface studies, which I describe in this

Chapter, focus on lakes as a reference point and they concem either the study method tools or the observation data processing. The outcome of these studies provides further understanding of the liquid extends contribution to the environment, their connection with the atmosphère and their contribution to Titan's methane cycle.

The lakes may also offer the proper habitability conditions for hosting biological building blocks on Titan, a possibility that motivated me to work on Titan's astrobiology, as it is described in Chapter 6. Moreover, these surface studies, mainly performed in collaboration with A. Coustenis, A. Solomonidou, E. Bratsolis and M. Hirtzig (but also with other international colleagues such as Jean-Pierre Lebreton and Hunter Waite), inspired me to propose the MEMS experiment for Titan, which is described in detail in Chapter 7.

Specifically, Titan's géologie environment is defined in Section 5.1 and the exchange processes as well as the relationship between the major géologie forces are described in Sections 5.2 and 5.3 respectively. In the Sections 5.4 and 5.5 that follow, I contribute to the study of the morphotectonic expressions of the Kraken Mare lake surroundings using géomorphologie numerical analysis tools. The Cassini RADAR image processing technique using a probabilistic filtering in which I got involved is discussed in Sections 5.7 and 5.9 respectively.

Titan's geological environment

Techniques and experiments adapted from Earth Science are also applied to distant rocky planets or satellites and provide essential information about the surface processes in other planetary environments. The Apollo Passive Seismic Experiment (PSE), a network of seismic instruments consisted of four seismometers deployed by the astronauts on the lunar surface between 1969 and 1972 provided the internai structure of the Moon [START_REF] Dainty | Seismic Scattering and Shallow Structure of the Moon in Oceanus Procellarum[END_REF][START_REF] Nakamura | Apollo lunar seismic experiment -final summary[END_REF][START_REF] Bulow | New events discovered in the Apollo lunar seismic data[END_REF]. The recently selected Discovery mission to Mars, InSight (for Interior exploration using Seismic Investigations, Geodesy and Heat Transport), planned for launch in 2016, will perform geophysical experiments on the Red planet to détermine its interior structure and reveal its thermal history.

This also can apply to Titan. After the Voyagers encounters of Titan, the existence a global deep was suggested (Flasar, 1983;Lunine et al., 1983). Both methane and ethane can exist in ail three States (vapor, liquid and ice) in Titan's surface température and pressure conditions (94 K and 1.5 bar, respectively c.f. Fulchignioni et al, 2005) and they can cause précipitations perhaps responsible for some at least of the surface lakes. But, according to [START_REF] Lorenz | The life, death and afterlife of a raindrop on Titan[END_REF], methane drops would evaporate before reaching the surface and the existence of pure methane surface deposit is then doubtful. According to Flasar (1983), the partial pressure of methane would be lowered in a mixture of methane and its photochemical byproducts (ethane, propane). The Lunine et al. (1983) model suggested the presence of ethane surface océan of 75% ethane, 20% CH4 and 5% nitrogen about 1 km deep. [START_REF] Stevenson | Interior of Titan. Paper presented at the In ESA[END_REF] suggested that the hydrocarbon océan may be located under the surface of Titan, by assuming a porous enough surface bedrock.

The global océan hypothesis was rejected by post-Voyager ground-based RADAR echoes, which indicated that the surface of Titan is not homogeneous nor compatible with global liquid coverage. The high reflectivity echoes recorded by Very Large Array (New Mexico) radio-telescope suggested significant surface variability consisting of 'dirty' ice (Muhleman et al., 1990). Steward Observatory and the Multiple Mirror Telescope data also showed a variation in Titan's albedo [START_REF] Lemmon | Titan's Rotation: Surface Feature Observed[END_REF], while Arecibo echoes in 2001 and 2002 marked areas of specular reflections, similar to small lake-like features [START_REF] Campbell | Radar Evidence for Liquid Surfaces on Titan[END_REF].

Moreover, Titan's atmosphère is transparent through methane Windows in the near infrared (see § 1.1.4) and the Infrared Telescope Facility on Mauna Kea Earth-based observations hâve confirmed RADAR measurements (e.g. Griffith et al., 1991, Griffith 1993).

This data hâve triggered updates to previous models about the presence of surface liquids and according to [START_REF] Dermott | Tidal effects of disconnected hydrocarbon seas on Titan[END_REF], liquid hydrocarbons could be restricted in small surface basins.

Cassini-Huygens surface observations

In January 2005, the Huygens probe performed the only in situ measurements to date of Titan's surface during its descent and landing [START_REF] Lebreton | An overview of the descent and landing of the Huygens probe on Titan[END_REF]. It observed the atmosphère and the surface during the descent from 170 km downwards with the Descent Imager/Spectral Radiometer (DISR) [START_REF] Tomasko | The Descent Imager/Spectral Radiometer (DISR) Experiment on the Huygens Entry Probe of Titan[END_REF] and measured its composition above and on the landing site with the Gas Chromatograph and Mass Spectrometer (GC-MS) [START_REF] Niemann | The Gas Chromatograph Mass Spectrometer for the Huygens Probe[END_REF]. The Huygens Atmosphère Structure Instrument (HASI) [START_REF] Fulchignoni | The Characterisation of Titan's Atmospheric Physical Properties by the Huygens Atmospheric Structure Instrument (Hasi)[END_REF] probed the atmospheric properties (T,p and density) during the descent from 1600 km downwards and determined the surface température and pressure, while the the Surface Science Package (SSP) [START_REF] Zamecki | Huygens' Surface Science Package[END_REF] investigated the local surface material.

Huygens measurements revealed that Titan's surface is mostly dry currently near the equator, but also compatible with the presence of wet ground close undemeath the surface as indicated by the increase in methane gas following évaporation from the location where the warm Huygens probe landed. The surface is also suggested to be covered by water and hydrocarbon ice, silicate rocks and organics/tholins frozen material [START_REF] Tomasko | Rain, winds and haze during the Huygens probe's descent to Titan's surface[END_REF][START_REF] Zamecki | A soft solid surface on Titan as revealed by the Huygens Surface Science Package[END_REF]Lebreton et al., 2009).DISR images showed a well-developed surface network of channels (Fig. 5.1) close to the probe's landing site [START_REF] Tomasko | Rain, winds and haze during the Huygens probe's descent to Titan's surface[END_REF]. Except for the Huygens observations of river-like features, Cassini/SAR has recorded a variety of fluvial channels networks on Titan's surface (Lorenz et al., 2008a) but with much lower spatial resolution. It has been suggested that the dry channels at the low latitudes are caused by earlier précipitation during a different season [START_REF] Turtle | Rapid and Extensive Surface Changes Near Titan's Equator: Evidence of April Showers[END_REF].

Figure 5. 1 -Titan's surface diversity, recorded from Huygens/DISR [START_REF] Tomasko | Rain, winds and haze during the Huygens probe's descent to Titan's surface[END_REF]. The panoramic view from 8 km shows that the surface is solid and inhomogeneous at the Huygens landing area (the circle in the middle of the image).

Since its First flybys of Titan from 2004, the Cassini orbiter maps remotely Titan with the Imaging Science Subsystem (ISS) [START_REF] Porco | Cassini Imaging Science: Instrument Characteristics And Anticipated Scientific Investigations At Satura[END_REF], the RADAR (Elachi et al., 2004) and the Visual and Infrared Mapping Spectrometer (VIMS) [START_REF] Brown | The Cassini Visual And Infrared Mapping Spectrometer (Vims) Investigation[END_REF]. The Composite Infrared Spectrometer (CIRS) (Flasar et al., 2004) is also used to map the surface températures. an inhomogeneous and complex topography (Elachi et al., 2005;Porco et al., 2005;Elachi et al., 2006). Titan's topography is relatively low and its lithosphère is too soft and flexible and/or heavily eroded for high mountains to exist [START_REF] Lorenz | Hypsometry of Titan[END_REF], as it was expected [START_REF] Perron | Dynamics of an ice continent on Titan[END_REF]). However, the joint instrumental coverage of Titan's surface delivers a long list of features, which resemble the terrestrial ones in smaller scale height.

Titan has mountains (e.g. Radebaugh et al., 2007, Lopes et al., 2010), and chains of mountains known as ridges (e.g. Soderblom et al., 2007, Mitri et al., 2010), faults (e.g. [START_REF] Radebaugh | Régional geomorphology and history of Titan's Xanadu province[END_REF], canyons (Lorenz et al., 2008a), volcanic-like features [START_REF] Bames | Cassini observations of flow-like features in western Tui Regio, Titan[END_REF]Elachi et al., 2006;[START_REF] Stofan | Mapping of Titan: Results from the First Titan radar passes[END_REF][START_REF] Lopes | Much like Earth: distribution of géologie processes on Titan from Cassini RADAR data[END_REF], network river drainages [START_REF] Tomasko | Rain, winds and haze during the Huygens probe's descent to Titan's surface[END_REF]Elachi et al., 2006;[START_REF] Stofan | Mapping of Titan: Results from the First Titan radar passes[END_REF][START_REF] Bames | Near-infrared spectral mapping of Titan's mountains and channels[END_REF]Lorenz et al., 2008a), lakes (Mitri et al., 2007;Stofan et al., 2007), impact craters [START_REF] Wood | Impact craters on Titan[END_REF], numerous longitudinal dunes (Elachi et al., 2006;[START_REF] Lorenz | The Sand Seas of Titan: Cassini RADAR Observations of Longitudinal Dunes[END_REF]Radebaugh et al., 2008) and tropical oases [START_REF] Griffith | Possible tropical lakes on Titan from observations of dark terrain[END_REF].

Cassini/SAR has also made the first détection of large liquid deposits on Titan's surface ffom data obtained during the Cassini Tl6 (July 2006) flyby at the northem polar latitudes. The RADAR instmment observed more than 75 topographie dépréssions which hâve very low RADAR reflectivity, circular or irregular shapes and filled with low dielectric constant material (Stofan et al., 2007).

Titan has indeed a very complex surface formed by multivariable geological processes. The very small number of impact craters, counted on the surface so far (Wood et al., 2010), indicates that it is an active planetary body where strong exogenic forces occur and contribute to the shape of its landform. The main source of these exogenic forces is Titan's thick nitrogen dominated organic atmosphère, which provides the érosion agents. In the relation between the atmosphère and the surface and how it affects the methane cycle, I essentially focus on the création, évolution and contributions of the methane lakes and also concem myself with the dunes, sometimes called "sand seas".

Surface material properties and Titan's substrate

To study the action of endogenetic and exogenetic forces on Titan's surface, we hâve to focus on the material of the ground and the substrate. Conceming the latter, a number of questions arise. We would like to find answers about its structure, its composition, whether it is globally the same and if it présents similarities with the terrestrial surface.

The Huygens in situ surface measurements were constrained in the région of the landing site (10.3°S, 192.3°W c.f. Lebreton et al., 2005). Huygens/SSP in situ data shows a damp porous ground material consisting of loosely packed grains, resembling mud to medium sand or clay consisting of water ice, clathrate and organics [START_REF] Zamecki | A soft solid surface on Titan as revealed by the Huygens Surface Science Package[END_REF]Lorenz et al., 2006a). The environmental conditions existing on the satellite indicate that Titan's crust mainly consists of mixtures of water ice, organic ices, tholins and nitriles (e.g. Soderblom et al., 2007) in température and pressure of about 90-94 K and 1,5 bar respectively (Fulchignoni et al., 2005;[START_REF] Jennings | Titan's Surface Brightness Températures[END_REF]Cottini et al., 2012b).

The discovery of lakes on Titan's polar régions by both RADAR and ISS provides information about the surface material and the interior and the surface-atmosphere-interior exchange processes. Morphotectonic analyses on Earth hâve shown a relation between the formation of lakes and the géologie interior forces in several cases. The question is if such a mechanism is responsible for the formation of surface topographie dépréssions on Titan, where the liquid deposited through the atmosphère is sequestrated creating lakes.

To address this issue, we hâve to compare the mechanic response of the terrestrial silicate substrate to tectonic stresses with Titan's icy bedrock response. Analogous tectonic behavior provides similar results on the surface. If the two planetary bedrocks are submitted to comparable mechanical processes, similar interprétation techniques can be applied. The study and the corrélation of landform, tectonism and erosional processes is in the scientific field of Morphotectonics (e.g. Scheidegger, 2004).

In the paper of Solomonidou et al. (2012), we compare silicate and icy tectonism.

According to Collins et al. (2009) the icy crusts of the outer planets moons exhibit similar brittle response to terrestrial silicate crust when stresses are applied. However, while water ice and silicate rocks hâve similar ffictional strength on small planetary bodies with diameter less than 100 km (Beeman et al., 1988), in larger bodies when plastic yielding becomes important, the silicate rock is about 10 times stronger than the ice crust (Melosh & Nimmo, 2011). This means that silicate rocks can be deformed for a longer period of time and under higher pressure of about one order of magnitude, without reaching the fracture point.

Therefore, we conclude that the tectonic behavior between silicate crust and the planetary ice at 100 K is comparable. This fact encouraged me to help adapt tools from the terrestrial geomorphology for studying Titan 's surface, an effort that I describe in the next section.

Exchanges between atmosphère and surface in Titan's environment

The environmental exchanges between the atmosphère and the surface form spécifie types of surface features, which are either lakes (or seas, or rivers, or oases) that are mainly located at the polar régions in the current season, or dunes, which are distributed at the équatorial latitudes. As we hâve discussed in the previous Chapter, (Fig. 4.1) methane possesses a key rôle by driving an active hydrologie cycle (Atreya et al., 2006). Thus, the link between the atmosphère and the surface is the existence of lakes and dunes in addition to the erosional processes that reshape the already existed structures and resurface the entire terrain.

Lakes are mainly located in the polar régions of the moon, where the environmental conditions maintain the liquid State. On the contrary, the equator is dry enough to support the dune fields. However, recent radiative transfer analyses of VIMS data by [START_REF] Griffith | Possible tropical lakes on Titan from observations of dark terrain[END_REF] reported the existence of methane tropical liquids with a depth of 1 m within the latitudes from 20°N to 20°S. They suggested that these deposits are supplied by subsurface sources such as subterranean rivers. Therefore, the study of the surface liquids provides insights on the methane cycle.

5.2.1

Vast dune fields on Titan's surface

Cassini/RADAR images hâve shown many dune formations at équatorial latitudes mainly from 0°N to 30°N and in some isolated régions up to 55°N (Radebaugh et al., 2008). Due to their extent the équatorial dunes hâve been interpreted as sand seas [START_REF] Lorenz | The Sand Seas of Titan: Cassini RADAR Observations of Longitudinal Dunes[END_REF] and they are mainly concentrated in zonal east-west direction. The presence of these vast dune fields indicates the wind blow direction which is towards east, but they wrap around the topographically high features they meet like mountains and craters (Radebaugh et al., 2008).

Dunes are 1-2 km wide and 1-4 km apart, having up to 150 km height and more than 100 km in length (Elachi et al., 2006). Except from the RADAR swaths, the presence of the dunes has been confirmed by Cassini/ISS images (Porco et al., 2005) and the Cassini/VIMS spectra [START_REF] Soderblom | Topography and geomorphology of the Huygens landing site on Titan[END_REF][START_REF] Bames | Spectroscopy, morphometry, and photoclinometry of Titan's dunefields from Cassini/VIMS[END_REF]. These linear dunes differ in composition from their surrounding material (Radebaugh et al., 2008) and are clearly the resuit of the action of aeolian exogenic forces.

Liquid séquestration

The diversity of Titan's surface does not exclude the existence of surface liquid hydrocarbons. Theoretical studies and observations indicate that the presence of stable liquid phase on Titan's surface is plausible. The surface température ranges at about 90-94 K while the pressure at about 1,5 bar (Fulchignoni et al., 2005;[START_REF] Jennings | Titan's Surface Brightness Températures[END_REF]Cottini et al., 2012b) and compared to the triple points of main atmospheric constituents organic liquids can exist on the surface. Indeed, the triple points of nitrogen, methane, ethane and propane are 63.148, 90.68, 90.348 and 85.47 respectively (Jacobsen et al., 1986;[START_REF] Younglove | Thermophysical Properties of Fluids. IL Methane, Ethane, Propane, Isobutane, and Normal Butane[END_REF].

The channel networks at the équatorial latitudes that Huygens/DISR observed resemble the common pattern of terrestrial river channels. Their presence indicates that liquid provided by heavy précipitation can flow on Titan's surface although, no rain was observed during the descent of Huygens [START_REF] Tomasko | Rain, winds and haze during the Huygens probe's descent to Titan's surface[END_REF].

Several channel networks deposit in large lakes. Hence, the channel liquid and the lake liquid should hâve many similarities. The Chemical composition of Titan's lakes is considered to be a mixture of ethane, methane, propane, higher order hydrocarbons and nitriles [START_REF] Cordier | An Estimate of the Chemical Composition ofTitan's Lakes[END_REF]2012). Cassini/VIMS measurements of the Ontario Lacus during the December 2007 T38 flyby indicate that it is mainly consisted of liquid ethane (Brown et al., 2008), but without giving a mixing ratio. According to Hayes et al. (2008) the imaginary dielectric properties of the liquid hydrocarbons of Ontario Lacus observed by Cassini are consistent with the experimental values of liquid natural gas, but not similar to the tholin ones. By analyzing the same VIMS spectra [START_REF] Moriconi | Characterization of Titan's Ontario Lacus région from Cassini/VIMS observations[END_REF] suggest that the région that surrounds the lake feature is dominated by propane, butane and acetylene. The existence of propane consistent with [START_REF] Cordier | An Estimate of the Chemical Composition ofTitan's Lakes[END_REF]2012) The study of Titan's liquids, as well as their relation to their environment, provides evidence about their formation, their évolution and their influence to the local area. However, Titan's exogenic and endogenic forces act upon the liquid deposits and shape their morphometry. By understanding the impact of Titan's géologie forces in the liquids of Titan,

we will be able to détermine links between these forces and the methane cycle. In the following sections, I will describe the techniques adapted from terrestrial Geomorphology and image processing that we hâve applied to Cassini data towards this direction.

Lakes on Titan

Titan's lakes with a mean diameter greater than 200 km are listed in the Table 5.1 below.

Most of the other lakes are quite smaller with diameters less than 100 km. The northem great lakes of Titan, Kraken Mare, Ligeia Mare (the main target of the discovery candidate TiME mission) and Punga Mare, are large enough to be characterized as seas (according the IAU Committee on Planetary Nomenclature) and are illustrated in the Figure 5.3. Ontario Lacus is the largest lake located at the Southern hemisphere.

Titan's lakes hâve been divided into three classes according their RADAR retums: (a)

the dark -liquid filled lakes, (b) the granular -partially filled lakes and (c) the bright -empty lakes (Stofan et al., 2007). The dark lakes completely absorb the incident RADAR beam.

SAR microwave radiation pénétrâtes the liquid layer and interacts with the lakebed in the partially filled lakes. The granular lakes can be considered as the transition feature between dark lakes and empty basins. Empty lakes look brighter in SAR images compared to their exteriors and hâve a depth of 200-300 m (Hayes et al., 2008).

The area I hâve selected to study is the northem shoreline of Kraken Mare, where the island Mayda Insula is located. I hâve decided to work on this image because of the presence of this island. It présents an interesting landform inside the lake to study any morphometric corrélations with the topography of the nearby shorelines.

Kraken Mare is the largest liquid surface basin of Titan. It covers an area of about 400,000 km2 (Turtle et al., 2009), bigger than the Caspian Sea on Earth. Kraken Mare has been selected as a target of a future lake lander from the Titan-Satum System Mission (TSSM) c.f. TSSM Final Report JPL D-48148 NASA Task Order NMO710851. Due to gravity and morphology, the surface liquid moves downwards to lower places, forming streams or rivers. On its way, the stream collects more liquid from the atmospheric deposit (rain, condensation, etc) and from other streams joining in. Usually, the river discharges its load into basins, which are lakes or seas. The main attribute of this work is to quantify the hierarchy of stream segments in the observed drainage network according to the ordering classification System proposed by Horton & Strahler (Horton, 1945;[START_REF] Strahler | Hypsometric (area-altitude) analysis of erosional topology[END_REF].

This classification provides information about the degree of the évolution of the drainage network, describes the géologie history of the basin and eventually tells us the story of the lakes.

In the drainage System, the channel segment, which begins from the head of the stream, is assigned the value 1 and called first-order stream. When two first-order streams corne together, they form a second-order stream, two second order streams form a third order stream, and so on. Having measured the length of the stream segments and its basins, it is possible to estimate the liquid budgets of this drainage and the hydrocarbon lake supply.

Several stream ordering Systems hâve been proposed so far in the literature [START_REF] Shreve | Statistical law of stream numbers[END_REF][START_REF] Woldenberg | Horton's Laws Justifïed in Terms of Allometric Growth and Steady State in Open Systems[END_REF][START_REF] Scheidegger | Horton's law of stream order number and a temperature-analog in river nets[END_REF][START_REF] Doomkamp | Numerical analysis in geomorphology: an introduction[END_REF]), but Strahler and Shreve methods are used widely today. In this latter method the order of channels increases only when two streams of the same order join together, while in the [START_REF] Shreve | Statistical law of stream numbers[END_REF] model the order number increases for ail links of the network.

The statistical analysis of the streams number and lengths reveals some interesting aspects of the évolution of this spécifie area and studies the balance between the exogenic erosional forces and the endogenic tectonic movements. In addition, the measurement of the drainage density provides a useful numerical measure of the landscape dissection and runoff potential. Here, I présent these considérations in association with the morphotectonics influence on Titan's surface. about 90 km by 150 km across and it is centered at about 79°N and 310°W.

At the West-Southwest side next to the island, there is a surface of low radar reflectivity. It is not as dark as the lake at the other side of the island and no significant topography is observed on it. A similar plateau is also observed at the South coast of the island, but not at the East. This area has an increased radar brightness compared to the lake.

The same features hâve been previously observed at the edges of the lakes in Northern Titan and they might be due to a reflection from the lake bottom where it is sufficiently shallow so that the bottom écho is not completely attenuated (Stofan et al., 2007). These features can be indicators of ground élévation.

Moreover, the East and West coasts seem scarped, while canyons in an average length of 20 km exist. Mountain ridges are extended to the coastline as well, while their slopes are large and they may hâve a tectonic origin, as I hâve discussed above. Furthermore, remarkable topography with ridges has been developed in the Western part. The Northern région seems to be smoother than the other régions.

The question is whether the local topographie features are somehow related and whether there is any evidence of tectonic activity in this région. By studying the development of the drainage System, I can get constraints on the influence of the antécédent geological events and subséquent géologie changes.

For data acquisition, I cooperated with Dr. Randolph Kirk from the Astrogeology Science Center of the US Geological Survey already since 2007.

Methodology and application

The stream order classification is a statistical relationship of the stream segments in drainage basin [START_REF] Scheidegger | Horton's law of stream order number and a temperature-analog in river nets[END_REF]. In this work we use the Horton and Strahler's ordering System [START_REF] Horton | Erosional development of streams and their drainage basins; Hydrophysical approach to quantitative morphology[END_REF][START_REF] Strahler | Hypsometric (area-altitude) analysis of erosional topology[END_REF] which is a classical approach. Horton's laws are certain topological invariants that are needed to be satisfied, whereas these laws connect the numbers and the lengths of streams of different orders in a drainage basin.

1. Law of Stream Numbers: The numbers of streams of different orders in a given drainage basin tend closely to approximate an inverse géométrie sériés in which the first term is unity and the ratio is the bifurcation ratio.

2. Law of Stream Lengths: The average lengths of streams of each of the different orders in a drainage basin tend closely to approximate a direct géométrie sériés in which the first term is the average length of streams of the first order.

After précipitation, surface liquids form at high latitudes shallow rills. The lowest order streams, the rills, are the smallest outlying tributaries on the edges of the network, according to the Horton and Strahler's ordering scheme. At each unique point, where two tributary streams join, a new stream originates. Whenever two tributaries of the same order meet, the outgoing stream has one order number higher than that of its ancestors. If two tributaries of different orders meet, the outgoing stream has the same order number with the higher ordered tributary. Eventually, ail streams in the network combine to form the highest order (main)

stream. The number of streams of order u is Nu, while <LU> is the average length of streams of order u. Horton's Laws [4] State that the bifurcation ratio Rb and the length ratio Rl, given by: (Eq. 5.1) ^K+l (Eq. 5.2) <Lu> are constant, or independent of u.

The Hortonian System ordering has a recursive nature since each arc's order dépends on the order of its inflow tributaries and has been implemented in various Geographical Information System (GIS) tools [START_REF] Gleyzer | A Fast Recursive GIS Algorithm for Computing Strahler Stream Order in Braided and Nonbraided Networks[END_REF]. The algorithm behind these stream ordering GIS procedures is not revealed. This is not a problem in terrestrial geomorphology, since the plethora of observations in a région optimizes the GIS application. However, in the case of Cassini/SAR images of Titan, the resolution does not allow recursive procedures and I made the decision to do it manually. Therefore, I hâve decided to use Computer-aided design software (CAD). The results of this procedure are illustrated in Figure 5. The unnumbered lines in Figure 5.6 are the lst order streams whereas the 2nd order streams are marked by the label "2" and 3rd order streams by the label "3". The basins' borders of the 3rd order streams are designed with thick dark lines and mentioned by roman numbers.

Since there is no topographie map with contour lines available, but only a SAR image, I hâve designed the streamlines following the fluctuation of the color's intensity. Hence, it is possible to find more streams, especially of the lst order group in an image taken in higher resolution. The outlined drainage and its development can be used as a tool to study the évolution of the local surface.

Scale 20km

Table 5.2 below contains the results of the analysis of the 3rd stream order basins, according the Horton-Strahler's classification System. The Bifurcation Ratio Rb is also calculated according the Equation 1. Usually, the bifurcation ratio Rb has a range of 3 to 5 for welldrained networks, while values greater than 5 refer to deep and narrow basins. From the morphologie point of view, I hâve divided the island's 3rd order drainage basins in two different régions. The criterion of this grouping is the degree of the 3rd order channels' development. The first group contains the 3rd order basins labeled as IV, V, III and the second one contains the basins labeled as II and I (Fig. 5.6). The longer the main branch of a stream network is, the more evolved its respective basin will be. Better évolution of a drainage basin means that the whole network has been operated for a long time without being influenced from other géologie or atmospheric events. In the région of interest, the third order basins I and II hâve their main stream well developed along to the basin's axis, compared to the other three basins.

The total length of the streams in the 3rd order basins is listed in Table 5.3. Lu is the length value for each branch of the drainage System. The lengths of the 3rd order branches of basin I, III, IV and V are shorter by a factor of about 0.8, 5.8, 4.4 and 1.4 respectively, compared to the longest 3rd order stream of basin IL where only 2nd order streams hâve been developed. I divide these streams in two groups, the Northern and the Southern, according the orientation of their flows. The last column of the Table 5.4 contains the total number of the streams of the entire island. The last column of Table 5.4 contains the percentage of these 2nd order streams compared to the total stream number of the island. Consequently, according the drainage development, I divide Mayda Insula into two morphotectonic régions:

1. The eastem part (basins I and II), where the drainage networks are well developed, 2. The western part (basins II, IV and V), where the drainage networks are in initial stages of their évolution.

The average lengths of the Table 5.3 length values are calculated in the Table 5.5 below, in order to find the length ratio Rl according the Equation 2. <LU> is the average length of streams of order u in meters. The stream length ratios Rl, the area A of each 3rd order basin and the Drainage Density D are listed in Table 5.6. The drainage density is calculated by the équation:

D = 2*
(Eq. 5.3) From Table 5.6, the length ratio Rl between the 3rd and 2nd order of streams of basins I to V ranges 0.2 to 2.1. When the ratio is greater than 2, then the 2nd order streams joint the 3rd ones in acute angles, like most of the tributaries in an intense topography.

Interprétation of the results

The study of drainage basins morphometric parameters reveals much information regarding the évolution and the dynanties of the "hydrological" processes on Titan surface, despite the fact that the liquid material is not water but hydrocarbons. The corrélation between the geomorphic features that are expressed by these parameters and the local topography can provide a qualitative description of the région 's development Stream order classification is a simple qualitative basis for understanding the development degree of a spécifie area through time. Generally, a basin is better drained when its stream order is higher. The 3rd order basins ofMayda Insula cover more than half of the total area of the island. On the other half of the island, low order basins exist, indicating a lower degree of development. Additionally, the 3rd order basins I, III, IV and V are located at the island's edges, wltile basin II, the best geomorphologically developed one, is situated in the centre of island's East part.

The 2"d order streams of the well-drained basin II are quite long and deployed at the V4 of the surface of the island. The longest 2ud order stream ofthis basin seems to be shaped like a meander, another clue of the well-developed stage of basin II. The 3rd order streams of basins III and IV are not so much developed as their 2"d order stream ancestors.

The lst order streams fow parallel or are conséquent to the original slope of the surface. These rills are steep-sided, hâve short lengths and are V-shaped. The number of nd these streams ( 52), as Table 5.2 shows, is greater than the total stream number of the 2 (16) and 3rd ( 5) order at the 3rd order basins I, III, IV, V. Because of the almost vertical entrance of most of the tributaries into the main stream (Figs. 5.6-7), the drainage pattern can be described as rectangular.

The drainage density D, listed in Table 5.6, extends to zéro, although there is a considérable degree of basin development. This can be explained by the existence of intermittent and ephemeral streams in the région as well as by the low image resolution.

Although the designed drainage in Figures 5.6-7 seems to be dry, it can still operate (or operated in the past) and may evolve on the island's surface.

The heterogeneity of Mayda Insula 's surface may hâve been caused by two factors:

the atmospheric or the tectonic activity, operating separately or together. Tectonic activity indications on Titan's surface hâve been reported by Soderblom et al. (2007). In the case of tectonic quiescence, either vertical or horizontal, the landscape will be erodedfollowing the Davies érosion cycle [START_REF] Davies | The geographical cycle[END_REF].

The bifurcation ratio is a numerical expression of the development degree of a basin. The bifurcation ratio for basin II is Rb=2.17, which means that the basin is well drained. The bifurcation ratio of the other 3rd order basins is greater than that of basin II, due to their lower drainage development. The ratio of basin V is greater than 5, which on the Earth is an indicator of a newly or an elevated surface.

The Horton law of stream lengths shows a géométrie relationship between the average length of streams of a spécifie order and the corresponding order. The stream length ratio Rl is the measure of this relationship. Wlten the ratio Rl is greater than 2, it means that the 2nd order tributaries enter into the 3rd streams at acute angles, as most of the streams on steeper slopes. The steeper the slope, the more acute the angle of the stream entrance is [START_REF] Horton | Erosional development of streams and their drainage basins; Hydrophysical approach to quantitative morphology[END_REF], which is the case of Rl(2,1) (Table 5.6). Our morphometric analysis shows the existence of steep slopes on the Western part of Mayda Insula. The stream length ratio Rl(3,2) (see Table 5.6) ranges between 0.2 and 2.1 that means smoother slopes than in the previous case. The magnitude of slope is one of the characteristics that détermine the relief. Since the island basins hâve different slopes, they hâve experienced different élévations. Summarizing the results of the drainage network analysis and taking into account the stream length ratio Rl between 2nd order and lst order streams, we can assume that the 2nd order basins I, III, VI and V are elevated after the drainage development of basin II.

Statistical analysis of Mayda Insula surface édifices

To hâve a deep insight of the topography of Mayda Island, I studied the close shoreline areas as well. Three régions with significant topography hâve been distinguished in the Cassini's RADAR image, marked with A, B, C in Fig. 5.4.

Statistical analysis of the angular distribution of the mountain ridges

In the following Tables 5.7-9, the angles of the mountain ridges in degrees are listed, for the three régions respectively, showing the orientation of their axis, according to the N-S direction. The angle distributions of the mountain ridges' orientations are presented in Figs. 5.9-11. The normal curve is also included. The angular distribution of the ridges (Fig. 5.9) shows two preferred directions: (a) 20°to 40°e astwards and (b) 20°westwards.

West Région (B):

The angular distribution of the mountain ridges in Figure 5.10 demonstrates an orientation at 30°eastwards.

Southwest Région (C):

The angular distribution of the mountain ridges in Figure 5.11 shows an orientation of 20°e astwards.

The prominent mountain ridges orientations of the area of interest are two. The first one is located in ail the three régions (A, B and C in Figure 5.4) at 20°to 30°eastwards. The other one has been observed at the island at 20°westwards. The next interesting resuit of this statistical analysis is that at région A an opposite orientation exists. There are almost 30% of the ridges that are N-W oriented. Due to the low radar resolution, only a sample of the mountain ridges is examined. A high-resolution image could provide the accurate orientation of the ridges.

Corrélation between ridge distribution and drained basins of Mayda Insula

The drainage basins are defined by the landscape and each basin has the same orientation as its main stream. Intense topography exists in the western part of Mayda Insula with a radial development of the 3rd order basins (Fig 5 .7). Particularly:

• 3rd order basins:

Basin I: The central stream (3rd order) of the basin is northwest oriented.

Basin II: As I hâve shown previously in the numerical analysis, basin II is the bestevolved 3rd order basin in length and occupied area. The central stream of the basin as well as the longer 2nd order channels are meander, are northeast oriented like the relevant mountain ridges. The drainage follows the régional topography. By taking into account that the northeast is dominant and that the basin II is well drained, I conclude that this direction is the older orientation of the landscape in the Insula and the near-by shoreline.

Basin III: The central 3rd order stream of basin III has south direction.

Basins IV and V: The 3ld order channels of basins IV and V présent western orientation.

• 2nd order basins:

As I hâve mentioned before, the west part of the Mayda Insula consists of many drainage basins of 2nd order (Fig. 5.8). Next section is dedicated to these basins.

5.5.3

Statistical analysis of the angular distribution of the island's 2nd order streams in poor drained basins

As discussed above, the 2nd order streams are poorly drained in the Island of Mayda Insula (Fig. 5.8). In the other part of the island, except for the 3rd order basins, the 2nd order streams hâve been developed vertical to the coastline. The orientation of these streams is listed in Table 9, according to the N-S direction. The angle distribution of Table 5.13 is designed in Fig. 5.13, with the normal curve. There are two prominent orientations of the 2nd order channels in poorly drained basins. The Northern streams are oriented at 80°W and the Southern streams at 60°E. Comparing the two histograms of the Mayda Insula (Figs. 5.10 and 5.13), the streams in poorly drained régions seem to be related to the axis of the mountain ridges of the surface.

5.6

Cassini/RADAR image processing Titan's lakes were First recorded by the Cassini Titan Radio Détection and ranging Mapper -RADAR (Stofan et al., 2007). The RADAR instrument onboard the Cassini orbiter is designed to map Titan's surface as well as other Satumian moons, its rings and Satum itself.

The main RADAR advantage is that the emitted microwave radiation pénétrâtes the thick atmosphère almost unaffected. The idea is to combine RADAR recorded echoes with other remote sensing observations, especially VIMS, in order to retrieve a full view of Titan's surface. HGA is pointed to Earth for tracking, the RADAR is not operating. Five antenna microwave beams, one at the time (B1 to B5 in Fig. 5.12) are used in order to obtain the maximum path coverage. The RADAR transmits a set of puises for a given time period and when the first écho burst retums the receiver is switched on. During each sequence the instrument collects 1 GB of data. These data are compressed to a spécifie format the Block Adaptive Quantizer (BAQ) (Elachi et al., 2004).

Figure 5. 12 -Cassini RADAR antenna beams. The top panel shows the nadir pointing for altimetry, while the bottom panel shows the imaging side-pointing (Elachi et al., 2004).

SAR can probe the surface with two resolutions: the High SAR Resolution (HiSAR) at 350-720 m per pixel and the low SAR resolution. The linear polarized electric field vector of the Cassini/RADAR is oriented to be approximately parallel to Titan's surface during SAR operation [START_REF] Hayes | Transient Surface Liquid in Titan's Polar Régions from Cassini[END_REF] Figure 5. 13 -The sequence of the radar operational stages during a Titan flyby with closest approach at an altitude of 1000 km (Elachi et al., 2004).

The combination of the five individually illuminated sub-swaths dérivés the total width of RADAR swath and ranges from 120 to 450 km at the spacecraft altitudes from 1000 to 4000 km. Each nominal Titan flyby produces a SAR swath 5000 km long, which is about 1.1% of the satellite's surface. When the flyby has C/A at higher than 4000 km, no SAR images acquisition occurs (Elachi et al., 2004). However, the backscatter recordings suffer from both speckle and instrument noise.

The speckle noise is a multiplicative granular noise that appears when the surface features are rough compared to the illuminating cohérent radiation [START_REF] Goodman | Some fundamental properties of speckle[END_REF]. It is caused by the constructive and destructive interférence of the RADAR retums, which are scattered by many elementary reflectors [START_REF] Lee | Speckle filtering of synthetic aperture radar images: A review[END_REF][START_REF] Patel | Component-based restoration of speckled images[END_REF]. The speckle noise exists in ail types of cohérent imaging Systems and it reduces the resolution of the image as well as the detectability of the target [START_REF] Goodman | Some fundamental properties of speckle[END_REF]. Its presence in SAR images overlays real structures and causes grey value variations even in homogenous image parts, making automatic segmentation difficult.

When suppressing multiple views of the same surface element, the speckle noise follows a Gaussian probability distribution (Chitroub et al., 2002). The multilook technique averages incoherently the independent neighboring pixels to estimate the characteristics of the same ground area.

5.7

Filtering and segmentation of the Cassini/SAR images on Titan

The liquid of the Titan exhibits very low reflectivity and therefore the lakeshores can be easily distinguished from the surroundings. Long-term observations of the lakeshores' évolution may show temporal variations of the liquid volume and provide evidence about the methane cycle. However, the automatic segmentation of the shores becomes problematic by the existence of the RADAR speckle noise.

In order to reduce this speckle noise in Cassini/SAR images, we hâve applied a despeckle filter that yields the SAR restored images. In the next two subsections, I describe the application of the filtering procedure in SAR images as well as a segmentation method.

The filtering technique, based in probabilistic methods, is described in (Bratsolis & Sigelle, 2003) and the application we hâve performed in Cassini/SAR images in [START_REF] Bratsolis | A despeckle filter for the Cassini synthetic aperture radar images of Titan's surface[END_REF] which I hâve co-authored.

Total Sum Preserving Regularization (TSPR) filtering

The filter is based on a membrane model Markov random field approximation optimized by a synchronous local itérative method, the outcome of which is a total sum preserving regularization (TSPR) for the pixel values of the image. The image is considered as a random element drawn from a pre-specified set of possible images. The problem is ill posed, since we try to recover the real image from the recordings. An itérative algorithm is then applied to optimize the image reconstruction.

To retrieve the maximum information of the

Supervised segmentation

The TSPR filter can be used as intermediate stage for the extraction of meaningful régions that correspond to structural units in the scene or distinguish objects of interest. The speckle noise has not been completely removed ffom the filtering and we apply an image segmentation method in order to divide the image into spécifie régions where the pixels share common properties.

The segmentation method classifies régions of an image following a spécifie criterion.

The supervised method of minimum Euclidean distance uses the mean values of each member and calculâtes the Euclidean distance from each classified object to the nearest class segmenting the image into different régions of interest or different labels.

In our case, we hâve selected three classes: the dark lakes (black), the granular lakes (grey) and the background (light grey). For each value of normalized backscatter, we subabstract the mean values that correspond to the région of interest. After using the despeckling filter TSPR we apply the segmentation method. The resulted image is illustrated in Figure 5.16 below. The dark part of lakes are indicated in blue, in green it is the granular part of lakes and in red the local background. The temporal variation of the dark spots can provide information on the évolution of the lake System and consequently help us to better understand the methane cycle on Titan (Atreya et al., 2006) and therefore the mechanisms linked with the lake surface features, their origin and fate, through a global temporal and spatial coverage [START_REF] Hayes | Transient Surface Liquid in Titan's Polar Régions from Cassini[END_REF]. The proposed filtering and segmentation method would be a help fui tool in enhancing the retum of the analysis of ail SAR data acquired on Titan and other objects as well as in the exploitation of such data from future missions to Titan.

I apply the TSPR filter and supervising segmentation procedure on a different surface feature where the contrast between the RADAR echoes is not obvious. I hâve selected the SAR image of Sinlap crater (11.3°N, 16.0°W)-PIA 07368 (Fig. 5. 17-left). This crater has a diameter of about 60 km and the bright surrounding material indicates impact origin. The central panel of Figure 5.17 shows the results of the despeckle filtering and the supervised segmentation procedure for the Sinlap crater. The segmentation outcome of our filtering procedure in Sinlap crater (Fig. 5.17-left) shows régions of different surface types according the backscatter recording after the speckle removal.

Future perspectives

With the filtering and segmentation procedure described above, accurate isolation of distinct surface features from their surroundings can be achieved. As we can see from Figures 5.16 and 5.17-right, the shapes of the surface édifices can be determined and then classified.

The technique I described above removes the speckle noise from RADAR images and provides the shapes of régions of interest. We can then optimize the results of the direct comparison of Cassini/RADAR images with the images retrieved from ISS and VIMS of the same area. Several similar efforts hâve been reported so far (Bames et al., 2007a;Soderblom et al., 2007a;[START_REF] Mouelic | Mapping and interprétation of Sinlap crater on Titan using Cassini VIMS and RADAR data[END_REF][START_REF] Tosi | Corrélations between VIMS and RADAR data over the surface of Titan: Implications for Titan's surface properties[END_REF].

One of the most challenging projects I plan to be engaged in is to superimpose VIMS images in which the atmospheric contribution is constrained by M. Hirtzig's radiative transfer code over the despeckled and segmented RADAR images of the same région. We hâve already been cooperating with Dr. R. Lopes to this direction.

In book sériés of Springer Editions, edited by Prof. J. Seckbach20 [START_REF] Coustenis | Life in the Satumian Neighborhood[END_REF]. In this work, we look for similarities as well as différences among the aforementioned satellites, 20 http://www.springer.com/series/5775 which may impact their habitability conditions. Our main purpose in this book chapter is to underline the presence of liquid water considerably further from the Sun than previous models of the Habitable Zone envisaged and to emphasize that both Titan and Enceladus host unique conditions for the emergence of biological building blocks in their environment and to recommend a future mission dedicated to these moons to perform astrobiological studies.

Comparative planetology among the icy moons of the outer planets allows us to better bring forward common characteristics of the atmosphère and the surface that allow us to look into the origin and évolution of these bodies. On this subject, we hâve written a paper which was published in the Journal of Cosmology [START_REF] Solomonidou | Water Océans of Europa and Other Moons: Implications For Life in Other Solar Systems[END_REF].

Accordingly, in this chapter, I describe the habitability of the icy moons of our Solar System more from the astronomical point of view and looking at them as potential habitats.

Introduction and context

The scientific field of Astrobiology (also bioastronomy or exobiology) investigates the origin and the évolution of life on Earth in the past, présent and future. Extending to the terrestrial analogues, Astrobiology explores the possibility of life forms existing in extraterrestrial environments hosting suitable conditions for life emergence and/or sustainability.

Astrobiology adapts knowledge obtained from different scientific disciplines namely astrophysics, geology, physics, chemistry, geochemistry, biology and more, in order to understand the appropriate aspects conceming the création of our Solar System and the initiation of life. It focuses on extraterrestrial environments, posing the unanswered question on the origin of life on Earth and elsewhere and how humans can detect it (Des [START_REF] Marais | Remote sensing of planetary properties and biosignatures on extrasolar terrestrial planets[END_REF]Raulin, 2007).

Planetary habitability is the ability of a planetary environment to support and sustain life forms. The habitability potential of a planet or a satellite dépends on a combination of factors, which are considered to be essential for life appearance, évolution and maintenance.

Crucial factors are, among other, the orbital properties of the planetary body, its stability, its bulk composition, the existence of an atmosphère and a surface, as well as the proper Chemical ingrédients.

The région around a star or a planet that is favorable for life in any form is called habitable zone (HZ) [START_REF] Huang | Occurrence of Life in the Universe[END_REF]. Within the habitable zone a planetary environment can hâve the required température, pressure and luminosity conditions to allow water to remain in a stable liquid State on its surface [START_REF] Hart | Habitable zones about main sequence stars[END_REF]. In such a zone, living organisms may arise and evolve.

Life emergence on Earth sets the habitability constraints. Terrestrial life is the final product of a long complex Chemical évolution, requiring at least four raw ingrédients: (a) the existence of liquid water, (b) a stable environment (c) carbonaceous matter as nutrients and

(d) energy. The fulfillment of these prerequisites over a long period of time can be considered as an indicator of suitable environments for hosting the proper biological building blocks, which may lead to the formation of primitive life structures [START_REF] Kasting | Habitable Zones around Main Sequence Stars[END_REF].

The basic définition of the habitable zone refers to the existence of liquid water on the surface. In this framework, the boundaries of the habitable zone of the Solar System (Fig. 6.1), extend from 0.95 AU to 1.2 AU.

The inner boundary of the habitable zone is the spécifie distance from the star where the surface liquid water of a planetary body has been vaporized due to runaway greenhouse effect. In this case, water is vaporized rapidly and dissociated by solar UV radiation in the upper atmosphère. Subsequently, hydrogen under high température conditions, escapes to space, breaking the water cycle [START_REF] Kasting | Runaway and moist greenhouse atmosphères and the évolution of earth and Venus[END_REF]. The outer limit of the habitable zone is the spécifie distance from the star where the local atmospheric conditions hâve failed to produce a satisfactory greenhouse effect that could keep the température at the planetary atmosphère above the freezing point (Lammer et al., 2009). Altematively, it is the distance from the star where CO2 begin to condense [START_REF] Kasting | Habitable Zones around Main Sequence Stars[END_REF], but CO2 ice clouds in the atmosphère still allow for the surface température to be above the freezing point of water [START_REF] Forget | Warming Early Mars with Carbon Dioxide Clouds That Scatter Infrared Radiation[END_REF].

However, what makes a planetary body habitable is more complex than fulfilling the distance limitation. Lammer et al., (2009) defines 4 classes of habitable planets. The first Class I consists of Earth-type planets in which life can evolve following the terrestrial analogue. These planets host the proper atmospheric conditions to maintain their surface liquid water stable over a long period of time in geological tenus, in combination with active plate tectonics.

Class II habitats are planets in which life may evolve like in Class I, but due to different stellar and geophysical conditions, they hâve followed different evolutionary paths [START_REF] Taylor | Comparative planetology, climatology and biology of Venus, Earth and Mars[END_REF]. Early Venus and early Mars belong to this class. According to Pioneer Venus and Venus Express measurements, Venus [START_REF] Kulikov | Atmospheric and water loss from early Venus[END_REF][START_REF] Svedhem | Venus as a more Earth-like planet[END_REF] which lies below the inner threshold of the habitable zone, at 0.72 AU, has kept its surface water during its early stage, before experiencing a runaway greenhouse effect [START_REF] Kasting | Runaway and moist greenhouse atmosphères and the évolution of earth and Venus[END_REF][START_REF] Cockell | Life on Venus[END_REF]. Mars orbits at 1.52 AU and managed to stay warm between 3.5 and 4 billion years ago to allow for the existence of liquid water on the surface [START_REF] Craddock | The case for rainfall on a warm, wet early Mars[END_REF][START_REF] Ehlmann | Subsurface water and clay minerai formation during the early history of Mars[END_REF] producing characteristic landforms such as océan shorelines, fluvial channels and valleys [START_REF] Carr | Océans on Mars: An assessment of the observational evidence and possible fate[END_REF][START_REF] Mcewen | A Closer Look at Water-Related Géologie Activity on Mars[END_REF].

Liquid water can also exist under the surface in several planetary bodies, below their icy crust, as in the case of the icy satellites of the outer planets. According to Lammer et al., (2009), the planetary bodies where the internai océan interacts directly with silicates, like Jupiter's Europa, belong to the Class III habitats. In this structure, the silicates provide various Chemicals and energy through hydrothermal or volcanic activity.

The Class IV habitats are the planets or satellites which hâve huge water underground deposits without interacting with silicate material. These subsurface océans of the icy satellites of the gas giants may be in direct contact with heat sources below their icy crust or encapsulated between two ice layers, or liquids above ice. This is the case of Jupiter's Ganymede and Callisto and Satum's Titan. In this context, the détermination of the habitable zone limits can be expanded. The question is if these underground water deposits can support any life forms. Except for water liquid, life requires other essential éléments such as nitrogen and phosphorus in addition to hydrogen and oxygen. If this situation maintains for long timescales, the liquid underground water may become capable of sustaining life. Similar conditions hâve been considered for the terrestrial case where a coupled sea/ice System could provide the necessary conditions for life emergence on the primitive Earth (Trinks et al., 2005). In the paper by [START_REF] Solomonidou | Water Océans of Europa and Other Moons: Implications For Life in Other Solar Systems[END_REF], which I hâve co-authored, we review the surface features of Ganymede, Europa, Titan and Enceladus, their internai structure and their astrobiological potential.

The rôle of a planetary field is also crucial for habitability. A strong intrinsic magnetic field protects the atmosphère from the direct interaction with the solar wind and thus, the quick loss through sputtering process is avoided. Additionally, the magnetic field traps highenergy particles into radiation belts, avoiding harmful radiation levels on the surface. The magnétopause of a magnetized planetary body, like the Earth or the possibly early Mars and early Venus, deflects the solar wind plasma globally except for the polar régions, where the solar plasma reaches the ionosphère. In the case of an unmagnetized (or weakly magnetized) body, like the Venus, Mars and Titan, the ionosphère shields the surface from the solar wind flow (e.g. Lammer et al., 2009). (image crédit: http://media.egu2012.eu/media/filer_public/2012/04/19/eguJuice_press.pdf).

Titan, Enceladus and Europa host the proper ingrédients for developing (today or in the past) life-friendly conditions and, therefore, the interest for an astrobiological investigation of these planetary objects is enormous. Since 2004, the discoveries of the Cassini-Huygens mission in the Satumian System hâve revolutionized our considérations as to whether these bodies could harbor life (now or in the future) or at least provide us with valuable information on the origin and évolution of life on Earth.

Titan as a possible habitat

Titan is a good candidate for astrobiological studies as its environment exhibits many similarities with the Earth's. Current investigations hâve shown that Titan fulfïlls many of life's prerequisites for an organic portfolio due to its atmosphère. Both the remote observations from the Cassini orbiter and the in situ measurements from the Huygens probe hâve significantly advanced our perspective of Titan's System and its potential to harbor the essential ingrédients for life. Titan is indeed a very complex world much like Earth (Fig. 6.2). A thick nitrogenbased atmosphère, four times denser than the terrestrial one, where a rich organic chemistry occurs, shrouds its surface. Besides this rich organic budget, we also find a geologically active surface including lake-like features filled with organic liquid (Stofan et al., 2007).

Additionally, past models and recent discoveries reveal that Titan probably contains a vast subsurface liquid water océan (Lorenz et al., 2008c;Beghin et al., 2009a;[START_REF] Iess | The Tides of Titan[END_REF].

From the outer limits of its atmosphère, to deep in its interior, Titan is a living planetary body with an evolving environment. Titan, along with cornets, is considered as a planetary scale natural laboratory for prebiotic chemistry and a possible habitat for extraterrestrial life (albeit probably different from the terrestrial one). Studying Titan could be bénéficiai to our understanding of the origin of life on Earth. The low solar influx, the composition of Titan's atmosphère and the possible presence of an internai water océan, give us the opportunity to study the conditions prevailing on the primitive Earth.

Titan's prebiotic atmospheric environment

Titan can be considered as a huge abiotic factory. The total organic abundance in the atmosphère, the lakes and dunes exceeds the carbon inventory in the Earth's globe (océan, biosphère and fossil fuel réservoirs) by more than one order of magnitude (Lorenz et al, 2008b).

Its atmosphère consists mainly of nitrogen (98%) and methane at about 2.2% in the upper layers [START_REF] Waite | Ion Neutral Mass Spectrometer Results from the First Flyby of Titan[END_REF] Plasma Spectrometer (CAPS) data strongly suggests that high molecular weight species (up to several 1000 Daltons) may be présent in the ionosphère (Waite et al., 2007;Vuitton et al., 2009). These observations open new avenues in our perception of the organic content and chemistry occurring in Titan's atmosphère. These detected compounds initiate the process of haze formation starting at about 950 km (Waite et al., 2007) to finally condense out on the surface after descending through the atmosphère.

Through photodissociation in this dynamic evolving environment, the second most abundant atmospheric constituent, methane, is destroyed irreversibly producing a variety of trace gases such as hydrocarbons (e.g. ethane, acetylene, and propane). Furthermore, CH4 in combination with nitrogen, gives rise to nitriles like hydrogen cyanide, acetonitrile, and cyanoacetylene. Many of these organic products hâve been detected in the stratosphère (between 70 and 500 km in altitude) by the Cassini Composite Infrared Spectrometer -CIRS (e.g. Flasar et al., 2005, Coustenis et al., 2007;2010), which has confirmed the Voyager 1/IRIS observations [START_REF] Hanel | Infrared observations of the Satumian System from Voyager 1[END_REF][START_REF] Kunde | C4H2, HC3N and C2N2 in Titan's atmosphère[END_REF][START_REF] Maguire | C3FL and C3H4 in Titan's atmosphère[END_REF][START_REF] Kim | The abundance of CH3D in the atmosphère of Titan, derived from 8-to 14micron thermal émission[END_REF][START_REF] Samuelson | CO2 on Titan[END_REF]Coustenis et al., 1989a;1995). These space détections hâve complemented ground-based (e.g. [START_REF] Hidayat | Millimeter and Submillimeter Heterodyne Observations of Titan: Retrieval of the Vertical Profile of HCN and the 12C/13C Ratio[END_REF][START_REF] Tanguy | Stratospheric profile ofHCN on Titan from millimeter observations[END_REF][START_REF] Kostiuk | Ethane abundance on Titan[END_REF], Roe et al., 2004) and space-bom discoveries of water vapor [START_REF] Coustenis | Evidence for water vapor in Titan's atmosphère from ISO/ SWS data[END_REF] and acetonitrile [START_REF] Bezard | Détection ofAcetonitrile on Titan[END_REF]. The Chemical inventory in Titan's atmosphère further includes oxygen compounds like CO (Lutz et al., 1983) and CO2 [START_REF] Samuelson | CO2 on Titan[END_REF]. Moreover, the presence of water vapor (Cottini et al., 2012a) and benzene (Coustenis et al., 2007;2010b) has been unambiguously confirmed by the CIRS instrument after the initial ISO discoveries (Coustenis et al., 1998;2003).

The CIRS instrument investigates Titan's atmosphère since 2004 and will continue to do so up to 2017, when the Cassini Solstice mission ends. The atmospheric Chemical composition provided by CIRS sets up constraints for astrobiological studies. Some of these gases can be considered as key molécules in terrestrial prebiotic chemistry, such as hydrogen cyanide (HCN), cyanoacetylene (HC3N) and cyanogen (C2N2). Moreover, HCN is considered as a prebiotic molécule, a precursor of life (see [START_REF] Raulin | Prebiotic Chemistry: Laboratory Experiments and Planetary Observation[END_REF] and references therein).

Determining the origin and évolution of these molécules on Titan could provide us dues on how life began on Earth.

Ail of the organic compounds, detected in Titan's atmosphère, were also produced in simulation prebiotic experiments. The first simulation of the reactions that took place in the primitive terrestrial atmosphère was the Miller-Urey experiment [START_REF] Miller | A Production of Amino Acids under Possible Primitive Earth Conditions[END_REF][START_REF] Miller | Organic Compound Synthesis on the Primitive Earth[END_REF]. Miller modeled the Early Earth atmosphère in a glass reactor with a mixture of methane, ammonia, hydrogen and water vapor in the presence of liquid water at 1 bar, following Urey's assumption. He applied electrical sparks to simulate lightning flashes which might hâve occurred in the early troposphère. The outcome of this experiment was the formation of key organic compounds like formaldéhyde and hydrogen cyanide which are considered as precursors of biological amino acids. More complex experiments were conducted later by [START_REF] Sagan | Long-Wavelength Ultraviolet Photoproduction of Amino Acids on the Primitive Earth[END_REF] using UV radiation as the energy source. Two types of products were formed (see [START_REF] Raulin | Prebiotic Chemistry: Laboratory Experiments and Planetary Observation[END_REF] and references therein): simple volatile organic compounds (i.e. HCN, HC3N, HCHO) which can participate in prebiotic reactions and macromolecular products, refractory organics, usually named "tholins" (Sagan & Khare, 1979). Tholins represent laboratory analogues of Titan's aérosols and are useful to interpret many observational data which require information on the aérosols. As experimental analogues of Titan's atmospheric particles, tholins also permit the study of the behavior of the aérosols in Titan's conditions in the laboratory. These organics release biological amino acids in an aqueous solution (Khare et al., 1986).

Titan lacks liquid water surface deposits and therefore is an idéal environment to study the processes of Chemical évolution under anhydrous conditions (Raulin & Owen, 2002).

Several laboratory experiments hâve shown the existence of hydrocarbon and nitriles [START_REF] Thompson | Plasma discharge in N2 + CH4 at low pressures -Experimental results and applications to Titan[END_REF][START_REF] De Vanssay | Experimental simulation of Titan's organic chemistry at low température[END_REF][START_REF] Coll | Review and Latest Results of Laboratory Investigations of Titan's Aérosols[END_REF]1999). Schaefer and Fegley (2007) predict that Earth's early atmosphère contained CH4, H2, H20, N2, and NH3, similar to the components used in the Miller-Urey synthesis of organic compounds, often related to Titan's and Enceladus' atmospheric inventory.

Furthermore, according to Trainer et al. (2006), the processes that formed the haze on Titan and on early Earth hâve many similarities with what could hâve served as a primary source of organic material to the surface.

The Huygens Aérosol Collector Pyrolyser (ACP) experiment provided the First direct in situ Chemical analysis of Titan's aérosols. The instrument collected haze particles in Titan's stratosphère and troposphère during the probe's descent phase, heated them at different températures and sent the produced gases for analysis to Huygens Gas Chromatograph-Mass Spectrometer (GC-MS). The results thus obtained indicated that the aérosols are made of a refractory nucléus, composed of H, C and N atoms (Fig. 6.3), and producing NH3 and HCN after pyrolysis at 600 °C (Israël et al., 2005). ACP results strongly support the tholin hypothesis (Nguyen et al. 2007 and references therein). They also suggest that Titan's aérosols may evolve if in contact with water ice on its surface, and could produce a variety of organics of biological interest, such as amino acids (Neish et al., 2010;[START_REF] Ramirez | The fate of aérosols on the surface of Titan[END_REF]. However, Huygens/GCMS did not detect a large variety of organic compounds in the low atmosphère (Niemann et al., 2005). Moreover, we do not know anything yet about the nature and quantity of the condensâtes, nor of the elemental composition or the molecular structure of the reffactory part of the aérosols.

The Cassini-Huygens mission has revealed the details of the organic and methane hydrologie cycles on Titan (Atreya et al., 2006;Raulin, 2008a;Lebreton et al., 2009).

Methane on Titan plays the part of water on Earth. It can exist in ail three forms : as a gas, liquid and solid, since the mean surface température is almost 94 K (Fulchignoni et al., 2005), When methane condenses in the form of haze, it is done irreversibly (Fig. 6.4). Taking into account the rate of methane photo-dissociation and the formation of organic products in Titan's atmosphère, methane should therefore hâve disappeared after 10-100 million years (Atreya et al., 2006). As a conséquence, we are either witnessing the extinction of methane or -more likely -looking at a source, which replenishes the methane gaseous abundance in the atmosphère.

approaching
Today, on Earth, methane is mainly a by-product of the metabolism of many living organisms. A legitimate question may be if the same occurs on Titan. Huygens/GCMS in situ measurements hâve shown that the carbon-13 isotopic ratio from methane is compatible rather with inorganic values methane, and therefore Titan's methane is not considered of biogenic origin (Niemann et al., 2005;2010). Although the sinks of atmospheric methane on Titan are relatively well understood, the major sources of replenishment in the atmosphère are still very model-dependent.

Comparatively to the Earth, Titan lacks both oxygen and hydrogen. However, with several percent of methane in dinitrogen, the atmosphère of Titan is one of the most favorable atmosphères for prebiotic synthesis. Analogies are obvious between the current organic chemistry occurring on Titan and the prebiotic chemistry which was once active on the primitive Earth, prior to the emergence of life (e.g. McKay & Smith, 2005). Although liquid water extents are absent on Titan's surface, both chemistries are similar.

A thick methane-induced organic haze was possibly formed on early Earth. Before the rise of the atmospheric oxygen in the terrestrial atmosphère 2.5 Gyrs ago, it is considered possible that the abundance of methane gas was 10 to 20 times higher than the today's value of l.ôxlO'6 [START_REF] Pavlov | Methane-rich Proterozoic atmosphère?[END_REF]. If the atmospheric CO2/CH4 ratio had become equal to 10 at the mid-Achaean era, an organic haze could hâve formed on this early environment (Pavlov et al., 2000;DeWitt et al., 2009). This hydrocarbon haze produced the anti-greenhouse effect which reduced the température of the atmosphère (Kasting & Howard, 2006). Titan expériences the same anti-greenhouse effect (McKay et al., 1999). The absence of vast amounts of CO2 on Titan is one of the major différences between the two atmospheric envelopes. On the other hand, hydrogen cyanide and other prebiotic molécules are among the starting materials for biosynthesis. The existence of hydrocarbons, and in particular acetylene and benzene, has really enlarged the borders of photochemical organic products to a degree that is still unknown today on Titan.

Especially, the presence of benzene (CôHô) is extremely interesting, as it is the only polycyclic aromatic hydrocarbon (PAHs) discovered on Titan today. PAHs could contribute to the synthesis of biological building blocks in liquid deposits on Titan's surface in low température. Recent laboratory experiments confirmed that aromatic compounds could be plausibly produced on icy surfaces (Menor-Salvân et al., 2008).

Benzene was first detected at 674 cm'1 in Infrared Space Observatory (ISO/SWS) data (Coustenis et al., 2003) with a mixing ratio of 4x10'10. It was then also detected in the thermosphère (950-1150 km) from the analysis of Cassini/INMS data (Waite et al., 2007) and firmly in the stratosphère (100-200 km) at ail latitudes by Cassini/CIRS (Flasar et al., 2005;Coustenis et al., 2007;2010b). Moreover, It has been tentatively identified on Titan's surface by Huygens/CGMS (Niemann et al., 2005).

Titan lacks oxygen and sufficiently high températures compared to the primitive Earth. The abundances of hydrocarbons are higher on Titan than those on Earth by a factor of about 102-104. Apparently, evolutionary pathways different from the ones operating on Earth must hâve been followed on Titan and polyphenyls could possibly be created (Delitsky & McKay, 2010).

Titan's active terrestrial-like surface

Titan is an evolving planetary body with several terrestrial shaped surface features distributed globally. The Cassini-Huygens mission has observed a small number of impact craters [START_REF] Wood | Impact craters on Titan[END_REF], which implies an active surface environment that erased almost ail traces of past activity and records. Indeed, as on Earth, sedimentological and meteorological processes, mark the surface of Titan. Erosional processes hâve been reported on Titan's surface (Jaumann et al., 2008;Lopes et al., 2010). Huygens/DISR recorded well developed rivers, like dedritic networks, close to the probe's landing site at 10°S and 192°W [START_REF] Tomasko | Rain, winds and haze during the Huygens probe's descent to Titan's surface[END_REF], see Fig. 6.5.

Figure 6. 5 -Channels on Titan's surface (black lines) from Huygens/DISR images during the probe's descent [START_REF] Tomasko | Rain, winds and haze during the Huygens probe's descent to Titan's surface[END_REF].

Particles from the complex layering of organic aérosols in Titan's atmosphère fall down and deposit on the surface. Cassini/RADAR observations link this aérosol rain to the génération of extensive organic dunes, which shape an équatorial belt (Fig. 6.6). Titan's dunes are believed to be composed of ice and organics grains that possibly dérivé from a combination of the surface ice and the organic aérosols dropped from the atmosphère [START_REF] Lorenz | The Sand Seas of Titan: Cassini RADAR Observations of Longitudinal Dunes[END_REF]Radebaugh et al., 2008). Hence, the upper atmosphère is linked intimately with the surface and the intervening atmosphère. Despite the low surface température at about 94 K (Fulchignoni et al., 2005;[START_REF] Jennings | Titan's Surface Brightness Températures[END_REF], the organics that accumulate on the surface can evolve once in contact with water ice and may form organic molécules of astrobiological interest.

Cassini/RADAR unveiled large liquid deposits (Fig. 6.7) on Titan's surface distributed at polar régions (Stofan et al., 2007;Hayes et al., 2008;2010) and the equator [START_REF] Griffith | Possible tropical lakes on Titan from observations of dark terrain[END_REF]. for comparison (Hayes et al., 2008).

With its lakes and seas, Titan is the only body in the Solar System hosting large liquid bodies on its surface except for the Earth. These very dark radar spots were finally proved to be filled with liquid, most probably with ethane rich mixtures (Brown et al., 2008). The features range in size from less than 10 km2 to at least 100,000 km2.

The cosmic rays that reach Titan's surface hâve a quite low flux [START_REF] Raulin | Organic chemistry on Titan versus terrestrialprebiotic chemistry: Exobiological implications[END_REF], which means that, in combination with the low température, the reaction rate in the upper layers of the liquid remains very low. The surface liquid deposits can then offer the proper stable environmental conditions for prebiotic chemistry which dépends on the long duration, the ffeezing degree, the dissolved organics and the sedimentary déposition to the lakes' bottom [START_REF] Tokano | Limnological Structure of Titan's Hydrocarbon Lakes and Its Astrobiological Implication[END_REF]. Pure methane lakes develop and freeze in short geological periods of time, and, therefore, such lakes hâve no prebiotic significance. Instead, deep lakes (> few hundred meters deep), consisting of a mixture of ethane, methane and nitrogen [START_REF] Cordier | An Estimate of the Chemical Composition ofTitan's Lakes[END_REF]2012), can favor stable composition at the bottom with accumulation of acetylene, as the latter sinks being denser than the liquid. The deeper layers do not show any significant movement and create an isolated environment with the proper constituents for prebiotic chemistry.

Cassini Imaging Science Subsystem (ISS) images hâve allowed for the compilation of a nearly global surface map and the monitoring of both the surface and atmosphère for activity. Repeated south-polar imaging by ISS revealed différences consistent with ponding of hydrocarbon liquids on the surface due to précipitation from a large storm [START_REF] Turtle | Rapid and Extensive Surface Changes Near Titan's Equator: Evidence of April Showers[END_REF].

The large lakes or seas on Titan's surface are primary targets of future lake landers. I hâve proposed a spécifie experiment to be incorporated as payload in a future lander mission.

The concept of this experiment is to use a sériés of micro-probes in order to sense the deep isolated layers of the liquid deposit as well as its floor [START_REF] Bampasidis | Sounding the interior of Titan's lakes by using Micro-Electro-Mechanical Systems (MEMS)[END_REF]. The experiment is fully described in the next Chapter.

Titan's internai océan

The source of the methane replenishment may lie beneath the lakes in an underground "aquifer" (or "methanofer") System hinted at strongly by Cassini data. The existence of an océan layer below the surface is crucial for defining the habitability potential of the moon.

According to Lammer et al., 2009, due to this encapsulated water layer, Titan can be considered as Class IV habitat.

Remote sensing data from the Cassini orbiter imply that Titan's interior is partially differentiated. Indeed, Cassini Radio Science (RS S) data provided variations of its degree 2 coefficient gravitational potential, which supports internai density variations [START_REF] Rappaport | Can Cassini detect a subsurface océan in Titan from gravity measurements[END_REF].

The Cassini Synthetic-Aperture RADAR (SAR) images, acquired between the October 2004 and May 2007 flybys, revealed a 0.36°faster than synchronous rotation of the moon. Lorenz et al. (2008c) then suggested that the crust is decoupled from the interior by a subsurface water océan.

HASI-PWA measured the electrical conductivity of the atmosphère during the descent of the Huygens Probe through Titan's atmosphère in January 2005. The sensor detected an extremely low frequency (ELF) radio wave during the descent which was oscillating very slowly for a radio wave (36 Hz) and increased slightly in frequency as the probe reached lower altitudes [START_REF] Hamelin | Electron conductivity and density profiles derived from the mutual impédance probe measurements performed during the descent of Huygens through the atmosphère of Titan[END_REF]. The ELF waves on the Earth are reflected by both the surface of the Earth and its ionosphère. However, Titan's surface is a poor reflector because of its low conductivity and so these waves penetrate the interior. The recorded wave could hâve been reflected by the liquid-ice boundary of a subsurface océan of water and ammonia predicted by theoretical models [START_REF] Simoes | A new numerical model for the simulation of ELF wave propagation and the computation of eigenmodes in the atmosphère of Titan: Did Huygens observe any Schumann résonance[END_REF]. [START_REF] Beghin | New insights on Titan's plasma-driven Schumann résonance inferred from Huygens and Cassini data[END_REF]2012) interpreted the ELF wave as a Schumann résonance between the ionosphère and a modestly conducting océan at 30-50 km below the surface.

Cassini's accélération measurements during flybys from 2006 to 2011 discovered large solid tides caused to Titan by Satum at about 10 m in height. If the interior is uniform, such solid tides should be only at 1 m in height. Therefore, this fmding implies the existence of an internai océan of liquid water, assuming a depth of 100 km beneath the surface [START_REF] Iess | The Tides of Titan[END_REF].

Titan internai structure models also support the presence of such an internai océan (Grasset & Sotin, 1996;Grasset et al., 2000;Tobie et al., 2005;Mitri et al., 2008). Thermal évolution models suggested that Titan may hâve an icy crust between 50 and 150 km thick, lying over a liquid water océan, a couple of hundred kilometers deep, with some amount (a few to 30%, but most likely -10%) of ammonia dissolved in it, acting as an antifreeze material. This corresponds to a pH of around 11.5. The pressure reaches -5 kbar at 200 km depth, and it could include hot spots reaching -20°C. The possible Titan's internai structure from the interior to the surface is as follows: (a) a silicate core, (b) a high-pressure water ice layer, (c) an ammonia-rich océan and (d) an ice layer covered by a crust of organics and ices (Tobie et al., 2005;Fortes et al., 2007;[START_REF] Fortes | Titan's internai structure and the evolutionary conséquences[END_REF], see Figure 6.8. However, the exact inner structure of Titan can only be identified by seismic experiments. For this reason I support the installation of a seismic network on Titan's surface (Bampasidis et al., 2011a;[START_REF] Bampasidis | MEMS-based seismic sensors on Titan and Enceladus[END_REF].

Subsurface

Mostty Solid Surface Liquid Metbane Such an internai liquid océan with the presence of methane clathrates could exsolute and produce overpressure and subsequently the ammonia-water could erupt to the surface leading to explosive cryovolcanic phenomena (Fortes et al., 2007).

Cryovolcanism is an activity that resembles terrestrial volcanic processes following a similar mechanisms and patterns but in cold environmental conditions with different initial and depositional products. An ice volcano erupts liquids and gases from the interior of an ice shell, when they are denser than the surrounding solid. The required heat for melting the material cornes from tidal friction. On Titan, fluid water mobilized and made buoyant by ammonia and/or methane and/or other materials, is ejected from the interior to the surface.

Cryovolcanism is a dynamic process that links the interior, the surface and the atmosphère, while it is believed to be a significant source of the methane in the atmosphère on Titan (Tobie et al., 2006).

Cassini data support the presence of cryovolcanism on Titan since traces of cryomagma flows are believed to exist across parts of the surface (Sotin et al., 2005;[START_REF] Lopes | Cryovolcanic features on Titan's surface as revcaled by the Cassini Titan Radar Mapper[END_REF]. Several régions hâve been proposed to change réflectance on Titan's surface so far, with the three more often cited: Tui Regio (20°S, 130°W), Hotei Regio (26°S, 78°W) and Sotra Facula (12.5°S, 39.8°W) [START_REF] Bames | Cassini observations of flow-like features in western Tui Regio, Titan[END_REF]Soderblom et al., 2007a;Nelson et al., 2009a;2009b). These features are wide crater-like with lobate flows originate from them (Soderblom et al., 2009;Lopes et al., 2010).

The theory of methane clathrates trapped in an internai liquid océan enhances the astrobiological interest. According to the "clathrate gun" hypothesis, potential movement and rise of the température in an underground liquid deposit could trigger the sudden release of methane from its clathrate buried in permafrost or seabeds (Kennett et al., 2003) or an océan like on Titan's case. Further température enhancement leads to further methane clathrate destabilization which could easily trigger cryovolcanic éruptions (Kennett et al., 2003). Such process on Titan could increase the température values, creating an environment more favorable for life to exist.

6.2.4

Is it possible to find life forms in Titan's environment?

Water is frozen solid on Titan's surface and much too cold to support terrestrial-type life. The list of liquid candidates, présent on the surface, includes liquid methane and related molécules like ethane. While liquid water is widely regarded as necessary for life, it is not a strict requirement. The new hydrogen findings on Titan are consistent with conditions that could produce an exotic, methane-based life form, but do not prove its existence.

McKay and Smith (2005) assumed the possibility for a different form of life to exist in liquid hydrocarbons of Titan. Such a methanogenic life form could consume H2 instead of O2 that could be measured in the lower atmosphère. Towards that hypothesis, Strobel (2010) and According to Strobel (2010), hydrogen flows down through Titan's atmosphère and then disappears on the surface, resembling the terrestrial oxygen consumption. On Titan, the hydrogen as byproduct of UV dissociation in the upper atmosphère should be distributed fairly evenly throughout the atmospheric layers. Strobel (2010) described a downward flow of hydrogen molécules to the surface at a rate of about 10,000 trillion molécules per second, similar to the escape rate. Hydrogen cannot be stored in an underground space on Titan. The surface is too cold to drive Chemical processes with a catalyst to convert hydrogen molécules and acetylene back to methane. Only an unknown minerai could provide the net energy needed for such reactions, acting as the catalyst on Titan's surface. The McKay and Smith ( 2005) assumption for methanogenic life, which consumes the hydrogen from the atmosphère, can explain its disappearance.

Methane-based life forms hâve not been yet detected anywhere, though there are liquid-water-based microbes on Earth that thrive on methane or produce it as a waste product.

At Titan's low températures, a methane-based organism would hâve to use a substance that is liquid as its medium for living processes, but not water itself.

The effectiveness of methane as a medium for life compared to water or ammonia is under considération. Water enables easier transport of substances in a cell than methane does, since it has higher solubility. On the contrary, methane's lesser Chemical reactivity favors the easier formation of large structures corresponding to proteins (Benner et al., 2004). In addition, the possible existence of cryovolcanic activity suggests higher températures within the océan and the volcanic conduit where heat transfer between the interior and upper layers would be critical in creating and sustaining any kind of subsurface oceanic life (Grasset et al., 2000).

Concluding remarks for Titan's astrobiological case

Taking into account ail the characteristics described above, the prebiotic potential of Titan is significant. Titan seems to be an idéal planet-size laboratory for increasing our knowledge of the évolution of the Earth 's atmosphère. However, the full extent of current géologie activity is still under investigation.

As we hâve mentioned in the Book Chapter by Coustenis at al. (2012), Titan, although it présents many similarities to the primitive Earth, should not be considered as a frozen analogue. It is an evolving planetary body with an active complex chemistry occurring from the upper atmosphère to its surface. One of the by-products of this chemistry, the macromolecular material, is included in atmospheric refractory organics, the Titan 's tholins, which are eventually deposited on the solid surface (dunes) or in the lakes. Both laboratory and theoretical works hâve provided an assumption of the composition of these compounds.

Laboratory experiments show that the accumulation of these aérosols on the surface, if they are in contact with liquid water, can release many compounds of biological interest, such as amino acids (Khare et al., 1986). Similar processes could be particularly favorable in areas of Titan 's surface where cryovolcanism may occur.

Titan lakes fulfill the habitability criteria as far as the presence of carbonaceous material and the stable environment. McKay & Smith (2005) suggested that methanogenic life on Titan's surface liquids could dérivé the energy they need by consuming hydrocarbons. The isolated environment of Titan's deep lakes can offer the proper conditions to increase the concentration of high order organics in their deep layers and their floor.

Therefore, a dedicatedfuture mission to Titan 's lakes could, in principle, investigate the presence of organic material as well as the thermodynamic conditions of their environment If no vertical convection occurs within the deep layers of the lakes, the local environment is favorable for maintaining the formation of biological building blocks and their Chemical évolution. The experiment I hâve proposed, which incorporâtes Micro-Electro-Mechanical Systems (MEMS) to be included as payload of a future Lake Lânder probe, could measure the température, the pressure of the liquid and produce a 3D topographie map of the bottom [START_REF] Bampasidis | Sounding the interior of Titan's lakes by using Micro-Electro-Mechanical Systems (MEMS)[END_REF]. Any température inversions would indicate vertical movement of the liquid. This experiment is described in detail in the following Chapter.

The existence of an internai water-ammonia océan is suggested both by modeling and by Cassini-Huygens mission results. The question is whether there are exchanges between this océan and the surface (and eventually atmosphère) System. Are the lakes such a link? The possibility of life signs in the internai liquid water deposit cannot be excluded.

Models hâve also predicted that during the first tens millions of years after Titan's formation, its global océan was in direct contact with the atmosphère on one side and with the bedrock on the other. This structure may provide conditions favorable for an efficient prebiotic chemistry towards the emergence of life. However, similar conditions are not incompatible with life as we know it on Earth (Fortes, 2000;Raulin, 2008a;[START_REF] Raulin | Planetary science: Organic lakes on Titan[END_REF].

Astrobiological studies on Titan can also provide insights regarding the future of life on Earth. On Titan a methane cycle occurs similar to the water cycle on Earth.

However, the source of methane atmospheric budget replenishment is still unknown. Titan can help us understand the climate change in a planetary atmosphère when the main compound of a global cycle is rapidly lost. As it is mentioned in the Nixon et al. white paper21, which I co-authored, climate change studies on Earth showed an increase in the mean surface température of the planet. If such an increase continues, the évaporation of the water océan will also increase. Water vapor will reach the stratosphère where it will be irreversibly photolyzed, as H2 escapes to space. CO2 levels will rise, and the Earth will heat up and dry out (e.g. Li et al., 2009). Therefore, Titan studies can provide crucial information about the atmospheric equilibrium on the Earth.

6.3

Enceladus, an active small satellite Enceladus is the sixth satellite of Satum in size, quite small compared to Titan, with a mean radius of 252 km (a tenth of Titan's). The first images of the satellite were taken by the Voyager space missions, which illustrated a highly reflective surface with few impact craters implying for a young surface. VI also showed that the satellite is orbiting within the extended Satum's E-ring [START_REF] Smith | Encounter with Satum: Voyager 1 Imaging Science Results[END_REF]. Its mass is 1.08xl020 kg, while its mean density is 1,610 kg/m3 [START_REF] Jacobson | The Gravity Field of the Saturnian System from Satellite Observations and Spacecraft Tracking Data[END_REF]. It orbits Satum at 3.95 Rs (Satumian radii, Rs=60,330 km, e.g. [START_REF] Postberg | The E-ring in the vicinity of Enceladus. II. Probing the moon's interior -The composition of E-ring particles[END_REF] between the moons Mimas and Tethys. Due to the small mass and the weak gravitational field, the satellite has a negligible atmosphère [START_REF] Parkinson | Enceladus: Cassini observations and implications for the search for life[END_REF].

The surface of Enceladus is fully covered by ice, having albedo values of about 1, the highest in the Solar System (Fig. 6.9). Its mean surface température is about 75 K. The surface features are mainly smooth and cratered terrains, rifts, ridges, grooves, escarpments and extensive linear fractures in the South pôle. Cassini/VIMS indicates that the surface consists of CO2 and organics (Brown et al., 2006). The Cassini-Huygens mission has significantly improved our knowledge of Enceladus during 19 close flybys so far. One of the most important discoveries of Cassini was the existence of large plumes ejected from its south polar région as firstly discovered by its Magnetometer (MAG) (Dougherty et al., 2006). This finding is relevant to astrobiology since it implies a complex organic chemistry occurring in the interior of Enceladus along with the appropriate thermal energy, in the presence of liquid water at short distances from the surface.

Enceladus plumes and habitability

By studying Enceladus, an example of active cryovolcanism in icy satellites, we can understand the processes that shape the surfaces of other icy moons. These processes include tidal heating, possible internai convection, cryovolcanism, and ice tectonics. Moreover, the plume source région on Enceladus samples a warm, chemically rich environment that may facilitate complex organic chemistry and biological processes. The geyser-like features on Enceladus are ejected from its south pôle from a sériés of sub-parallel fractures named as "Tiger Stripes" (Fig. 6.10), spewing a sériés of jets more than 1000 km into space, see Figure 6.11 (Porco et al., 2006;Spencer et al, 2006). Their shape is linear while their structure typically about 500 m in depth, 2 km in width and around 130 km in length, flanked on both sides by prominent 100 m high ridges. Cassini/CIRS mapped higher températures from these linear features compared to the surrounding area (Spencer et al., 2006). The mass production rate of the plume gas has been estimated to be at about 150 kg/s from occultation data (Tian et al., 2007). If this value remains from the formation of the satellite, it is sufficient to remove a significant fraction (~20%) of Enceladus' mass (Kargel, 2006).

Figure 6. 11 -Enceladus' plumes ejected from the south polar région by Cassini/ISS (Porco et al., 2006).

The Cassini Cosmic Dust Analyzer (CDA) showed that the plumes' columns consist of small ice grains with low concentration in sait far away from the satellite. Close to the surface, large grains enriched with sodium and potassium dominate the plume's column [START_REF] Postberg | A salt-water réservoir as the source of a compositionally stratified plume on Enceladus[END_REF]. ammonia and other various organic compounds, deuterium and 40Ar, as well as complex organics like benzene and other species such as methanol and formaldéhyde (Waite et al., 2009).

The Cassini Ultraviolet Imaging Spectrograph (UVIS) observed an occultation of the Sun by the plumes of Enceladus on 18 May 2010. Water vapor was confirmed as the dominant component, while the upper limit for N2 was less than 0.5%. Combination of the results of the three Cassini instruments (CDA, INMS and UVIS) supports the existence of an internai liquid réservoir with gas escaping through the narrow channels in the Tiger Stripes région [START_REF] Hansen | The composition and structure of the Enceladus plume[END_REF]. Cassini/CAPS detected positive water group ions [START_REF] Tokar | Cassini détection of Enceladus' cold water-group plume ionosphère[END_REF], water-cluster molecular ions [START_REF] Coates | Négative ions in the Enceladus plume[END_REF] and charged nanometer-sized grains [START_REF] Jones | Fine jet structure of electrically charged grains in Enceladus' plume[END_REF][START_REF] Hill | Charged nanograins in the Enceladus plume[END_REF] in the plumes of Enceladus. Moreover, Cassini /RPWS-LP detected a cold dense plasma in the same material [START_REF] Shafiq | Characteristics of the dustplasma interaction near Enceladus' South Pôle[END_REF].

The Chemical composition of the plume and surface material of Enceladus suggests the presence of a very hot interior, for such a frozen satellite, implying internai températures on the order of 500-800 K. Such hot interior décomposés ammonia into N2 and drives reactions with hydrocarbons (Matson et al., 2007).

Several models hâve been proposed so far regarding the geyser mechanism occurred on Enceladus' surface, which can be divided into two categories. The first assumes that an underground océan forms the plumes [START_REF] Schmidt | Slow dust in Enceladus' plume from condensation and wall collisions in tiger stripe fractures[END_REF]Tobie et al., 2008;Postberg et al., 2009;[START_REF] Behounkova | Tidally-induced melting events as the origin of south-pole activity on Enceladus[END_REF][START_REF] Matson | Enceladus: A hypothesis for bringing both heat and Chemicals to the surface[END_REF], see Figure 6.12, while the second one suggests that the ejected material originates from warmed, melted or crushed ice by tectonic motions (Nimmo et al., 2007), see Figure 6.13. Most of the models suggest the existence of a liquid water environment undemeath the Tiger Stripes.

Since Enceladus is not in hydrostatic equilibrium (Schubert et al., 2007;2010) a simple and very general stratigraphie interior is being suggested which consists of a 169 km rocky core overlain by an icy 82 km mantle [START_REF] Barr | Convection in Enceladus' ice shell: Conditions for initiation[END_REF]Fortes et al., 2007;Schubert et al., 2007).

Conclusions and questions about Enceladus' astrobiological potential

Enceladus, Saturn 's most active moon, is the place in the Solar System System where we hâve observational evidence of a habitable environment already occurring in terms of energy, organics, stable environment and liquid water. Despite its small size, the moon can possess sufficient dynamical energy to drive a plume éjection 1000 km in space out of the moon 's south pôle and eventually feed the outer E-ring of Saturn (Postberg et al., 2009). In a model by Cooper et al. (2009), the astrobiological parameters that support life on Enceladus are evaluated as higher than for Europa due to a less extreme State of oxidation and greater residual abundance of organics. The quest of habitability on Enceladus includes: the presence of liquid water (either in a subsurface océan, in the plume vent régions, or elsewhere); the extension and long-term stability of the water liquid; the interactions of the liquid réservoir and the surface; the degree of thickness and uniformity of the ice crust; the available energy sources and finally the presence of any biomarkers or biological building blocks.

6.4

Jupiter's satellites as astrobiological targets: Ganymede, Europa and (Kivelson et al., 2002), generated by currents flowing in a conducting liquid, which indicates the existence of the deep océan. It is covered by a thin exosphere consisting mainly of oxygen in the forms of O2, O and O3 [START_REF] Noll | Détection of Ozone on Ganymede[END_REF][START_REF] Calvin | Latitudinal Distribution of 0 2 on Ganymede: Observations with the Hubble Space Telescope[END_REF]Hall et al., 1998).

Most of Ganymede's surface coverage displays dark régions which are filled with impact craters [START_REF] Prockter | Dark Terrain on Ganymede: Geological Mapping and Interprétation of Galileo Regio at High Resolution[END_REF] and brighter régions covered by terrains curved by tectonic ridges and grooves (Fig. 6.14) (Pappalardo et al., 1998a;Patterson et al., 2010).

Several non-water ice materials hâve been found in the surface of the satellite such as MgSÛ4 and possibly Na2SC>4 (McCord et al., 1998a;2001). Galileo gravitational measurements showed that Ganymede is fully differentiated into a core and a mantle [START_REF] Anderson | Gravitational constraints on the internai structure of Ganymede[END_REF]. It consists of an iron-rich sulfide core in liquid State, a silicate-rich mantle covered by an icy crust up to 100 km (Fig. 6.15). A liquid océan possibly exists 200 km below the surface, between layers of ice (McCord et al., 2001). Réflectance spectra from Galileo/NIMS discovered hydrated minerais on Ganymede's surface, similar to Europa's (McCord et al., 1998a), derived from the encapsulated internai océan.

Europa as a habitat

Jupiter's Europa is the smallest of the four large Galilean satellites with diameter of 3,122 km, a mass of 48 x 1021 kg and a density of 3,010 kg/m3 (Fig. 6.16). It orbits Jupiter at 671,000 km and its surface has a mean température of about 100 K. As shown by the Galileo mission, the lack of many impact craters on its surface [START_REF] Pappalardo | Geological evidence for solid-state convection in Europa's ice shell[END_REF] indicates that it is relatively young and active,. It is mainly composed of water ice (Dalton, 2010;Dalton et al., 2010) and the most typical surface structures are linear chains, the lineae (e.g. Figueredo & Greeley (2004)). Réflectance spectra from Galileo's Near Infrared Mapping Spectrometer (NIMS) showed that the surface of Europa also contains hydrated sait minerais such as magnésium sulfates and sodium carbonates [START_REF] Mccord | Salts on Europa's Surface Detected by Galileo's Near Infrared Mapping Spectrometer[END_REF].

Hubble Space Telescope observations showed that Europa possesses a tenuous exosphere mainly consisting of molecular oxygen [START_REF] Hall | Détection of an oxygen atmosphère on Jupiter's moon Europa[END_REF] which has not biological origin. However, it may interact with the possible internai liquid océan and thus it may hâve biological significance.

The interior of Europa is differentiated as reported from Galileo gravitational data [START_REF] Anderson | Europa's Differentiated Internai Structure: Inferences from Two Galileo Encounters[END_REF]. It is primarily composed of silicate rock (Sohl et al., 2010) and an iron-rich core (Anderson et al., 1998). Galileo's magnetometer data showed that Europa has an induced magnetic field, which suggests the presence of a subsurface conductive layer (Kivelson et al., 2000). The layer is likely a salty liquid water océan decoupled from the icy crust no more than 100 km thick (Fig. 6.17) (Schenk & McKinnon, 1989;Zimmer et al., 2000;Schenk & Pappalardo, 2004). The big différence compared to Ganymede, the large satellite, is that the liquid layer is almost certainly in direct contact with the silicates.

Rocky Interior \ Liquid Océan Under Ice H2O Layer Jupiter probably generates large tidal waves on Europa, due to its small but non-zero obliquity, which keep the océan warm (Tyler, 2008). The composition of this internai océan is different from the terrestrial ones. The Earth's océan major component is sodium chloride, while Europa's should hâve a highly conductive component such as magnésium sulfate [START_REF] Mccord | Salts on Europa's Surface Detected by Galileo's Near Infrared Mapping Spectrometer[END_REF]Fanale et al., 2001), or sulfuric acid hydrate (Carlson et al., 2005).

Recent analyses of Galileo images suggest that ice-water dynamics are active. Surface features may be formed from exchanges between the ice shell and shallow subsurface water deposits. This mechanism provides transfer of energy and nutrients between the surface and the interior, increasing the associated habitability potential [START_REF] Schmidt | Active formation of'chaos terrain' over shallow subsurface water on Europa[END_REF].

6.4.3

Callisto, the heavily-cratered moon Callisto is the third largest moon in the Solar System after Ganymede and Titan with a diameter of 4821 km, while its mass is at 107.6 x 1021 kg and its density at 1830 kg/m3. Its surface has a mean température of about 134 K (Fig. 6.18). Of the four Galilean moons, Callisto orbits farthest from the giant planet at 1883 x 103 km [START_REF] Anderson | Shape, Mean Radius, Gravity Field, and Interior Structure of Callisto[END_REF]. Due to its orbital distance, Callisto is less affected by the magnetosphere of Jupiter compared to the other Galilean satellites (Cooper et al., 2001) and it has a tenuous carbon dioxide exosphere [START_REF] Carlson | A Tenuous Carbon Dioxide Atmosphère on Jupiter's Moon Callisto[END_REF].

The surface of Callisto is heavily cratered which means that no surface processes occur on it [START_REF] Greeley | Galileo views of the geology of Callisto[END_REF]. Callisto is very similar to Ganymede, but it is only partially differentiated and does not hâve an intrinsic field [START_REF] Spohn | Océans in the icy Galilean satellites of Jupiter?[END_REF]. The thickness of Callisto's icy lithosphère is between 80 and 150 km (6.19). It possesses a deep water-ammonia océan similar to Ganymede with a thickness between 70-90 km beneath the crust indicated by Galileo magnetometer and gravity data (Khurana et al., 1998;Zimmer et al., 2000;[START_REF] Spohn | Océans in the icy Galilean satellites of Jupiter?[END_REF][START_REF] Kuskov | Internai structure of Europa and Callisto[END_REF]. Beneath the océan the interior is composed by rock and ice and the rock percentage increasing by the depth without having a core [START_REF] Nagel | A model for the interior structure, évolution, and différentiation of Callisto[END_REF]).

Habitability of Ganymede, Europa and Callisto

Surface properties of the satellites provide indications about their interior structure. Callisto has a dark, geologically old, heavily cratered ice surface and from its rotational data that its density is equally distributed [START_REF] Anderson | Shape, Mean Radius, Gravity Field, and Interior Structure of Callisto[END_REF]. The rotation measurements of the three inner Galilean moons (Ganymede, Europa and Io) indicate differentiated interiors. Moreover, Ganymede's surface features indicate possible past tectonic processes. Europa probably hosted more recent tectonic movements, with a thinner ice crust. The nearer a moon is to

Jupiter the hotter its interior. Callisto's internai structure may be the original structure of ail the moons.

The gravitational pull of Jupiter induces stresses to its moons and as a resuit their interiors are heated. Except for Callisto, the tidal heating melts the interior ice, allowing for rock and iron to sink downwards and for water to lift to the surface. Ganymede has a thick ice crust. Europa expériences stronger tidal heating and has a thinner crust. The geological and structural data of Ganymede and Europa provide promising worlds for astrobiological studies.

Gravitational forces from Jupiter and other Galilean satellites cause tidal friction on Europa and provide the proper energy for the internai océan to remain in liquid State. Large amounts of energy at the bottom the océan could form hydrothermal vents like the ones seen on Earth. Such vents can provide the proper température status and ingrédients to support life forms. [START_REF] Greenberg | Transport Rates of Radiolytic Substances into Europa's Océan: Implications for the Potential Origin and Maintenance of Life[END_REF] suggested complex aérobic organisms could exist on Europa, providing evidence of oxygen concentration within the océan greater than that of the Earth's.

Since the salts' concentration is large, only extremophile organisms (like halophiles) could survive (Cooper et al., 2001;Marion et al., 2003).

Ganymede is colder than Europa, which lowers possibly its habitability potential.

However, Ganymede possesses an intrinsic magnetic field which protects the satellite from incident radiation and thus, provide the necessary tools to concentrate biological building block ingrédients (Trinks et al., 2005). Its dark terrain is of astrobiological interest due to the existence of organic materials (McCord et al., 1998a). Additionally, nutrient-rich material provided from magmatic events could exist in the internai océan (Barr et al., 2001).

Europa does not hâve an intrinsic magnetic field, but its icy crust provides some natural shielding to the sub-surface liquid océan, protecting it from direct exposure to the intense radiation environment of Jupiter's magnetosphere. Galilean moons due to the presence of the thick atmosphère. Erosional processes act and shape the surface globally, resulting a fiat topography [START_REF] Lorenz | Hypsometry of Titan[END_REF].

The surface of Titan and the other outer planets ' satellites appears initially as an unfavorable place for life, at least for terrestrial-type life. The absence of liquid water on the surface makes them unlikely to support any terrestrial-type life. If liquid water exists under the surface, it is not in direct contact with a silicate core (with the exception of Europa), which is isolated from the subsurface océan by a layer of a high-pressure ice phase.

However, exotic types of life could be found in the internai large réservoirs of the outer planets. Ganymede, Callisto, Europa, Titan and Enceladus hâve encapsulated water océans, which may maintain an exotic type of life, using liquid hydrocarbons as solvents (McKay & Smith, 2005). According to Stoker et al. (1990), terrestrial bacteria can satisfy their energy and carbon needs by consuming tholins, the laboratory analogue of Titan 's aérosols.

On the other hand, living organisms hâve been found on Earth under extreme environmental conditions in deep oceanic layers. Such environments, with low-temperature and high-pressure conditions, are the cold seeps and the hydrothermal vents (e.g. Ritt et ai, 2010). Single-cell Archaea and Eubacteria microbes hâve been discovered to consume methane and hydrogen sulfide from seep and hydrothermal vents. Additional laboratory experiments and in situ studies of deep subglacial isolated lakes in Antarctica (Kapitsa et al., 1996) would improve our understanding in this field, as the physical properties of deep subglacial lakes resemble those found on outer planets' moons (Bulat et al., 2009).

The satellites of the giant planets like Ganymede, Callisto, Europa, Titan or Enceladus are possible habitable environments and astrobiological targets. According to current models of internai structure, the existence of subsurface océans is expected for most of the icy moons of the outer planets (e.g. Sohl et al., 2010, Schubert et al., 2010 andreferences therein). The common properties that need to be satisfied on ail bodies in order to sustain a liquid subsurface océan are [START_REF] Solomonidou | Water Océans of Europa and Other Moons: Implications For Life in Other Solar Systems[END_REF] The stability of the crust against convection (e.g. [START_REF] Mckinnon | Geodynamics oflcy Satellites[END_REF][START_REF] Rainey | Océans in ice-rock bodies: conditions for the existence of subsurface liquid water[END_REF].

The internai océan could provide the proper isolation to the satellites of the outer planets to form possible habitable environments. The discovery of hydrocarbon lakes on Titan's surface and the possible existence of subsurface liquid océans in Ganymede, Callisto, Europa, Titan and Enceladus suggest reconsidering the habitable zone définition and limits (Lammer et al., 2009).

However, life forms that do not influence the atmosphère of their host planet on a global seule will not necessarily be remotely détectable. In the Solar System's neighborhood, such potential habitats can only be investigated with space missions (see Chapter 7). For the Jovian and Saturnian satellites the proposed future missions will address the question of the hypothesized internai liquid water océan for their icy moons (see following Chapter).

Although it is a difficult challenge to search for life forms in the environments of outer planets' moons, such studies will profoundly improve our understanding of the origin and the évolution of life and the habitable environment on our own planet.

Future exploration of the satellites of the outer planets -MEMS devices for icy moons

The gas giants are massive planets, mainly composed of hydrogen and hélium, but with no well-defmed planetary surface. Our Solar System has four gas giants: Jupiter, Satum, Uranus and Neptune. They are also known as Outer Planets22 since they orbit beyond the asteroid belt. Recently, many extrasolar gas giants hâve been detected, for which our Solar System giant planets serve as archétypes.

Jupiter A new NASA mission is on its way to the Jovian System: The Juno mission23, which was launched in August 2011 and will enter into orbit insertion at 2016. The eight instruments of Juno will study Jupiter's gravitational and magnetic potential as well as its composition.

However, this mission lacks a lander vehicle for the in situ exploration of an icy moon. payloads were designed to study the moons' surfaces, interiors, their gravitational fields and magnetospheres and their atmospheric environments. ATHENA, a big X-Ray telescope and NGO, a concept of three satellites to detect gravity waves. The complété instrument payloads, as well as, the instrument contractors hâve not yet been decided. A NASA contribution towards the payload is under considération. The participation of NASA will be limited to a few instruments perhaps and to co-investigators.

7.1.2

The Jupiter Icy Moons Explorer mission (JUICE)

Science goals

The Jupiter Icy Moons Explorer mission (JUICE) will visit the Jupiter System focusing on the investigation of Ganymede mainly, but also Callisto and to a lesser extent of Europa as potential habitats and on the exploration of Jupiter. JUICE will study the conditions that may hâve led to the emergence of habitable environments among these icy satellites and their internai océans. The mission will also focus on characterizing the diversity of processes in the Jupiter System, including gravitational coupling between the Galilean satellites and their long term tidal influence on the System as a whole (ESA/SPC(2012)12)24.

Mission profile

The mission25 will be launched in June 2022 on an Ariane 5 rocket carrier and will perform a 7.3-year cruise towards Jupiter, based on an Earth-Venus-Earth-Earth gravitational assist. The Jupiter orbit insertion will occur in January 2030, and will be followed by a tour around the Jupiter System. It will comprise a transfer to Callisto (11 months), a phase studying of Europa (with 2 flybys) and Callisto (with 3 flybys) lasting one month. Then, a "Jupiter high-latitude phase" will begin that includes 9 Callisto flybys (lasting 9 months) and the transfer to Ganymede which will last 11 months. In September 2032 the spacecraft will insert into orbit around Ganymede, starting with elliptical and high altitude circular orbits (for 5 months) followed by a phase at a medium altitude (500 km) circular orbit (of 3 months) and by a final phase at a low altitude (200 km) circular orbit (for 1 month). The end of the nominal mission is planned to be in June 2033.

Spacecraft description

The spacecraft is three-axis stabilized and powered by solar panels, providing around 650 W at end of mission. Communication to Earth is provided by a fixed 3.2 m in diameter high-gain antenna, in X and Ka bands, with a downlink capacity of at least 1.4 Gbit/day. ). ESA plans to set up web sites focusing on the JUICE mission, which will be upgraded as the mission progresses and I will contribute towards this effort. Ail of this I will further develop in the next chapter. As regards the future missions, my contributions were in particular centered around providing an atmosphere-surface interaction combined approach for the possible instrumentation aboard the orbiters (through my expérience with

Cassini/CIRS and Cassini/VIMS, see Chapters 4 and 5) and the HASI instrument aboard

Huygens, but also through novel ideas that I will describe more in detail hereafter on future in situ instrumentation in case of a balloon or a lake lander, in particular.

7.2

Mission studies to Saturn and Titan: KRONOS, TSSM and TAE Satum is the second largest planet of the Solar System after Jupiter. Its radius is 60,330km, while its mass is 568.46x1024 kg. It orbits around the Sun at about 10 AU and it takes almost 29.5 years for a full révolution. It has an intrinsic magnetic field at 0.2 gauss at the equator, which is perfectly symmetric. This field créâtes the second largest magnetosphere in the Solar System after Jupiter [START_REF] Blanc | Solar System Magnetospheres[END_REF][START_REF] Belenkaya | Définition of Saturn's magnetospheric model parameters for the Pioneer 11 flyby[END_REF][START_REF] Gombosi | Satum's Magnetospheric Configuration, Saturn from Cassini-Huygens[END_REF]. As the Cassini magnetometer discovered, the main source of neutrals in the Satumian magnetosphere is Enceladus (Dougherty et al., 2006). Its magnétopause extends up to 20 Satumian radii at the orbit of Titan.

The Satumian atmosphère is dominated by hydrogen at about 96%, while the rest 3% the Kronos mission (Marty et al., 2009a;b), and to Titan and Enceladus, the TandEm mission (Coustenis et al., 2009a). The TandEM concept was later merged with a NASA Flagship mission (Titan Explorer, [START_REF] Leary | Titan explorer flagship mission study[END_REF] to form the TSSM concept.

I contributed to both proposais and was therefore included in associated publications or communications. Hereafter, I describe first the mission concept for the study of Satum with probes and then the TSSM mission, a very complété and challenging mission concept and compare it to other smaller mission proposais for a balloon or a lake lander.

The KRONOS mission

The Kronos mission focused on studying Satum as a whole System, from its magnetosphere and its atmospheric environment to its interior. The basic concept of this mission was to probe globally the giant planet with two probes and remote sensing techniques through a carrier which would also serve as relay and dérivé the Chemical and isotopic composition of the giant The mission would measure the Satumian gravity with high accuracy in order to yield constraints about its interior. Gravitational data would also détermine if the interior rotâtes as a solid body or not. The knowledge of the composition and the conductivity of the interior would help understanding better the asymmetric magnetosphere of the planet. Kronos would also observe directly the Satumian rings.

Data from the Kronos mission compared with the results from the Juno mission could hâve provided the opportunity to leam about the dynamics and the évolution of our Solar System's giants. Additionally, since advanced space télescopes like KEPLER31 are now able to observe extrasolar planets, parallel studies could also give detailed information about the physics, chemistry and thermodynamic State of other solar Systems.

The mission architecture

The Kronos mission proposai consisted of a spacecraft, which carried two atmospheric probes launched into the atmosphère of the planet. Each probe was planned to make in situ température measurements down to the pressure limit of 10 bars. The spacecraft design is similar to the one from the Juno mission. Two other ring probes hâve also been proposed to be included in order to study the rings and their particles within a distance of a few kilometers, similar to that of the Galileo entry probe. Structure Instrument (ASI) on board the entry probes design would beneflt the Huygens Atmospheric Structure Instrument (HASI) héritage [START_REF] Fulchignoni | The Characterisation of Titan's Atmospheric Physical Properties by the Huygens Atmospheric Structure Instrument (Hasi)[END_REF].

I hâve investigated some of the possible modifications on the HASI instument to adapt to the higher pressures of the deeper Satumian atmosphère and focused on the température sensors which would perform direct measurements during the descent phase after the release of the probe's shield. Each HASI température sensor was a dual element platinum résistance thermometer [START_REF] Fulchignoni | The Characterisation of Titan's Atmospheric Physical Properties by the Huygens Atmospheric Structure Instrument (Hasi)[END_REF] designed to meet the Cassini/Huygens mission requirements.

However, the environmental conditions in Satum differ significantly from Titan's. Prof. X. Moussas and I were the co-leaders of the TandEM Outreach working Group.

We hâve organized several outreach activities and participated or attended in many outreach events globally. A detailed list and a description of our outreach projects and their évaluation is in Chapter 8. In collaboration with A. Coustenis and X. Moussas I will continue to support the outreach campaign of the Titan future mission.

TandEM mission concept

The mission concept of TandEM (Fig. 7.6) consisted of two moderate-size spacecraft the Titan-Enceladus Orbiter with the Enceladus penetrators and the Carrier for the Titan in situ investigation éléments which were the Titan Montgolfier-hot air balloon and the three lander/mini-probes (Coustenis et al., 2009a).

After its arrivai, the orbiter would perform several flybys of Titan and Enceladus and carry the Enceladus landers/penetrators. The carrier spacecraft wwoud release the Titan hot air balloon and up to three entry probes. To facilitate the télécommunication with the in situ vehicles, the carrier could be converted into an orbiter and perform science measurements like occultations using the telecom link between the two orbiting spacecraft. The telecom relay between the Titan in situ éléments and the Earth orbiter would be provided by the orbiter. According to the initial mission scénario, the Titan-Enceladus orbiter would perform 12 Enceladus flybys as well as 2 Titan flybys and 3 or 4 more Titan flybys before entering into Titan Orbit Insertion (TOI). Then, the orbiter would enter in Titan polar orbit, in order to support the in situ mission éléments. For enchanced performance, the spacecraft would be maneuvered to a polar orbit before the delivery of the balloon and the entry probes by the carrier.

The Montgolfière consisted of a balloon and a gondola. The balloon would be filled by Titan's gas, heated by the energy from the two Multi-Mission Radioisotope Thermoelectric Generators (MMRTGs) which were located in the gondola with the scientific equipment.

Titan's dense nitrogen atmosphère with températures less than 100 K at the first 70 km above the surface is favorable for a hot air balloon flight. The average operation altitude of the Montgolfière is 10 km above the surface.

The Titan mini-probes design would follow mainly the Huygens probe expérience as well as the GEP-ExoMars (Geophysics Package). The parachute System would hâve been similar to the one used for the Huygens Entry Descent and Landing (EDL) sequence. These probes would be powered by Radioisotope Thermoelectric Generators (RTGs). Figure 7.7

shows the design of a pénétration module.

Détachable Propulsion Stage

Figure 7. 7 -TandEM pénétration probe (Coustenis et al., 2009a).

These penetrators could be released by the Montgolfière or the carrier. They could include seismometers as payload and thus construct a seismic network on the surface of Titan.

Two penetrators would be delivered to Enceladus surface from the Titan-Enceladus orbiter.

TandEM, along with the LAPLACE mission to Jupiter [START_REF] Blanc | LAPLACE: A mission to Europa and the Jupiter System for ESA's Cosmic Vision Programme[END_REF] proposais An international team of scientists and engineers with 16 US and 15 European members, the Joint Science Définition Team (JSDT), was then formed in order to develop a mission to Titan that would address the key questions left by the Cassini-Huygens mission.

The JSDT merged the NASA Titan Explorer study for a mission to Titan [START_REF] Leary | Titan explorer flagship mission study[END_REF] 33 and the TandEM study into a single one named "Titan and Satum System Mission" (TSSM) .

The TSSM JSDT also examined the Decadal Survey of 2003 documents of the US National Academy of Sciences, which regarded Titan and Europa as the highest priority targets for future missions. TSSM would be TiRS and HiRIS (see description heareafter, in Table 7.2).

My other interest in future Titan exploration, more focused, is the lakes study. Cassini has identified their existence, but crucial information about their nature, structure, origin and évolution is not to be addressed by the current mission. Cassini monitors remotely the lakes and only in situ exploration could calculate their liquid volume, identify their bottom topography and their contribution to the active methane cycle. The TSSM concept also includes a lake lander, which could provide the relevant information. For this purpose, I hâve proposed the emplacement of micro-probes in the lake lander payload. The main idea is to use Micro-Electro-Mechanical-Systems (MEMS) technology is to establish a network of micro probes spread in the interior of a lake. These devices are lightweight, low cost and highly reliable. The concept and the constraints of MEMS experiment are described in detail in the following section.

The TSSM concept

The current TSSM concept foresees a carrier to be launched on an Atlas V 551 who would reach Satum by using an ion propulsion System, the Solar Electric Propulsion (SEP) to gain thrust from ion beams as well as gravity assist trajectory, capable of accomplishing outer planet missions. TSSM basic concept consists of three space vehicles (Fig. 7.8):

a) The orbiter, which will tour mainly around Titan for two years, after performing seven close-up Enceladus' flybys. On Enceladus, the orbiter will sample the plumes and make subsurface measurements.

b) The hot air balloon (Montgolfière) will probe both Titan's atmosphère and surface at a latitudinal bin of circa 20°around the equator, being in low altitude orbit of 10 km for at least six months.

c) The Lake Lânder will perform the first extraterrestrial océanographie experiment by landing and floating on the largest of Titan's lakes, the Kraken Mare.(however the lake to be explored could be a different target depending on the mission arrivai time).

d) The TSSM orbiter after its arrivai to Satum will perform a 24-month System tour including sixteen flybys of Titan and seven close flybys of Enceladus. Four of these Enceladus flybys will be over the pôles at 100 km, one at 300 km and two at 1000 km.

During this tour, the orbiter will encounter several icy moons and it will observe Satum. It will release the hot air balloon during its first flyby of Titan, while the lake lander will be released at the second flyby.

e) After this phase, the orbiter will enter into highly elliptical orbit to conduct a two month concurrent aerobraking and aerosampling phase in Titan's atmosphère, reaching the altitude of 600 km. Then it will execute a final periapsis raise bum to achieve a 1500 km circular, 85°mapping orbit. The duration of this phase will be 20 months.

f) Finally, the orbiter will perform a maneuver, which will start the orbit decay assisted by the influence of Satum's perturbations and Titan's atmospheric drag. In this stage it will observe the magnetospheric interaction between Satum and Titan. The TSSM orbiter is a three-axis stabilized spacecraft. The basic éléments of the spacecraft as well as the positions of the Montgolfère and lander on the mothership are illustrated in Figure 7.9, while its payload is listed in Table 7.2. The orbiter is also the carrier of the two other vehicles of the mission. The Montgolfière is a hot air balloon to be released by the orbiter during the first flyby. It will be powered with 100 W by. Figure 7.10 shows the entry System of the balloon capsule, while Figure 7.11 illustrâtes the position of the instrumentation in the gondola. From the altitude of 40 km the balloon will start to be filled by air and 1.7 W from the a Multi-Mission Radioisotope Thermoelectrical Generator (MMRTG) will warm the air for buoyancy at about 10 km above the surface. The Montgolfère will then navigate with the winds at 1-2 m/s. The balloon's gondola will be equipped with scientifïc instrumentation in order to perform atmospheric measurements, imaging, spectrometry, subsurface radar profiling, and electric and magnetic field measurements (Table 7. The Lake Lânder will follow an entry, descent and landing (EDL) procedure like the Huygens EDL, while the orbiter will be in continuous contact with it. The target location will be the Northern largest lake of Titan, Kraken Mare. Figure 7.12 illustrâtes the aeroshell of the Lake lander, while Figure 7.13 shows its payload. The probe will be powered by batteries like Huygens and its nominal mission is designed to last for 9 hours, 5 hours in the descent phase and 3-4 hours after the touchdown on the lake's surface. The floating configuration of the lake Lander is depicted in Figure 7.14 and its instrumentation is listed in Table 7. The TSSM mission is planned to be supported by ground based radio télescopes [START_REF] Coustenis | Earth-Based Support for the Titan Satum System Mission[END_REF], following the héritage of the Huygens data retrieval in 2005 [START_REF] Witasse | OverView of the coordinated ground-based observations of Titan during the Huygens mission[END_REF]. Titan Lake Probe (Fig. 7.18) is a lake lander, similar to what we had proposed in TSSM, which could be considered as part of a larger mission or as a stand-alone mission. The main objective of this proposai is to investigate the lake deposit and the physical properties of the liquids like the TiME concept (Waite et al., 2010). [START_REF] Hall | Titan Aerial Explorer (TAE): Exploring Titan By Balloon[END_REF]. TAE was a helium-filled high-pressure balloon, which was planned to fly in the lower atmosphère of Titan at an altitude of 8 km from 3 to 6 months on Titan's équatorial latitudes. The data transmission of about 1 Gbit (total) to Earth would be established directly, avoiding the intervention of an orbiter. As a follow-up to the TSSM mission with similar concept and objectives, my interest in TAE was obvious, as I will describe hereafter.

I hâve also created the web page of the mission proposai using flexible web tools in order to achieve a quick impact to the scientific community and the public, which is: http://users.sch.gr/gbabasides/joomla/index.php

The TAE concept

The Titan Aerial Explorer (TAE) mission is a super-pressure hélium balloon, 4.6 m in diameter and spherical in shape. It combines the mobility and coverage of an orbiter with the capability for high resolution and in situ observations demonstrated by the Huygens lander.

The balloon would be powered by two Advanced Stirling Radioisotop Generators (ASRGs) (Figure 7.19) but without propulsion [START_REF] Hall | Titan Aerial Explorer (TAE): Exploring Titan By Balloon[END_REF]. Although TAE aérostat and TSSM Montgolfière air platform hâve similar objectives in Titan science, they présent différences in architecture. In Table 7.5 below, I summarize the results of the comparison between TSSM Montgolfière and TAE. TSSM balloon is twice in diameter compared to the TAE one and it will be filled by Titan gas, which means that there is no need to carry spécial tanks, and therefore it has more space for scientific instrumentation. Instead, TAE balloon is a high-pressure hélium filled balloon and has to carry the gas from the Earth. This fact influences the amount of data which will be collected and transmitted to Earth from both vehicles. The Montgolfière gondola will host 8 scientific instruments which will provide data volume three orders of magnitude larger than TAE.

Both vehicles will be powered by Radioisotop generators: TAE by ASRG and TSSM Montgolfière by MMRTG. ASRGs are préférable power sources compared to MMRTGs, since they combine reduced mass, increased power output and reduced cost. Nevertheless, the Montgolfière architecture can easily adopt ASRG.

One more différence between the proposed balloon concepts is that the Montgolfière will use the orbiter in order to transmit data to Earth, while TAE will establish a direct-to-Earth link. By comparing the two, it is obvious that both the Montgolfière and TAE can provide satisfactory proposais for future Titan exploration. However, the the former could be considered as a more complété option since it carries more instruments onboard including a magnetometer and its imaging spectrometer provides a wider range of wavenumber.

Furthermore, if the Montgolfière concept manages to adopt ASRGs instead of MMRTG, it will enhance its energy budget. Finally, the scientific value of the TSSM hot air balloon can be increased by the possible emplacement of a simple seismometer in its beat shield. TAE made the first sélection from 47 to 14 mission concepts, and was judged excellent in science, but finally it was not selected. The major concem was about the implémentation and especially the availability of ASRGs, provided from US.

7.3

Micro-Electro-Mechanical-Systems (MEMS) for Titan Lakes

In order to overcome portability and consumption restrictions in any future outer planetary missions' requirements, new alternative technologies should be considered. Towards this direction, the adaption of Micro-Electro-Mechanical-Systems (MEMS) devices has been proposed, as part of the science surface properties package of a future probe to Titan like the TSSM Lake Lânder [START_REF] Bampasidis | MEMS Techniques for Sounding the Interior of Titanic Lakes[END_REF]201 lb).

In this section, I describe how I propose MEMS devices to operate as microlaboratories by including RF wavelength emitters and température and pressure sensors.

These micro-machines could obtain the 3D sounding of the liquid deposit, its Chemical composition and detect the presence of any biomarkers in a broader area. The température and pressure micro-sensors could provide the vertical pressure, température and density profile of the liquid deposit [START_REF] Bampasidis | Sounding the interior of Titan's lakes by using Micro-Electro-Mechanical Systems (MEMS)[END_REF]. Obviously, the utilization of any innovative instrumentation requires complété knowledge of the environmental issues and parameters and, of course, deep understanding of their complexity.

The MEMS pattern, due to its very small shape and size without reducing their operational performance, is the idéal payload for a future Lake Lânder probe in Titan and outer planetary space missions in general. As this unique technology is new in the space field, further analysis and testing nécessitâtes. MEMS implémentation in Titan's exotic environment is a great challenge for science, engineering and space physics.

As a MEMS device can be considered any small-size product within the range of a micron to a centimètre which also combines mechanical, as well as, electrical structures.

Although H. C. Nathanson in late 1960s constructed the first MEMS device [START_REF] Nathanson | The Résonant Gâte Transistor[END_REF], only after the 1980s minor integrated circuits were massively produced in microscopie scale (Stark & Bernstein, 1999).

Due to their small size and shape (20 pm to 1 mm), MEMS hâve unique properties, which increase the performance of every scientific experiment for a wide range of uses and applications. MEMS can provide an excellent rate of the optimum shape to the rendering performance relation. MEMS devices are able to easily locate impossibly approachable places by traditional instrumentation. They can successfully substitute quite larger in size and shape conventional devices, giving the user the opportunity to explore the structure of the nature in minor scales. Therefore, because of the minimization of the cost in their manufacture and operation stages, MEMS instruments are préférable for scientific applications, improving the quality and widening the range of the obtained data.

MEMS devices in the space scene

The advantage of implementing MEMS techniques in Aerospace and Space Systems become obvious when the reduced requirements in size, mass, power issues of many Aerospace/Space applications are taken into account [START_REF] Janson | MEMS, Microengineering and Aerospace Systems[END_REF]. Minimizing the electronic instrumentation assembly is a promising approach when it employs MEMS. The MEMS revolutionizing technology includes ail the properties of the larger devices, which are replaced in an inexpensive way. Therefore, the presence of MEMS in space Systems offers important fmancial advantages to space agencies by severely shrinking the cost of any mission. Indeed, MEMS are reducing the weight of future spacecraft without diminishing the scientific performance of the instrument.

Among many other applications, MEMS devices hâve already been used in the adaptive optics imaging technology in some major télescopes (Krishnamoorthy & Bifano, 1995;[START_REF] Roggemann | Micro-electro-mechanical déformable mirrors for aberration control in optical Systems[END_REF][START_REF] Dayton | Démonstration of new technology MEMS and liquid crystal adaptive optics on bright astronomical objects and satellites[END_REF][START_REF] Zhou | Characterization of contour shapes achievable with a MEMS déformable mirror[END_REF]. Moreover, the Jet Propulsion Laboratory (JPL) has already developed a new MEMS déformable mirror (DM)

System for NASA's adaptive space-based télescopes and in particular for Terrestrial Planet Finding (TPF) mission [START_REF] Stewart | Design and development of a 329-segment tip-tilt piston mirrorarray for space-based adaptive optics[END_REF].

Likewise, several proposais hâve recently been reported for future space missions and télescopes, exploiting the new leading edge MEMS technology. Especially, MEMS techniques hâve been incorporated as detectors, spectrophotometer and an infrared caméra of the new proposed telescope called MTEL (MEMS Telescope for Extrême Lightning) in order to observe the extreme lightning occurring in the upper atmosphère [START_REF] Nam | A telescope for observation from space of extrême lightnings in the upper atmosphère[END_REF].

MEMS devices seem also appropriate for observing fast moving objects and transient events by the proposed space-based telescope called Obscura [START_REF] Park | Obscura telescope with a MEMS micromirror array for space observation of transient luminous phenomena or fast-moving objects[END_REF]. MEMS hâve also been implemented in Space Infrared telescope for Cosmology and Astrophysics (SPICA), the Japanese coronagraph which is planned to be launched in 2017 [START_REF] Swinyard | The space infrared telescope for cosmology and astrophysics: SPICA A joint mission between JAXA and ESA[END_REF].

Although the use of MEMS in space applications is in fairly early stage, they will certainly optimize the scientific value of any mission.

Scientific Goals and Objectives

I proposed the implémentation of MEMS technology as part of the payload of a future lake lander in Titan. I hâve presented this proposai at several meetings and hâve had the opportunity to discuss it with international experts like Dr. Hunter Waite who was interested for a submarine in a Titan lake and Dr. Jean-Pierre Lebreton Huygens Project Scientist/Mission manager.

A MEMS experiment would expand the scientific potential of a future in situ space mission to Titan, as these devices can provide miscellaneous measurements, ail included in a single capsule. These micro-machines will be able to:

1. Construct a topographie map of the lake floor 2. Probe the température and the pressure of the liquid layers 3. Contribute to comparative Planetology

Investigate the presence of any biomarkers in a broader area

My study includes several micro probes in the concept of the mission, in order to establish a broad network in the lake, when these devices are released from the lander during its descent. The MEMS network can provide a three dimensional sounding of the lake's bottom topography. This unique map will be constructed by the RE signais transmitted from each device and received by the sensor on-board the Lake Lander. The floor of the lake hosts dues about the processes, which created and shaped the lake and any signs of exchange material with the interior.

As far as the structure of the lakes on Titan is concemed, the contribution of this experiment will be important as well. Firstly, the depth of the lake will be identified, enabling researchers to test their models. Nowadays, only assumptions hâve been used to détermine the lake depth either by studying the surrounding topography of the lake in accordance with terrestrial models, or correlating the lake size to the depth of the Earth's 20 largest lakes (Hayes et al., 2008;Lorenz et al., 2008b;[START_REF] Paillou | Microwave dielectric constant of liquid hydrocarbons: Application to the depth estimation of Titan's lakes[END_REF][START_REF] Griffith | Possible tropical lakes on Titan from observations of dark terrain[END_REF]. When the volume of the liquid inventory is calculated, it can be considered in the schéma of the methane cycle occurring on Titan (Atreya et al., 2006). The existence of hydrothermal vents on the lake floor could also provide information on the methane cycle of Titan and its hidden source, which replenishes the atmospheric abundance.

MEMS data can also provide information about the interaction of Titan's liquefiers with the lower atmosphère. Particularly, a methane dominated lake affects seriously the local atmospheric status due to the cooling effect of the methane évaporation. This cooling, causes a réduction of the lake surface température and as a resuit an atmospheric downward heat current balances the whole pattern and saturation with methane of the local air mass (Mitri et al., 2007). On the other hand, a lake abundant in ethane is independent of evaporative cooling and it could contribute to the atmospheric convection by remaining in liquid State during the winter like their large terrestrial analogues (Brown et al., 2009).

Measuring the local température and pressure profile with the MEMS sensors, scientists will be able to evaluate such current models and détermine the local atmospheric heat flux. The incognizance of both composition and structure of Titan's lakes increase the significance of such an in situ experiment. The température and pressure sounding of the liquid deposit with MEMS experiment will provide the vertical distribution of thermodynamic variables (pressure, température, density) beneath the lake surface. This set will supplément HASI vertical distribution profiles (Fulchignoni et al., 2005;[START_REF] Harri | Vertical pressure profile of Titanobservations of the PPI/HASI instrument[END_REF] and improve the study of Titan as a System.

MEMS experiment will be the first extraterrestrial océanographie experiment. New opportunities will be opened for comparative Planetology between Earth and Titan especially looking in terrestrial seas and lakes and their équivalents on Titan.

One of the perspectives of the MEMS experiment would be to investigate the existence of biomarkers located inside Titan's lakes. Molecular biological markers, or biomarkers, are organic structures, sedimentary grown, which contain important information about the precursors of origin of life [START_REF] Peters | The Biomarker Guide[END_REF]. Due to the geologically fresh surface of Titan, with few impact craters [START_REF] Wood | Impact craters on Titan[END_REF], most of the dues of the moon's past hâve been erased. However, in isolated environments, like a deep lake, past geological traces can be detected, remaining there under the protection of the liquid.

Similar organic products, like the terrestrial Stromatolites, which emerged one or two billion years ago in Earth [START_REF] Grotzinger | Stromatolites in precambrian carbonates: Evolutionary Mileposts or Environmental Dipsticks?[END_REF], might be well-preserved at the bottom of a deep lake in Titan, encapsulating evidence from past geological environments (Fig. 7.21).

These patterns could be easily identified during the topographie sounding of the floor of the hydrocarbon lake with the MEMS experiment. Therefore, this innovative micro-system could become a significant tool to astrobiological research as well. Except for Titan, the proposed instrumentation can be incorporated in the payload of any future space mission to other planetary bodies which consist of liquid layers like Europa and even to Enceladus. To sum up, as Titan resembles Earth more than any other planetary body in the Solar System, any results from Titan can be used for studying the terrestrial environments and help scientists to test and evaluate their models for Earth. The microcapsules will be released from the lake lander during the last stage of the Entry, Descent and Landing (EDL) procedure as well as after the touchdown. Each capsule will include a MEMS Unit (MEMS-U). The MEMS-U consists of one température and one pressure sensor for direct measurements and the RF signal transmitter. A similar design is depicted in Figure 7.22. The RF receiver unit would be on board the lake lander probe (Bampasidis et al.,201 lb). The fimctional specificity of the température and pressure sensors is that they should interact with the environment being outside of the MEMS capsule. The capsule will encounter harsh and unknown environmental conditions which may affect the performance of the experiment. In order to implement such a payload, several issues should be addressed which are mainly the absorption of the RF signais by the medium and the operation in low températures range. Therefore, the physical properties of the liquid should be discussed.

The question is if the transmitted RF will be attenuated from the medium before reaching the receiver. If the liquid material behaves as a partial conductor (when consisting of ionic contaminants for instance), it will prevent RF propagation in long distances. Such liquid conductors set the independent operation of these RF emitters inside the liquid under discussion. Media with high permittivity and permeability cause a strong atténuation of the RF signal power. One scénario to avoid the atténuation is to connect the devices by a wire to the mother ship. This pattern can also support the microprobes with energy originated by the carrier.

The only information about the composition of the Titan lakes is provided by the remote measurements from Cassini. However, the exact blend of the lakes is still unknown.

As I hâve described in Chapter 5, the main constituent is expected to be ethane, but more hydrocarbons could be mixed in lesser quantities as well. Modeling results about the composition of Titan's lakes show that the major lake liquid constituents are hydrocarbons, mainly ethane, methane and nitrogen (Table 7.6) [START_REF] Cordier | An Estimate of the Chemical Composition ofTitan's Lakes[END_REF]2012). Hence, we should consider the conductivity of such solvent, which is crucial for the proposed experiment. Nitrogen is a non-electrical conductive and it is used as a réfrigérant in superconductive applications. Alkanes do not conduct electricity and do not get polarized by an electric field, since they are covalent compounds and do not contain freely moving ions.

Aliphatic hydrocarbons hâve non-polar liquids and weak hydrogen bonding liquids. These liquids are non-polar which means that they allow for the propagation of RF waves throughout the medium.

An electrical conductive mixture can be produced if the hydrogen and nitrogen are connected by hydrogen bonding. To participate in hydrogen bonding, the nitrogen should be polar and charged, which is not the case for free nitrogen. Hence, no hydrogen bond exists.

The possibility that the HCN solves in a minor percentage in the liquid ethane and nitrogen does not lead to hydrogen bonding too. Additionally, nitrogen's solution with HCN does not produce free charges.

In conclusion, according to the models that calculate the composition of Titan's lakes, the RF signais emittedfrom the MEMS-U will be properly propagated in the liquid.

The mean surface température in Titan's environment at the equator is approximately 94 K, as it has been determined by the Huygens probe (Fulchignoni et al., 2005) and 2 or 3 K lower in the polar régions as measured remotely by CIRS [START_REF] Jennings | Titan's Surface Brightness Températures[END_REF]. MEMS devices, will be released by the Lake Lânder at low height and either they will flow on the liquid's surface or dive in the liquid. They hâve to be qualified for such low températures.

From experimental expérience using MEMS devices at low températures, the microdevices are already adaptable in cryogénie conditions. In the SPICA coronagraph [START_REF] Swinyard | The space infrared telescope for cosmology and astrophysics: SPICA A joint mission between JAXA and ESA[END_REF] the micro-scale machines are designed to operate in cryogénie conditions, while the whole telescope will be cooled down to 4.5 K [START_REF] Enya | A Micro Electrical Mechanical Systems (MEMS)-based Cryogénie Déformable Mirror[END_REF].

Any température measurement is correlated on the variance of an attribute of the exposed material with any ambient température change. MEMS resonator-based oscillators hâve been proposed to be used as Complementary metal-oxide semiconductor (CMOS)

température sensor with a resolution of 0.008 K [START_REF] Jha | Cmos-Compatible Dual-Resonator MEMS Température Sensor with Milli-Degree Accuracy[END_REF]. The MEMS pressure sensing technique is based on the Silicon piezoresistive effect, where any pressure change will cause deflection and internai stain change, which will resuit an output voltage variation (Lin & Yun, 1998;[START_REF] Li | The effect of diaphragm on performance of MEMS pressure sensor packaging[END_REF]. Usually, a Wheatstone bridge circuit delivers such voltage measurements, in order to correlate them with the applied pressure.

RF MEMS Switches are currently under development at the NASA Goddard Space Flight Center35 (Fig. 7.23).

Description of the experiment

The MEMS experiment will be developed in two phases. Phase One is the deployment of the MEMS micro-shells network (Fig. 7.24). Phase two starts when the devices begin to émit (Fig. 7.25). The deployment will take place during the descent phase of the EDL procedure of the probe. Few meters before the touchdown, the probe will release the MEMS shells to the atmosphère which will begin simultaneously to operate as micro-laboratories. This procedure will be continued after the landing, when the probe releases more MEMS directly to the liquid. Some of these devices will float on the surface of the liquid and other will dive without stopping transmitting.

Schematically, this deployment will provide a broad network of floating and submarine RF sources, while the signal receiver will be aboard the Lake Lânder. In a sense, this ambitious experiment can easily accomplish both the 3D sounding of the lake as well as its thermal profile. Subséquent to Phase One, these devices will float or dive into the liquid hydrocarbons and continuously transmit data to the Lânder depending on the duration of their power supplies.

Phase B: Operation The MEMS experiment will vertically measure the température and the pressure of the liquid deposit. Moreover, the recorded reflections of the RF signal emitted by MEMS-U will construct the topographical map of the bottom of the lake.

Although it is désirable to hâve many devices in order to establish a network, the major issue of portability cornes up. However, MEMS devices suit perfectly this kind of requirements because of their small shape, size and weight.

Seismometers on icy moons

The icy moons of the Outer Solar System possess evidence about the origin and évolution of their Systems and the Solar System in general. In terms of surface morphology, internai structure as well as their environmental uniqueness, Satum's moons Titan and Enceladus and Jupiter's Ganymede and Europa are the best représentatives to that perspective. In this section, I would like to propose a seismic experiment for the icy moons as a payload of future missions' landers. I also suggested pre-target areas with internai dynamic potential and multivariable surface expressions. Through these experiments, it will be possible to identify active régions along the satellites, which will provide important information regarding the fluid transfer processes towards the surface, as well as, determining the presence of a subsurface liquid deposit (Bampasidis et al.,201 la).

Scientific Goals and Objectives

The scientific goals of the future seismic experiment are: Both low and high frequencies can be recorded by broadband seismometers used on Earth. Newer seismographs measure ground movements smaller than 1 nm. There are several kinds of such instruments available, depending on the surface's location. With the usage of seismic instruments, we will extend our knowledge of planets' interior at the Solar System.

Additionally, both these space and terrestrial observations will help us understand better the origin and trace the evolutionary path of our own planet.

Table 7.7 below lists the main seismic instrument's requirements considering the spécifications of the shortest modem seismic instruments which operate in extreme conditions on Earth like the deep océan floor [START_REF] Romanowicz | The Monterey Bay broadband océan bottom seismic observatory[END_REF]. Similar sensors can be easily found in the global market. accelerometer can be used as a transducer to sense any weak or strong ground motion in a low frequency range (up to 100 Hz) without using extra power to operate [START_REF] Garcia | Seismic Accelerometer. Instrumentation and Measurement Technology Conférence[END_REF].

Once the sensor records a motion, a signal will be sent at the Main Processor Unit (MPU) of the System. The MPU is responsible for the operation, the administration and the maintenance of the instrument and contains the core of the application software. After receiving and recording the signal from the MPU, it will be transferred to the next component of the instrument, the Multiplexer (MUX) by using a line interface circuit, which provides the connectivity between MPU and MUX. MUX data files will be compressed, converted to the right format and prepared for transmission. Because of the continuous fonction of the seismometer, a service for bulky data is needed.

The transmission frame protocols will be defined similar to the protocols of the lander's instrumentation. Figure 7.26 shows the basic fonctional procedure of a seismic instrument to be operated on icy moons according to the requirements mentioned above. The configuration of the System is illustrated in a simplified plot. Ail the components inside the dashed line in The experimental constraints and conditions of the implémentation of seismic instrumentation on the surface of icy moons and the possibly emerged problems and difficulties are listed in Table 7.8 below. with molecular electronic transducer technology (MET), c.f. Hongyu Yu (pers. comm.) and [START_REF] Knapmeyer | Planetary Seismometers: An OverView[END_REF]. It performs high sock tolérance and therefore seems idéal when the geophysical network scénario development will be adopted. A seismic network has been proposed (see Fig. 7.28) to be part of the geophysical payload of forthcoming missions (Lorenz et al., 2009;[START_REF] Lange | Concept Study for a Titan Geophysical Network[END_REF]. A seismic network can consist of an array of sensors based on the promising MEMS technology as well as several piezoelectric transducers and new developments based on laser-interferometric sensing or convection in electrolytic liquids. [START_REF] Stein | An Introduction to Seismology, Earthquakes, and Earth Structure[END_REF]. A network of geophones can receive signais of seismic waves and thus records not only the ground movement but also the interior structure.

This network, which resembles a geophone array (Fig. 7.29) widely used on the Earth, can detect ground motions caused by natural (passive experiment) or controlled sources (active experiment). Space seismic exploration is high priority in NASA programs. Insight mission to Mars newly selected by NASA as the next Discovery mission will consist of sophisticated geophysical instrumentation which will sound Mars deep interior.

Possible locations for ground-based experiments

The idéal location for the location of seismic equipment will hâve to ensure that its recordings will represent accurately and separately every ground vibration. For this purpose, the seismic instruments on Earth should be placed in a hole of approximately 0.5 m depth. Thus, less ground noise will be recorded. Noise on an icy surface can be originated by the atmosphère due to its seasonal and diumal effects as reviewed in [START_REF] Hirtzig | A review of Titan's atmospheric phenomena[END_REF] for Titan's case.

Noise can be produced also from the local wind and meteor impacts.

In the case of Titan, its dense atmosphère protects the surface ffom meteoroid impacts and few craters hâve been observed so far (Elachi et al., 2005;[START_REF] Wood | Impact craters on Titan[END_REF]. Therefore, the seismometers will record merely interior events. On the other hand, the same instrumentation will also measure vibrations caused by impacts on Enceladus, Ganymede and Europa.

The proper contact between the feet of the seismograph and the local surface at the landing area stands for another issue to be confronted. Any particles and dust can easily perturb the seismic sensors during their fonction and cause random errors at the sensor's record. This type of surface features has been already found on Titan. The Huygens probe landed on a relatively soft solid surface analogous to wet clay, lightly packed snow and wet or dry sand [START_REF] Zamecki | A soft solid surface on Titan as revealed by the Huygens Surface Science Package[END_REF]. Moreover, the Descent Imager and Spectral Radiometer's (DISR) surface images showed rounded stones approximately 15 cm in diameter to lie on top of a finer-grained surface in variable spatial distribution [START_REF] Tomasko | Rain, winds and haze during the Huygens probe's descent to Titan's surface[END_REF]. If the icy pebbles lying over the instrument's feet move or/and melt, the equilibrium position of the seismic equipment will be disturbed and the sensors will misplace their initial orientation. The measurements should be corrected if such a micro-movement of the probe is noticed.

Since 2004, the Cassini-Huygens mission has made exciting discoveries at the Saturnian System and especially at its satellites, Titan and Enceladus. Titan may be représentative of many planetary bodies and consequently, the understanding of its internai geology will enlighten the inner structure of an entire class ofplanets and moons (Coustenis & Taylor, 2008). Similarly, Enceladus, despite its small size, it is extremely interesting due its huge geysers, which mainly feed the E-ring of Saturn. These plumes show that the moon is geologically active, hosting a possible internai water océan which is the source of its vents (Dougherty et al., 2006;Collins & Goodman, 2007).

Likewise, the Galilean satellites of Jupiter présent many similarities with rocky planetary bodies in terms of their surface features and atmospheric environments. Europa hosts a relatively stable environment, an internai océan and in combination with its young surface, the moon houses a great astrobiological potential [START_REF] Kargel | Europa's Crust and Océan: Origin, Composition, and the Prospects for Life[END_REF]. On the other hand, Ganymede seems to be a partially geologically active satellite with an intrinsic magnetic field (Pappalardo et al., 1998a;Patterson et al., 2010).

Future missions at icy moons of the outer planets will be a great opportunity to study them comparatively and give proofs of active planetary Systems. Jupiter's Ganymede, Callisto and Europa will be the main targets of a future ESA mission at the outer planets to study the emergence of habitable worlds and the Jovian System as a whole. The Jupiter Icy moons Explorer (JUICE) mission will study in depth the icy satellites of Jupiter. It will focus on detailed investigations on Ganymede, Europa and Callisto in order to understand their surface geology, detect the presence of a liquid océan in their interior and study their habitability potential.

Saturn's Titan and Enceladus could be also studied in the future to improve our understanding of their atmospheric circulation, the surface geological processes and tectonics and their internai océans, among other. The icy moons of the gas giants can also provide hints regarding the emergence oflife and the climate of the Earth.

To study Titan as a System, the Titan Saturn System Mission (TSSM) or smaller more focused concepts hâve been proposed based on interdisciplinary scientific interest in order to understand the nature of the geological playground, surface formation and its interaction with the atmosphère, meteorology and aeronomy. They could détermine the present-day structure and levels of activity of Titan and Enceladus. Such a mission could also identify beat sources, internai réservoirs of volatiles (in particular methane and ammonia) and eruptive processes and astrobiological processes which possibly occur in its organic environment The mission would also investigate the relation and the interaction between Titan and Saturn and between Enceladus and Saturn.

The utilization of Micro-Electro-Mechanical Systems (MEMS) technology will give to the future missions the opportunity to improve the exploration. By exploiting their small shape, these prospective devices can reach distant areas where traditional instrumentation is unable to visit. MEMS are some of the most innovative technologies available today. The challenges inhérent in using MEMS applications in exotic space environments will surely advance the current technological innovations. Irrespective of that, the impact of this kind of experiment in space science will be enormous, even when operating at rudimentary level.

Although several technically originated issues arise from MEMS implémentation due to the lack of long time utilization in space missions, they can be compounded by the fact of the réduction of piece-part instrumentation. MEMS can operate in extremely unfavorable conditions by overcoming any technological limitations, thus establishing a benchmark for future experiments.

The MEMS experiment I propose for the lake lander will monitor extensively the ground beneath the liquid and quantitatively define precisely its Chemical composition. As a conséquence, a MEMS experiment in Titan's lakes will appreciably enlighten our view of the largest Saturnian moon, unraveling a crucial link for understanding atmospheric dynamics and geological processes on it. This conceptual project, after surviving a rigorous testing procedure, can be part of any planetary mission establishing a path for forthcoming surveys.

The seismic experiment which I also would like to propose can identify the existence of liquid internai deposits, with a great astrobiological potential. Isolated environments that consist of water and organics, components that hâve already been identified on most icy satellites, providing idéal conditions for the survival of biological building blocks. Since only the seismic instrumentation can map the subsurface layers, détermine their composition and structure and measure their thickness, I would suggest mounting it in a future mission as part of a landing probe payload.

I believe that these sophisticated experiments should be part of the framework of a future mission on icy moons. As micro-devices are in compliance with the enhanced technological capabilities and meet the NASA and ESA régulations, they seem to be the key for solving the issues of investigation of the remotest régions of Titan, Enceladus, Ganymede and Europa.

Chapter 8

Titan and the Cassini-Huygens mission as educational resources

Growing and living in a technology-dominated world, it is indispensable for people to become confident in their ability to understand Mathematics and Science [START_REF] Fumer | Making Connections: Using GIS to Integrate Mathematics and Science[END_REF].

Alternative educational resources, different from traditional tools and techniques, can be used as supplementary material in order to attract them to these disciplines. Astronomy and Space

Science are excellent fields to use for this purpose. However, the public is occasionally In this Chapter, I describe the incorporation of Titan and the Cassini-Huygens mission as resources in formai and informal/outreach educational activities. I also présent the results of a project conducted in a secondary éducation group of pupils in Greece, using Astronomy and especially the Cassini-Huygens mission as a tool for attracting them to science tasks, testing and enhancing their knowledge, abilities and perspective.

The outcome of the outreach activities hâve been presented in the European Planetary Science Congress (EPSC) in 2006 in Berlin, Germany [START_REF] Moussas | Outreach ofAstronomy with emphasis to the Solar System by the Space group in Greece[END_REF], in the EPSC in 2009 in Potsdam, Germany (Coustenis et al., 2009d;Moussas & Bampasidis, 2009), in the International Astronomical Union Symposium (IAU) in Paris, France in 2009 [START_REF] Moussas | The gears of the Antikythera Mechanism: an educational pathfinder to the solar System[END_REF], in the EPSC in Rome, Italy in 2010 [START_REF] Coustenis | Titan trace gaseous composition from CIRS at the end of the Cassini-Huygens prime mission[END_REF] and in the European Geophysical Union (EGU) General Assembly in 2010 (Moussas et al., 2010b) and 2012 [START_REF] Moussas | The oldest computer in today's éducation: The great attractor of children to science, the Antikythera Mechanism, as an educational instrument[END_REF]Solomonidou et al., 2012), both in Vienna, Austria.

Introduction

It has been reported that the majority of students of ail levels are usually unable to change their previous beliefs and intuitions about physical phenomena, even though they hâve attended adéquate science courses offered by their national curricula [START_REF] Swann | The teaching of physics[END_REF][START_REF] Halloun | The initial knowledge State of college physics students[END_REF][START_REF] Barrington | Attitudes toward science and science knowledge of intellectually gifted and average students in third, seventh, and eleventh grades[END_REF][START_REF] Halloun | Schematic Concepts for Schematic Models of the Real World: The Newtonian Concept of Force[END_REF]. Many students continue to use the so called quasi-Aristotelian theory for analyzing nature, being led by their everyday expérience, notwithstanding the proved scientific explanations [START_REF] Disessa | Unleaming Aristotelian physics: a study of knowledge-based leaming[END_REF][START_REF] White | Sources of Difficulty in Understanding Newtonian Dynamics[END_REF].

Additionally, pupils often do not succeed in solving mathematical problems, because they do not understand the framework of a problem and therefore cannot demonstrate a correct solution [START_REF] Frykholm | Connecting Science and Mathematics Instruction: Pedagogical Context Knowledge for Teachers[END_REF].

A more effective way of describing Physics is required in order to overcome people's previous beliefs [START_REF] White | Sources of Difficulty in Understanding Newtonian Dynamics[END_REF]. Educators hâve to inspire pupils in demanding modem leaming environments and give them the opportunity to engage in productive discussions pertaining to challenges of Science and Technology. Motivating pupils to use contemporary scientific and technological achievements could accomplish significant improvement of teaching effectiveness. The teacher should go beyond the traditional way of teaching and focus on a more student-oriented approach [START_REF] Heacox | From Differentiating Instruction in the Regular Classroom: How to Reach and Teach Ail Learners[END_REF]. Obviously, a strong motivator is needed, which can make the children visualize Science and understand the scientific phenomena as well as the embedded mathematics.

The past academie year (2011-12), I had the opportunity to teach Science lessons in a State school of Athens, Greece. My students had not attended Science lessons before and it was a great opportunity for me to incorporate the Cassini-Huygens advents in the secondary éducation school curriculum. I tried to test whether such a topic could enhance the effectiveness of my teaching and eventually evaluate the results in order to check the degree to which I succeeded the target.

Astronomical issues can be used as an essential pedagogical tool towards this goal. The Astronomy field has already been successfully selected for conceptual change in Science [START_REF] Vosniadou | Designing curricula for conceptual change restructuring: lessons from the study of knowledge acquisition in astronomy[END_REF]. The Astrophysics Laboratory's Space Group of the University of Athens, and especially my Thesis supervisors Prof. X. Moussas and P. Preka-Papadema hâve a long-term expérience in educational and outreach activities. Together with my supervisors, we considered the possibility of implementing the Cassini-Huygens mission concept, which is the most successfiil human exploration project, to attract students and the general public in the disciplines of Science and Mathematics [START_REF] Moussas | Outreach ofAstronomy with emphasis to the Solar System by the Space group in Greece[END_REF]). An idéal educational and outreach resource for our purposes should be intriguing and both Titan and the Cassini-Huygens mission are indeed attractive objects.

From the ancient years, people hâve been always intrigued to seek for terrestrial-like analogues and habitable worlds throughout the Solar System and beyond. As I hâve described in the previous Chapters, Titan satisfies these prerequisites since its environment présents many similarities to the terrestrial one and it is also one of the prime targets for astrobiology research.

Since 2004, Titan is under deep investigation by the Cassini-Huygens mission, which will last until 2017. The mission's discoveries are often in the headlines worldwide. This material is easily accessible through the World Wide Web, so students, teachers and the layman public can obtain a wealth of information about the Satumian System.

The Cassini-Huygens mission is a state-of-the-art resuit of efforts in Science and

Engineering and its findings hâve advanced significantly our knowledge of the Satumian System. It is a perfect example of the contemporary technology achievements of the application of Science in real life situations. The mission is in the forefront of NASA outreach activities. NASA has already scheduled the lOth édition of "Cassini Scientist for a day", a global annual contest37, which gives the opportunity for students of ail grades (7-18 years old) http://satum.jpl.nasa.gov/education/scientistforadaylOthedition/ to interact with the Cassini material and the local Cassini experts. My involvement in this event will be described later in this Chapter.

Titan and the Cassini-Huygens mission are exceptional tools for astronomical educational and outreach activities since they combine most of the prerequisites of idéal educational resources: intriguing concept, accessibility, interactivity, illustrative knowledge transitivity, applicability to real life situations and graduai approach to knowledge. Finally, it is worth noting that the interdisciplinary character of the Cassini-Huygens mission is an additional positive element for Science éducation and outreach.

Titan and the Cassini-Huygens mission in informal/outreach activities

The starry sky has always seemed mysterious and familiar at the same time to humans from the ancient era to modem times. Today, from ground and space télescopes, a wealth of data revolutionizes our knowledge for our living planet and the entire Universe in a way never imagined before. This digital évolution, currently available worldwide, allows the humanity to benefit from ail the space applications. The introduction of young people and the general public into the space field seems essential for not only cultivating astronomical culture and skills, but also for becoming more familiar with scientific work.

In collaboration with the Space Group of the University of Athens, I hâve organized, participated or attended numerous outreach activities focalized on the general audience in order to conceptualize astronomical phenomena and especially Satum and Titan to enhance people's knowledge and perceptions. These public events encourage and help young people to love Astronomy and the Solar System, to understand why we need to study Titan and the Satumian System and to develop their critical thinking, self-expression and Creative talents.

An overview of these events is listed in Table 8.1 below.

Engaging Cassini-Huygens and Titan in outreach activities has proved very efficient.

The concept is to use ail the available material to show to the layman public that the human progress is strongly related with the advents in Science, Mathematics and Technology.

Beyond its tremendous importance for both Science and Technology, the Cassini-Huygens mission can be successfully employed as an attractor of the general public to Science, Mathematics, Engineering and Technology during outdoor and indoor events.

For this purpose, we perform a Virtual time travel from the roots of the Astronomy to the contemporary space era. Two are the major milestones of this joumey: the Antikythera Mechanism, which is the most sophisticated astronomical device survived from the Antiquity and the Cassini-Huygens mission (Moussas & Bampasidis, 2009;Moussas, 2010a;2011a;2012). Through an introductory lecture, the public becomes capable to visualize the concept of the spacecraft, its operation and usage and leam the high engineering capabilities of it.

People can easily visualize the celestial movements of the Solar System and comprehend the fundamental principles that govem the Universe.

I hâve strongly supported this spécifie outreach field because of my broad knowledge of Titan science and the Cassini-Huygens mission as well as my personal interest in both of them. Usually in the framework of each event, I perform also a sériés of simple science experiments, which are the offspring of the collaboration among the Space Group, the British Council in Athens and the Science Communication Unit of the University of West England in Bristol38.

Prof. X. Moussas and I were the co-leaders of the TandEM Outreach working Group.

Distinctive activities/events hâve been designed for this study. Our plans of outreach activities included interactive web pages, images and video related to the mission (from its developmental stage until the nominal operation at Titan and Enceladus). We will exploit the wide spread of social media especially to interact with the young people. Conventional activities hâve been also considered such as newsletter release, TV broadeasts, spécifie publications, lectures, CD/DVD release, card games, collaboration with communities.

The planned outreach activities will be divided into three phases: (a) before the launch of the mission, (b) during the cruise to the Satumian System and (c) during the nominal mission. As for the first phase, the purpose of the outreach is to keep the layman public interest alive. Towards this goal, the discoveries of Cassini-Huygens mission, which is planned to operate up to 2017, will be used. During the spacecraft cruise to the Satum System, the outreach activities and events will exploit the first inflight data of the mission from its flyby maneuvers and the Juno results from Jupiter.

Going beyond the level of démonstration, with the contribution of local scientifïc and amateur astronomers who are considered to be effective Astronomical and Science ambassadors [START_REF] Berendsen | Conceptual Astronomy Knowledge among Amateur Astronomers[END_REF][START_REF] Gibbs | Effectiveness of Amateur Astronomers as Informai Science Educators[END_REF] Through these events, people are empowered to conceptualize various astronomical assets by making their own observations and leam the usage of scientific instrumentation, like the telescope and the coronagraph. By these means, Science and Astronomy approach the public and equip young people with critical thinking skills (Coustenis et al., 2009d;[START_REF] Coustenis | Titan trace gaseous composition from CIRS at the end of the Cassini-Huygens prime mission[END_REF]. In particular, children can expérience astronomical observations by watching the Sun during the daylight whereas during night, they watch the Moon, the planets, the stars and the constellations.

In addition, children are strongly encouraged to design posters and perform oral présentations in subjects related to Physics and Astronomy. In the meantime, scientists can easily emphasize the importance of funding Science as an investment for society and as a key factor for improving the quality of human life.

Through these activities, pupils enjoy exploring the fundamentals of Science and Technology, which might hâve a positive long-term impact on their careers and lives [START_REF] Fumer | Making Connections: Using GIS to Integrate Mathematics and Science[END_REF]. The general public has the opportunity to meet experts in space science, scientists, astronomers and engineers and inquire them about astronomical assets conceming the space research, the Earth and the évolution of the Solar System and Cosmology.

At these events, visitors get an up-close view of scale models of spacecraft constmcted for young children and are informed about space agencies' missions and their future projects for the exploration of the Universe. Through these events and exhibitions we aim to empower social interactions, encourage the school community to get involved into Space Science and Astronomy and discuss the merits of major global issues like global warming and pollution.

The majority of these outreach events hâve been presented in numerous conférences worldwide. Table 8.1 below lists the outreach activities I hâve been engaged. The feedback from the people attended these activities and especially the children encourage me to adapt similar techniques to formai éducation practices, in which I am engaged since 2001. Due to their interdisciplinary value, Titan and the Cassini-Huygens mission can be easily incorporated into the national science curriculum material. The astronomical knowledge has significantly increased with time and people hâve become familiar with scientific facts. Since serving the social benefit is the major commitment of Astronomy, it seems crucial in the contemporary harsh environment to approach the layman public [START_REF] Clough | The Story Behind the Science: Bringing Science and Scientists to Life in Post-Secondary Science Education[END_REF]). People's perspective of Science and Astronomy in particular is a significant concem of modem scientific communities [START_REF] Wallace | A Study of General Education Astronomy Students' Understandings of Cosmology. Part IV. Common Difficulties Students Expérience with Cosmology[END_REF].

By general consensus, conventional methods for making science accessible to the public hâve failed. However, due to the large impact of pseudoscience on people of different social and cultural backgrounds, it is impérative to teach real science to the public in order to clarify any misconceptions.

Astronomers and scientists hâve to produce innovative methods to attract people and bring Astronomy at the forehead of éducation. This can only be achieved by following successful pedagogical principles and tools, aiming to make Science accessible and promote lifelong Science leaming [START_REF] Linn | WISE design for knowledge intégration[END_REF]. An example of this kind is the use of films [START_REF] Efthimiou | Avatars of Hollywood in physical science[END_REF], or user-friendly computer software [START_REF] Fumer | Making Connections: Using GIS to Integrate Mathematics and Science[END_REF][START_REF] Keating | The Virtual Solar System Project: Developing Conceptual Understanding of Astronomical Concepts Through Building Three-Dimensional Computational Models[END_REF], visual représentations [START_REF] Gazit | Emerging Conceptual Understanding of Complex Astronomical Phenomena by Using a Virtual Solar System[END_REF] or the wealth of the Internet [START_REF] Cohen | Internet Links for Science Education: Student-Scientist Partnerships[END_REF]. The outreach activities focus on both children and the general public. The real focus, however, should be on the development of Creative and critical thinking throughout the society.

Conceming the educational process, the interaction of the scientific community with school educators improves their teaching skills and effectiveness by including astronomical assets to the design of their lesson plans. In the contemporary school environment, Astronomy is indeed an efficient tool which allows the educators to self evaluate their work on Science lessons and intrigue their students in a leamer-friendly environment.

Pupils understand the existence of multiple pathways towards knowledge, develop the ability to express ideas about Physics and Astronomy and eventually think outside the box.

Adequately trained teachers readily act as local science distributors to their pupils and community. Moreover, more students hâve decided to follow a scientific career, being inspired by the interdisciplinary value of astronomical issues.

Advances of Science hâve always proved to be bénéficiai to the humanity and people's quality of life. The outreach projects try to strengthen public interest in Astronomy, by providing alternative paths of leaming and expanding their expectations. The participation of people in the events we hâve organized or participated is increasing, showing their interest in Space Research matters (Fig. 8.2). Through Science, people move beyond the barriers imposed by pseudoscience and start negotiating with their issues using a thoroughly scientific way. Teachers can enrich their inventory with the proper tools to effectively empower their pupils with ail the necessary science and mathematical skills. Indeed, Astronomy and Astrophysics stand as excellent attractors of pupils to Science and Technology.

Titan and Cassini-Huygens mission in formai éducation

Astronomy can be used in formai éducation, not only as an independent discipline but also as an attractor to Science and Mathematics and as an évaluation tool for testing pupils'

knowledge and perspective about fondamental Science. In order to attract young students from ail levels (7-18 years old) A. Solomonidou, Prof. X. Moussas and I became the coordinators of the very successful NASA student contest the "Cassini Scientist for a day" (Solomonidou et al., 2012) the concept of which is described in the next section. This international contest is focusing on the Cassini mission for the last 10 years with Satum's Titan always being in the center of attention. The goal of this research was dual. I first applied this project to the pupils in order to introduce them to Science issues, which they would face in the next grade. In my point of view, creating positive perspective to Science before pupils begin their science lessons is crucial and Titan and Cassini-Huygens mission provide the essential tools to achieve it.

The second goal of this research was to produce a negotiated interaction between the students and an exchange of information. Through this procedure I was able to evaluate the contribution of each pupil in his/her group and, eventually, self-evaluate my project and the impact of Titan and the Cassini-Huygens to the students. In the forthcoming academie year, this small-scale study will be distributed among the teachers throughout Greece.

8.4.1

NASA -Cassini Outreach school contest: Scientist for a day NASA holds a very active, productive and student oriented Outreach Group. For the last 10 years the members of this Group focus in maintaining in high levels the interest of teachers and students towards the joint NASA-ESA successful mission Cassini-Huygens with the school essay contest "Cassini Scientist for a day". This mission has attracted global attention and the Outreach Group focus in the international participation in every contest or event regarding Cassini.

The "Cassini Scientist for a Day" contest is a chance for students of ail grades ( 7

8.4.3

The method I adapted cooperative group-based activities by using worksheets for three 45-minute teaching hours. [START_REF] Cooper | Student involvement in leaming: Cooperative leaming and college instruction[END_REF] described this type of educational practice as a structured, systematic instructional strategy in which small groups work together towards a common goal. Although students in Greece are not used to working in groups [START_REF] Alexopoulou | Small-Group Discussion in Physics: Peer Interaction Modes in Pairsand Fours[END_REF], this particular State school has been selected from the Ministry administration as a pilot-school to adapt such educational practices. Hence, the target group is quite familiar with group-based work.

The pre-task activity consisted of an introductory lecture, which explained the tasks and linked the fortheoming activity to the Cassini contest. This lecture lasted for about ten minutes. After having attended the pre-task activity, the students were asked to form selfselected groups of 3 to 5 members. They formed 8 groups, 4 of them consisting of 5 students, 3 groups of 4 students and one group of 3 students (Table 8.2). Each student had his own worksheet in which he was asked to write down the response of the whole group to the activity's tasks. On the other hand, only one sheet with planetary facts was given to the group as supporting material. By this approach, ail students were forced to participate by filling their worksheets and interact with their team members using the common material. The worksheet was divided into 3 individual activities. A short description of the worksheet is given in the Table 8.3 below. Acting as a facilitator, I had a minor intervention, mainly observing and monitoring students' work rather than participating as a group member, so as not to influence the final outcome.

Before starting the activity 3, I performed a short démonstration of two experiments, similar to the ones I performed in informai outreach activities. In particular, a simulation of a rocket launch and one of zéro atmosphère pressure were demonstrated by using every day materials. An empty film canister played the rôle of the rocket body. When I lit half of a commercial fizzing antacid tablet and few drops of water, their Chemical reaction produced carbon dioxide, which in tum gave the canister the proper impulse to launch for a few meters.

Then, I put in a kitchen vacuum packer an inflated balloon or preferably marshmallows and then I pumped the air out from the container. The latter experiment demonstrates what happens to the astronauts in space. More details about both experiments can be easily queried in the internet. This type of experiments can be easily performed by the students as well,

showing them the strong connection between Science and real life situations.

The task cycle closed by asking the students to answer an on-line questionnaire related to the whole project. Most of the questions follow the 5-point Likert scale [START_REF] Likert | A Technique for the Measurement of Attitudes[END_REF], and evaluate the group activity by testing their satisfaction. The questionnaire can be accessed in the following website: http://users.sch.gr/gbabasides/joomla/index.php?option=com_ckforms&view=ckforms&id=2

Problems related to group-based activities

Working with groups has a bénéficiai impact in leaming process especially in Science disciplines as well as to the development of the communicative skills of the student [START_REF] Heller | Teaching problem solving through cooperative grouping. Part 1: Group versus individual problem solving[END_REF]. In this group-based project, I hâve faced the common difficulties mentioned in literature: the free-ride effect and how to evaluate each group member individually.

People participating in groups lower their sense of individual responsibility [START_REF] Armstrong | Natural Leaming in Higher Education[END_REF] and usually the free-rider effect takes place when skilled members do most of the work [START_REF] Kerr | The dispensability of member effort and group motivation losses: Free-rider effects[END_REF][START_REF] Karau | Social loafing: A meta-analytic review and theoretical intégration[END_REF].

The outcome of the group work in formai éducation can be easily evaluated whether it meets the initial goals or not. However, it is difficult to estimate the degree of participation of each individual member of the group in the final resuit. The same issue has been reported in other disciplines such as English as a foreign language [START_REF] Foster | A Classroom Perspective on the Negotiation of Meaning[END_REF], in which group-based activities are a widespread practice. For a successful outcome, students should be convinced for the bénéficiai purpose of the task and be serious and committed. Additionally, each group member must hâve a spécifie rôle in the task [START_REF] Ellis | Task-BasedLanguage Learning and Teaching[END_REF].

Data analysis

The activity was performed as planned within the time limit of three 45-minute teaching hours. Ail the students participated, submitted the worksheets and filled in the online questionnaire. Ail the answer sheets were presented in a whole-class context.

The favorite planet for a future travel was Mars which gained 23% of the students, while the Satumian System was less favorite (20%). Students explained that they selected Mars due to its proximity to the Earth, the results of current water research and the possibility of life existence on it. On the other hand their favorite object was the personal computer at 30%, followed by the mobile phone at 25%. It is worth noting that none of them questioned the possibility of a mobile phone functioning in space. These choices show the degree to which the everyday use of technology influences contemporary teenagers. Finally, although they are technologically influenced, 30% of them did not like the web search activities.

As far as the outcome of the online questionnaires is concemed, Table 8.4 below lists the results. This questionnaire illustrâtes the satisfaction of the students after finishing the group task. I hâve decided to give a more general census questionnaire in order to avoid répétitions ffom the contest. Most of the students were satisfied (91%) from the activity and the concept (79%), while the majority of them (96%) found the subject interesting. Almost half of the students (47%) used successfully the knowledge obtained from their elementary studies, which means that they believe they hâve a good basis for the forthcoming class. Sixty-five per cent of the students feel more confident to solve similar problems in the future. About 65% of the pupils would like to attend lessons devoted to Astronomy and prefer more astronomie subjects in Science lessons. The students' favorite activity was the design of a scaled solar System by a percentage of 31 followed by the experiment démonstration at 20%. Ail the students were satisfied by their participation in the NASA contest.

About the feedback for the teacher, the vast majority (94%) appreciated his facilitating rôle to the activity. Additionally, most of the students understood the task instructions (86%), which means that they are well structured and related to their âge and educational level.

Almost three quarters of the participants enjoyed the experimental démonstration, which helped them understand the concept of activity 3. Conceming the performance of group work, half of the students cooperated effectively and interacted with the rest team members.

Conclusions

As described above, the interdisciplinary value of Titan and the Cassini-Huygens mission is apparent and the public interest worldwide is enormous. This can be used as a magnificent educational instrument, not only to model physical phenomena, but also to become involved in the whole curriculum of the learning procedure. Titan and the Cassini-Huygens mission can motivate children and the general public to understand modeling, identify the scientific way of thinking and initiate or encourage their interest to scientific disciplines. The interaction of the public with the Cassini-Huygens concept emphasizes the technological applications and their large impact to our lives. Without doubt, Cassini-Huygens seriously and generously contributes effectively in promoting Science and Technology to the broader public.

The proposed group work based activity follows the concept of the formative assessment [START_REF] Ainsworth | Common formative assessments[END_REF] in which the teacher applies procedures during the learning process to improve his students' performance. This activity focuses on a qualitative feedback from the student. The high percentages of satisfaction, in combination with their will to implement more astronomie aspects in the Science lessons, encourages me to work towards this direction.

Although the involved students were familiar with working in small groups, only 50% of them felt that they had cooperated properly. This is explained by the fact that the groups were self-formed. Since the free-rider effect had appeared, it would hâve worked out more efficiently if I had organized the groups according to their level in order to form groups of equalpower [START_REF] Davies | Groupwork as a form of assessment: common problems and recommended solutions[END_REF].

This activity has significantly enhanced the students' self-confidence and their An additional, new contribution from our work was that the observed increase is followed by a strong decrease of the gaseous Chemical content within the next terrestrial year that remains to be confîrmed with forther data but which has strong implications on the seasonal variations and the behavior of the northem vortex.

To perform this analysis, I hâve made use of data obtained from Cassini/CIRS on which I hâve applied the radiative transfer code (Atmospheric Radiative Transfer for Titan -ARTT) I hâve upgraded. ARTT fonctions in an enhanced and optimized way, providing an improved fit to the data, thanks to the updated haze and spectroscopic parameters. Thus, ARTT can search for 18 additional molécules and isotopologues compared to the previous version, includes recently detected molécules (Nixon et al. 2008a;[START_REF] Raulin | Planetary science: Organic lakes on Titan[END_REF], Jennings et al. 2008) and allows for forther research conceming the presence of more complex and new weak species or the identification of more minor isotopologues.

Moreover, I was involved in the évaluation of CIRS data inferences (current pipeline, DS4000 and Grand Average) by testing the outcome of new calibration algorithms applied by the CIRS calibration Team in Goddard. I concluded that a new release of DS4000 database, the Phase-Corrected, should be exploited in future studies using CIRS data for ail users. If the latter database is not available, the DS4000 still constitutes an improvement with respect to the current pipeline (v3.2). The analysis of atmospheric and surface Titan data has also enhanced our understanding of the habitability potential of the satellite as it gives better access to the properties involved in the environment and the stability of the System for supporting the emergence of biological building blocks. The experiment I hâve proposed, to be included as payload of a future Lake Lânder probe incorporating Micro-Electro-Mechanical Systems (MEMS), could measure the thermodynamic parameters of the liquid and produce a 3D topographie map of the lake's bottom [START_REF] Bampasidis | Sounding the interior of Titan's lakes by using Micro-Electro-Mechanical Systems (MEMS)[END_REF]. I hâve also proposed a seismic experiment which will be able to identify the existence of liquid internai deposits, with a great astrobiological potential (Bampasidis et al., 2011a).

Finally, the interdisciplinary value of Titan and the frequent, amazing discoveries in the Satumian System by the Cassini-Huygens mission can be used as a magnificent educational instrument. I hâve adapted the achievements of the Cassini-Huygens mission on Titan to motivate children and the general public to get involved in scientific disciplines related to Astronomy.

A.l

The quest for the proper solution

In order to understand whether we can proceed to a direct calculation of the above équation, we should firstly formulate the mathematical framework of these problems. To begin with, the class X of fonctions f(x) is considered as Hilbert space. Both the set of the linear fonctionals and the set of the values of these fonctionals are also defined in the class X, such as:

Fm(f)-gm,n=l N.
Since X is a Hilbert space, the fonctionals should be continuous and they should be the scalar inner product of the fonction / and the fonction (pn E X such as [START_REF] Bertero | Linear inverse problems with discrète data: I -General formulation and singular System analysis[END_REF]:

UA)x =gn,n=l,...,N.

The solution of such an intégral in order to uncover the sought fonction /(x) can be accomplished by using numerical methods.

When one solves the linear Fredholm intégral équation of the First kind, one should take into account that this form is an ill-posed problem. As firstly defined by Hadamard When ail these conditions are satisfied, these problems are called well-posed. On the other hand, when part of the boundary data is not available due to the insufficient number of measurements or heterogeneities and singularises, the problem becomes ill-posed [START_REF] Xin | Globally convergent numerical methods for coefficient inverse problems[END_REF]. In these kinds of problems the solution is extremely sensitive to arbitrarily small oscillations of the System and hence, the data is not really a solution in a physical sense.

Indeed, the solution amplifies the noise producing a large and wildly oscillating fonction which eventually covers the physical solution. As a conséquence, a unique existing solution is entirely corrupted by a small error in data, which apparently refers to a loss of information.

The whole issue is linked with the compactness of the operator associated with the Kernel fonction which cannot hâve a continuous inverse [START_REF] Kress | Linear intégral équations[END_REF]. In our case of Titan's inffared radiation, the high frequency components in the fonction /(x) are smoothed out by the intégration with Kernel K(y,x) in équation (4). Inversely, the calculation of the f(x) fonction fforn g(_y) will amplify the high frequency in g(y) [START_REF] Hansen | Numerical tools for analysis and solution of Fredholm intégral équations of the first kind[END_REF][START_REF] Bertero | Introduction to Inverse Problems in Imaging[END_REF].

Therefore, an approximate, but stable and reliable, and eventually meaningfol solution should be quested by applying additional parameters to the problem which describes the expected physical properties of the System. We should also seek to filter out the high frequency components of the data in order to avoid instabilities on the solution.

To do so, the intensity recording of the instrument 7(^u,v) or the g(>>) fonction should hâve a linear relation with B(z,T) or f(x) fonction which will be correct only when: a) g(y) fonction is related to /(x) by form ( 5)

b) The recordings on the detectors should hâve the same or a constant ratio to the yn value of g(y) fonction. Then, the équation ( 5) can be written as: 8n(y) = f p(y" -y) f K{y,x)f{x)dx dy n = \,...,N (6) where P{yn -y) = ô(y-yn).

In this point of our discussion, it should be mentioned that the latter ( 6) équation ignores the noise contribution to the measurements. Eventually, each value of g(y) fonction of the instrument's recordings is a linear fonctional of the f(x) fonction. In other words, g(y)

numbers are the norm [/(*)] from the vector space Â"(j;,x)to its field of scalars. Therefore, the équation ( 4) can be solved by numerical analytical methods.

A.2 Linearization

The radiance and the transmission fonction Tr(fj,,v,z) are known values by the CIRS recordings and spectroscopic databases respectively. The Planck fonction B(v,T(z)) remains the unknown parameter. Therefore, the équation ( 4) should be solved to dérivé the 5(v,T(z))

values. The solution of the radiative transfer équation ( 4) can be approached as follows (Vinatier, 2007): /"-/v5«4WW

but since AI = Iv -I°Vl and if we define K" as:

1 BB(yt9T) a/.y w r)/ A/= V-^AT ATjATr(j) j-1 ^A/,. dB{vi9T(Zj)) krrf = -L - i-AT(vi,z i) ij AT, dT, r 1
we can hâve the linear form: AI = KAT (7)

A.3 The Ky matrix

To begin with, we assume that the température is a depended quantity by the altitude rt and the latitude . Thus, we can define the array of n-element vectors:

-T, where i = and j = l,...,q, which means that we divide the atmospheric région in N layers.

Moreover, since the measurements are remotely recorded, errors due to orbiter's position in respect to the planetary body should be also considered. Furthermore, the atmospheric pressure is either not accurately known on a constant gravitational potential surface or outside the sounding région. Therefore, we introduce a single shift §) in ail tangent heights at a given tangent point latitude(ç?,),, when 1=1,...,p. CIRS spectral radiance recordings dépend on frequency v, altitude rt and latitude cpt (Achterberg et al., 2008). Then, we defïne the m-element vector:

where k = 1 . h = \y^"(p,\

Secondly, the quantities to be retrieved are incorporated to a single column vector x the éléments of which are the arrays of température Tj, the neperian logarithm of the K -jjLj,.. , M|,..., iM),...,Up J

Having a reference set of atmospheric parameters and in our case the température T°{zj)

derived by the a combination of CIRS limb and nadir data and Voyager 1 radio occultation results (Flasar et al., 2005) or by the Huygens Atmospheric Structure Instrument (HASI) (Fulchignoni et al., 2005), we can consider a forward model h(x) and by expanding it by applying the First order of Taylor sériés we dérivé:

h(x) = h(x°) + K(T -T°) h(x)-h{x°) = KAT Ah(x) = KAT
where /z(x)are the measurements and /z(x°)is the model using the reference values T°{zj)

and which is the same expression as (7).

A.4 Numerical Solution

If we followed the classical least squares approach, since the problem has become at the linear form of ( 7):

AI = KAT
we would minimize the residual:

|A7 -ATAZ'|2

AI and AT are the perturbation quantities with respect to a reference température T°(z) , the radiance /0(v,) is calculated using the reference value of T°(z). The matrix Kÿis the functional dérivative at vt with respect to T at level z.

ar-ry-ry

A/W(v,)-r(v,)

The solution should confront the same problems with the original équation. The regularization technique is usually applied in solving inversion problems and Tikhonov's approach in Hilbert spaces is one of the most famous [START_REF] Tikhonov | Solution oflll-posed Prohlems[END_REF] by introducing the parameter y as a constraint which détermines the relative weight which is applied. The regularization can be applied in both the intégral and linear System derived from integral's discritization since the effect on the smooth solution is the same [START_REF] Hansen | Numerical tools for analysis and solution of Fredholm intégral équations of the first kind[END_REF]).

Let's expand the perturbation in température AT in a set of basis vectors such as [START_REF] Conrath | Thermal Structure and Para Hydrogen Fraction on the Outer Planets from Voyager IRIS Measurements[END_REF][START_REF] Hanel | Exploration of the Solar System by Infrared Remote Sensing[END_REF]:

AT = Fa
where by F is the matrix the columns of which are the basis vectors and a is the vector of the expansion coefficients.

Hence, we minimize the fonction:

||A/-AÀ7f + ||}of which means that according to Tikhonov's approach we can minimize the following penalty AT -FFTKT\yE + KFFTKT J1 AI (8) which can be written as:

AT = SKt [yE + KSKt J1AI (9) where S = FFT, a two-point corrélation matrix of the basis vector F (e.g. [START_REF] Conrath | Thermal Structure and Para Hydrogen Fraction on the Outer Planets from Voyager IRIS Measurements[END_REF]. AT = VA1 (10) We then separate the stratosphère into N levels, which in Titan's case are 200 (Achterberg et al., 2008). Now, we are able to apply a constrained inversion algorithm, by assuming local thermodynamic equilibrium (LTE).

Once a System is supposed to be in LTE, although the intensive parameters of the System still change within space and time in each of the 200 levels, they vary slowly enough to support the thermodynamic equilibrium assumption. Then, its thermodynamic condition is determined only by its température which describes both the distribution of the velocity of the contained particles and the radiance of the radiation as a function of frequency. Thus, the radiation field within each atmospheric layer approaches the black body émission.

Hence, the recorded energy can be approximated as a function of the thermal radiation and the vertical distribution of the absorber. The total contribution of the outgoing infrared radiation of the atmosphère consists of the partial contribution of each atmospheric level. The région of the maximum contribution is located where the optical depth equals unity.

Due to the nonlinearity of the problem an itérative process of ( 10) is necessary. Only the corrélation matrix S = FFT of the basis vector is présent in the final expression and the basis vector. The error covariance matrix due for the température retrieval due to the instrument's noise propagation is [START_REF] Hanel | Exploration of the Solar System by Infrared Remote Sensing[END_REF]:

R = VEVt

If we assume that the existing random measurement errors at two random points in the spectrum are uncorrelated, the diagonal éléments of the (mxm) matrix R equal the square of the Noise Equivalent Spectral Radiance (NESR) of each measurement associated with the CIRS instrument (Achterberg et al., 2008). In fact, NESR is the signal of the instrument when the signal-to-noise ratio equals unity [START_REF] Hanel | Exploration of the Solar System by Infrared Remote Sensing[END_REF]. The rms error estimate for the i-th level is:

«r, =(2vîE/ j For the CIRS averaged spectra, the effective NESR is derived by the: NESRji)

Vn

where N is the number of spectra. The error estimate aT refers only to the précision or reproductivity of the retrievals in the presence of random measurement error [START_REF] Conrath | Thermal Structure and Para Hydrogen Fraction on the Outer Planets from Voyager IRIS Measurements[END_REF].

The functional dérivatives or the contribution fonctions or kernels are determined from the équation mentioned in §4.4.2, the Ktj = B(z)W(z,/*)•

However, if we assume that the température dependence of the Planck fonction is strong compared to that of the transmittance the Kÿ can be written as [START_REF] Hanel | Exploration of the Solar System by Infrared Remote Sensing[END_REF]:

K cB[vj,7'(Z;)] ST " cT, âZj

The vertical profiles of the température retrieval correspond to a spécifie spectral query in FP1, FP3 and FP4 part of the recorded spectrum. We directly apply this profile to fit the V4 methane band in the FP4 data and then we apply it to the other two focal planes, as it is described previously in the direct solution of the radiative transfer équation. With the 2006 upgrade, the user became capable of choosing between constant-toheight abundances and vertical distribution mixing ratios for the model calculations.

Additionally, in this release, the software simulation is divided into 3 separated spectral régions, following CIRS data focal planes, FP1, FP3 and FP4 data in order to save valuable computational time. Moreover, two more molécules were added to the source code of ARTT: a) Benzene (CôHô) from Dang-Nhu et al. lists [START_REF] Dang-Nhu | Spectral intensifies in the v4-band of benzene at 15 pm[END_REF] was added to the code, followed its détection in Titan's atmosphère (Coustenis et al., 2003).

b) Monodeuterated acetylene C2HD V4 band at 451-581 cm'1 and V5 at 600-761 cm'1 also was added from LISA (Laboratoire Interuniversitaire des Systèmes Atmosphériques) spectroscopic database (http://www.lisa.univ-parisl2.fr/).

Acetylene-dl was firstly detected on Titan in 2008 (Coustenis et al., 2008).

The latest upgrade before our effort was the 2007 one. In this upgrade, two more molécules b) The 13CH3D isotopologue of monodeuterated methane was added at 1148 cm'1 following its détection [START_REF] Bezard | Détection of 13CH3D on Titan[END_REF].

Hydrogen cyanide (HCN) spectroscopic Unes fforn GEISA 97 (Jacquinet- [START_REF] Jacquinet-Husson | The 1997 spectroscopic GEISA databank[END_REF] were replaced from HITRAN 2004 (HIgh-resolution TRansmission molecular AbsorptioN database) (Rothman et al., 2005) and later by HITRAN 2008 files (Rothman et al., 2009).

B.2 ARTT Installation

In order to install the Atmospheric Radiative Transfer code for Titan (ARTT) the following software is prerequisite:

• Fortran compiler

• Gnuplot

• the make command

• an archive manager

As far as the hardware is concemed, the workstation should hâve:

• at least 2GB ram capacity • distrib: This folder contains the vertical profiles as derived by General Circulation Models (GCM).

• profils: This folder contains the thermal profiles.

• src: This folder contains the source files needed for the installation of ARTT.

Adopt GEISA linelists

The spectroscopic files, adopted from GEISA database are accessible in GEISA website. The steps are the following order:

We hit the Geisa-2009 link, then the Interactive access link and fmally the Database Extract. The latter opens a webpage, which is divided with three different boxes. In the first box we select the molécules or the isotopologues we would like to download, in the second box we define the spectral range we are interested in and in the last box we check the columns we want. The following procedure should be done for every different molécule or isotope separately.

Each GEISA entry has its own id number instead of their name and therefore we should find the correct code from the GEISA list. This list is also posted in GEISA web site.

As far as the spectral range is concemed, it is not necessary to download different files for every spectral région (FP1, FP3, FP4). ARTT has the ability to read from the same file using the right parameters, besides the spectral range.

From the available parameters for of GEISA line lists, we select the A, B, C, D, F, G, I, parameters to be included in the download: This program is executed in a terminal by command "gfortran -o read_hitran read_hitran.f" where read_hitran is the name of the exécutable file which will be created, while and read_hitran.f is the name of the source file containing the code above. Then we should execute the command "./read_hitran". The output of this program is the HITRAN file in the correct ARTT format. Important note: during the execution of the read_hitran program we will be asked to enter the name of the output file, we must not enter the same name as the input file.

Add a new molécule in ARTT After retrieving the desired spectroscopic files from the online databases, the incorporation procedure is the following:

1) We create a copy of the installation folder of a previous installation and we will make ail the subséquent changes in this folder. It is essential to keep a backup of the previous installation folder in case of a data loss.

2) We add the data file from GEISA or HITRAN in the folder /db and we edit it in order to change the number corresponding to the molecule's database id (it is the number in the final column in most cases). We replace this number with a molécule id, which doesn't exist in our code. In order to fmd ail the existing molécule numbers it is better to search in the spectroFP 1-2-3.don files first. Altematively, we can edit the spectro.f file but it is a bit more risky. If we update a molécule, which already exists in our database, we don't hâve to change the molecule's id, since the new file will replace the old one.

3) The following procedure must be done for the files spectroFP 1.don spectroFP2.don spectroFP3.don:

• We open the file spectroFP*.don and we augment the number before the following text "nombre d'associations fichier-molecule" by one.

•

We add a new line with the idgeisa we added before. The wninf and wnsup columns must hâve values that cover the desired spectrum.

•

In the fichier column we put a number which must be unique and less than 65 (this number will be used later for the identification of the data file). It is crucial that the name of every molécule or isotope will be only 6 characters long.

4)

Then we open the /src directory and edit the files "trans_rad.fi, "linelists.fi and "distrib.fi.

5)

In the file trans_rad.f in the line 22 we change the value of the variable named mcorps in the new number of molécules in our radiative transfer code. In the file distrib.f we must change the format identifier 627 to ensure that ail the new molécules will be printed.

6)

In the file linelists.f we hâve to be very careful. We add a couple of lines like the following : fichier( 56 After the installation the number of molécules in the line 21 in the .mod files should be edited.

Note: This manual is based in Stefanos Stamogiorgos Bachelor dissertation.

Latitudinal variations in Titan's stratospheric thermal and Chemical structure hâve been reported in the past from Cassini data acquired by the Composite Infrared Spectrometer (CIRS) (Flasar et al. 2005;Teanby et al. 2006;Coustenis et al. 2007Coustenis et al. , 2008Coustenis et al. , 2010;;de Kok et al. 2007;Teanby et al. 2007Teanby et al. , 2008Teanby et al. , 2009a;;Vinatier et al. 2007Vinatier et al. , 2010;;Nixon et al. 2008a;Achterberg et al. 2008Achterberg et al. , 2011)). Here, we explore the thermal and Chemical évolution discemible within the timeffame from 2006 to 2012, thus complementing and refining previous reports from earlier stages of the mission. We study Titan's neutral atmosphère between ~120 and 300 km in altitude (the stratosphère and lower mesosphere) and latitudes from 50°S to 50°N.

CIRS is a FourierTransform Spectrometer, aboard the Cassini orbiter, consisting of Lwo interferometers probing the farinfrared (10-600 cm-1) and mid-infrared (600-1500 cm""1)

ranges with an apodized spectral resolution varying from 15.5 to 0.5 cm-1. CIRS scansTitan's atmosphère throughthree separate focal planes that share the same telescope: FP1 (10-600 cm-1), FP3 (600-1100 cm"1), and FP4 (1100-1500 cm-1) (Flasar et al. 2004). We focus here mainly on FP3 and FP4 high-resolution observations (0.5 cm-1) but also some medium resolution ones (2.5 cm"1).

Prior to 2006 (flybys T0-T9) insufficient data at the right conditions for our purposes (signal-to-noise, émission angle, distance, région, etc.) were acquired. During the flybys that followed, CIRS obtained a large number of spectra in FP3 and FP4 at high, medium, and low spectral resolutions (0.53, 2.54, and 15 cm-1), respectively, in surface-intercepting (nadir)

and horizontal viewing (limb) geometry conditions, albeit not always covering ail latitudes, so that in the northern hemisphere we hâve data only from late 2007 to 2010. An instrument anomaly followed by a reboot of CIRS took place in December 2006, thus depriving us of data from that time. We otherwise use ail available and exploitable data at high spectral resolution (0.5 cm-1) and perform averages of the spectra acquired during one or several flybys as necessary to attain a high signal-to-noise ratio for our calculations.

Tables 1 and2 list the CIRS FP3 nadir northern and Southern averages, respectively, at high resolution (0.5 cm'""1), while Tables 3 and4 list the associated FP4 averages. The average signal-to-noise ratio is shown as well as the corresponding Cassini Titan flyby(s) and the relative solar longitude (Ls). The sélections we made in order to enhance the signal-to-noise ratio cover different latitudes on Titan (we sum ail longitudes because no longitudinal variations were found in our studies; Coustenis et al. 2007Coustenis et al. , 2010)). The averaged spectra created from the highresolution nadir data in the several latitudinal bins (between 50°S and 50°N) contain a large number of spectra in general, with some exceptions (as indicated in Tables 1234). As discussed in Coustenis et al. (2007), we hâve taken care to correct the data for any latitude smearing effect observed for high émission angles and/or data taken at latitudes higher than 50°N. So that the latitudes sounded can be about 5°lower when the line of sight intercepts the surface at 55°N or 55°S, the actual latitude at the vertical of the stratospheric altitudes being observed is then rather 50°N and 50°S, respectively. The latitudes listed in the tables are corrected for this effect.

We hâve applied a radiative transfer code to analyze the CIRS data spanning the 2006-2012 ti mel ine, during whichthe Cassini mission performed 71 Titan flybys (Tables 1234). Besides the high resolution spectra, we also hâve analyzed six data sets at medium resolution (2.54 cm'1) taken in FP3 and FP4 in The Cassini/CIRS spectra are analyzed using a method extensively described in previous articles (Coustenis et al. 2007(Coustenis et al. , 2010) ) . Inbrief, we model Titans thermal infrared spectrum by use of a line-by-line monochromatic radiative transfer code updated from the one we used in previous Titan atmospheric struc ture retrievals (Coustenis et al. 2010, and référencés therein). In We first retrieve température profiles by the inversion of the observed v4 metliane émission band at 1305 cm-1 in CIRS FP4 following the method described in Achterberg et al. (2008Achterberg et al. ( , 2011) ) . We use the CH4 vertical mixing ratio profile in the stratosphère as measured by the Huygens probe (Niemann et al. 2010) and compatible with the CIRS inferences from FPI (Flasar et al. 2005), which yields 1.48% above the cold trap.

The initial guess température profile for these calculations was the 15°S profile from Flasar et al. (2005), in a process described in detail in Achterberg et al. (2008). We also made tests using the HASI measured température structure in the stratosphère and troposphère (Fulchignoni et al. 2005), as discussed in the next section. We then use the température profiles to solve the radiative transfer équation in the FP3 part of the spectrum for the mixing ratios of the various coinponents seen in émission at different epochs and latitudes (Tables 1 and2).

For ail the molécules analyzed here in "nadir" geoinetry con ditions, we adopt constant-with-height profiles above the con densation level relevant to stratospheric levels of 0.1-20 mbar essentially (roughly 80-280 km). As explained in detail in Section 4.4 of Coustenis et al (2010), the use of constant, vertical profiles does not allow for a good fit of the whole C2H2

émission band centered at 729 cm-1. We therefore first dérivé the C2H2 abundance tliat fits the Q-branch and then adjust tire fit in the wings of the band to infer the other mixing ratios of the weaker species in the 620-780 cm-1 région. In order to model the continuum observed in the spectra, we hâve used the haze dependence as reported in Vinatier et al. (2012a) giving the aérosol refractive index from CIRS spectra in the far-and mid-infrared régions near 15ÙS and 20°S, adjusted to match the level of radiance observed in between the molecular bands in our sélections.

TEMPERATURE EVOLUTION OF TITAN'S STRATOSPHERE DURING THE CASSINI MISSION

Figure 1 shows an example of fits obtained in tire FP3

(600-1000 cm-1) région and in the FP4 (1100-1500 cm-1)

région. The latter is used for ail spectral sélections to dérivé the corresponding température profile. constrained inversion algoritlim (Achterberg et al. 2008) as the a priori reference profile the HASI température structure (Fulchignoni et al. 2005). We lind that at the équatorial latitudes ail molecular abundances are affected by 10% or less, except.

for HCN (20%). 'l'Iris small impact is expected because the HASI profile was inferred above the Huygens landing site, which was at 10°S. In the northem latitudes, the divergence is more pronounced (as much as 20% of decrease in abundance)

for some molécules (the ones with the stronger émission bands: C2H2 and HCN), however, even at 50°N, ail the other molecular abundances vary by less than 10% whatever the initial température profile.

In the northem hemisphere, we find a decrease in température in the mesosphere (more sensitive to the seasonal insolation variations), starting from the stratopause above 0.1 mbar, where the température significantly decreases from the earlier to tire later mission dates. The warm north polar lower mesosphere lias cooled (as also suggested previously by Achterberg et al. 2011) and now we find the decrease to be about 12 K at 50°N

(Figure 2 . r* '\ 1 between 1 and 10 mbar, but this is within error bars. More importantly, the general shape of tlie thermal structure is changing and we find the température gradient to be weakening, resulting in a more isothermal profile at liigher latitudes witli a loss of tlie marked stratopause présent in earlicr years.

• 4 : 30°N o -* "i-rf- r-f ' - 1
These changes at northem latitudes suggest a weakening of the descending branch of the middle atmosphère méridional circulation (Achterberg et al. 2011 ), which implies less adiabatic heating and hence lower températures. Moreover, the changes in the circulation affect the distribution of aérosols (Vinatier et al. 2012a), wliich in turn affect the thermal structure. Since the model parameters for ail these calculations are the same (in terms of calibration, metliane profile, haze description model, etc.), we only consider tlie relative uncertainties, due to noise (rather small in our usually large averages, except for the early mission dates and sometimes in the north), the température profiles and any remaining uncertainties regarding some of the spectroscopic data (ail within 15%). Exceptions are CjHs and C2H4, which are affected by abundance retrieval difficulties and some insufficient laboratory spectroscopic data. In the case of C3H$, at 748 cm"1, a major hindrance is tlie rnixing with the C2H2 lines while tlie spectrum of C2H4, around 950 cm"1, suffers from some electrical interférence artifacts ("spikes") and also interférence from QsH» contribution and other émission bands. As a conséquence, these two molécules show ratios that hâve error bars on the order of ±30% and are not as reliable as the other inferences for tlie temporal or spatial variations.

At Southern latitudes (50°S and 30°S, Figure 3) as well as near the equator, we find the abundances of the gaseous species to remain rather constant during the Cassini mission within the error bars. In general (with the exception of C2H4 and C3H8, identified as complicated molécules to process, in particular at the mid and soutliern latitudes), the équatorial data yield abundances higher by about 20% than the 30 or 5Û°S inferences (which are almost at the same level for ail molécules). The mixing ratios at 30°N are in general higher than the équatorial values by lü%-20%. We found a definitive trend for increased gaseous content in the stratosphère when moving from the south to the north, and, in addition, significantly enhanced abundances during the whole mission duration at northern latitudes, observed in most cases at latitudes higher than 30°N.

At 50°N, ail abundances except for CGj are thus enhanced with respect to 55°S and the equator, consistent with our previous work with this model (Coustenis et al. 2007(Coustenis et al. , 2010) ) and with Voyager 1 results (Coustenis et al. 1991;Coustenis & Bézard 1995), almost at the same season as in mid-2010 (A. Coustenis et al., in préparation). In average, and at the maximum, the 50°N values are higher than the 30°N by 20%

and by 40% with respect to the equator. These enhancements appear to be greatest for the shortest-lived Chemicals. This is explained by a combination of chemistry and dynamics:

the vertical distributions are steepest for the shortest-lifetime species in a purely Chemical model, while the presence of a circulation cell in the atmosphère, with subsidence in the polar régions (at the north pôle during the period of the observations), causes the lower stratosphère to be greatly enriched in these species at the winter pôle compared to the equator. The study of the temporal évolution of this enhancement in the higher northern régions within the Cassini mission duration, as well as changes observed with time at southem polar régions, is described by Teanby et al. (2008Teanby et al. ( , 2010Teanby et al. ( , 2012)).

Furthermore, at 50°N, we find an indication for an increase Ethane shows variations with time only at this higher latitude.

We note that C3H8 seeras to hâve increased abundances around NSE, but the data analysis results for this molécule at lower We inferred the l2C/13C ratio in CH4 and QH2 to be 95±15 and the 14N/15N ratio in HCN to be 50±10 for the latitudes on which we investigâte here as an average. The 12C/13C ratio is consistent with results reported by Vinatier et al. (2007) and Nixon et al. (2008a[START_REF] Nixon | Isotopic Ratios in Titan's Atmosphère from Cassini CIRS Limb Sounding: C02 at Low and Midlatitudes[END_REF] and with the terrestrial inorganic standard value (88.9; Fegley 1995). The 14N/1SN ratio is also Since the C2H2 émission band at 729 cm"1 contains infor mation at several altitudes through the central and wing fines, we hâve also tested a vertical profile predicted for Ls = 358°b y the GCM models (Rannou et al. 2005), which does not match the shape of the C2H2 band, but is not too off with respect to the actual values.

At higher northern or southem latitudes, poleward of 70°N, the situationhas dramatically changed in recentyears with much stronger Chemical and température variations [START_REF] Vinatier | Optical constants of Titan's stratospheric aérosols in the 70-1500 cm'1 spectral range constrained by Cassini/CIRS observations[END_REF]Teanby et al. 2012) but at the latitudes considered here, we do not observe them which gives us hints as to the location and the extent of the winter polar vortex. (West et al. 2011) suggests that the dynamics during this period go through a rapid transition which should also affect the gas distribution. The rapid decrease after mid-2009, for which the most straightforward expia nation is that the vortex lias shrunk somewhat. would be consistent with the weakening thermal gradient we find here and tirât of the winds also reported by Achterberg et al. (2008Achterberg et al. ( , 2011) ) and Teanby étal. (2009b). The finding also ties into the location of tire maximum température gradient, which appears to be moving northward over the winter/spring season (Teanby et al. 2010; Figure 3, T panel). If 50°N is emerging from the vortex core, it would cause a large réduction in the abundances, hence explaining our observations. Thus, decreasing abundances at 50°N could be 'I'he Astrophysicâl Jouknal,760;144 (8pp), 2012 December 1 due to a weakening vortex with reduced latéral mixing across the vortex boundary (Teanby et al. 2010).

Another cause could be the spatial variations in the energy input to Titan's atmosphère (due to Titan's inclination) as a driver for changes in the advection patterns, which in tum (2006, 2007, 2008a,b, 2009), Coustenis et al. (2007, 2008), Vinatier et al. (2007, 2009), Nixon et al. (2008a[START_REF] Swinyard | The space infrared telescope for cosmology and astrophysics: SPICA A joint mission between JAXA and ESA[END_REF]), and Jennings et al. (2008) among others.

In a previous paper (Coustenis et al., 2007) Besides these well-known trace species, a firm détection of benzene (C6He). first reported in ISO data (Coustenis et al., 2003), was provided by CIRS at 674 cm-1. We use this band here to study its latitudinal variations. The détection of G,MD was also obtained from the analysis of the extra émission observed at 678 cm-1 (Cou stenis et al., 2003(Cou stenis et al., , 2008)). The D/H ratio on Titan was determined from the CFI3D band at 8.6 pm and found to be 1.17^® x 10~4 (Coustenis et al., 2007). [START_REF] Bézard | Détection of "CHjD on Titan[END_REF] furthermore reported the détection of ,3CH3D at 1148 cnf ' and determined a D/H ratio of 1,3 2/^;]3 x 10~4. Nixon et al. (2008a,b) and [START_REF] Jennings | Isotopic Ratios in Titan's Atmosphère from Cassini CIRS Limb Sounding: HC3N in the North[END_REF] used CIRS data to retrieve information on the l2C/,3C and other isotopic ratios in Titan. Finally, there is ongoing work on the analysis of CIRS limb data to infer information on the vertical profiles of several gaseous constituents of Titan's stratosphère at nine lati tudes between 56"S and 80°N (Vinatier et al., 2007[START_REF] Vinatier | Analysis of Cassini/CIRS limb spectra of Titan acquircd during the nominal mission. 1: Hydrocarbons, nitrilcs, and C02 vertical mixing ratio profiles[END_REF]Teanby et al., 2007Teanby et al., , 2008a,b),b).

The results published in Coustenis et al. (2007) andTeanby et al. (2008a) indicated that no longitudinal variations existed for any of the gases. On the other hand, the retrievals of the méridional variations of the trace constituées showed an enhancement for some of them towards the North pôle. However, due to the geometry of the observations from 2004 to 2006, the most northern latitude exploited was limited to about 60°N. Mol écules showing a significant enhancement at northern latitudes are the nitriles (HC3N, HCN) and the complex hydrocarbons (C4H2, C3H4). The results were tied to prédictions by dynamicalphotochemical models [START_REF] Hourdin | The GE1SA spectroscopic database: Current and future archive for Earth's planetary atmosphère studies[END_REF]lawas et al., 2008a,b;Crespin et al., 2008).

With respect to the results published in Coustenis et al. (2007),

we improve here on the analysis by exploiting a far larger number of spectra with a signifïcantly higher number taken at the higher resolution of 0.53 cm-1. The more complété coverage of Titan's disk, combined with the larger number of spectra at high resolu tion available, allows for the inference of more précisé abundances for the trace gases and for a more adéquate définition of méridional variations, in particular in the northern régions, which were not fully observed previously.

Nadir observations in the FP3 and FP4 GRS ranges

For a full description of the CIRS spectroscopic observations and the instrument capabilities we refer the reader to Flasar et al.

( 2004) and for the general procedure applied to the data in our studies, see Coustenis et al. (2007).

The GRS nadir spectra characterize various régions on Titan from roughly 85°S to 80eN with a variety of émission angles.

We hâve studied the émission observed in the CIRS FP3 and FP4

detector arrays (covering roughly the 600-1500 cm"1 spectral range) with apodized FWHM (full width to half maximum) reso lutions of 2.54 or 0.53 cm"'. The more recent Cassini/CIRS obser vations (T10-T44, until May 2008) hâve allowed for a more complété coverage of the higher northern latitudes on Titan, where previously only a few spectra were available (Teanby et al., 2006;Coustenis et al., 2007). We include here a large num ber of both mid-and high-resolution CIRS spectra at ail latitudes from 90°N to 90°S at émission angles smaller than 50°and we perform averages in order to focus on a given latitudinal range and also in order to enhance the S/N ratio (see Table 1). The airmass numbers and the S/N ratios given in Table 1 represent aver ages of the mean values in each spectral average. As said in the previous section, the lack of longitudinal variations in gas abundance that we hâve demonstrated in Coustenis et al. (2007) justi fies the latitudinal binning, while longitudinal variations in température [START_REF] Nixon | Isotopic Ratios in Titan's Atmosphère from Cassini CIRS Limb Sounding: C02 at Low and Midlatitudes[END_REF] are small enough to be ignored for our purposes.

The composite spectral averages thus produced in the 600-1500 cm"1 range covered by the FP3 and FP4 detectors (Fig. 1)

show several signatures of previously identified molécules: hydro carbons (CH", CH3D, C2H2, C2H", C2Hs, C3H4, C3H8, C.,H2i C*Hfi), ni triles (HCN. HC3N) and C02. We hâve included ail of these species in our model, but we hâve not included ail the isotopes of C, N and O discovered today for lack of complété spectroscopic data. In addition, C2HD, as observed at 678 cm-1 (Coustenis et al., 2008) is incorporated in our model but the associated abundance and its relative variations were extensively discussed in the earlier paper.

Small offsets in the wavenumber of the spectra were corrected by cross-correlation with a synthetic spectrum. These instrumental offsets were generally small compared to the resolution and were at most ±0.1 cm"1. However, these small adjustments reduced the misfit between model and data by up to 25%. Characteristics of the spectral sélections from Titan flybys TB-T44 used in this study for medium spectral resolution of 2.5 cm 1 and high spectral resolution, 0.5 cm '.

The latitudes in the first column refer to atmospheric levels at around 5 mbar and take into account the latitude smearing effect (see text). The latitude coverage for each sélection isactually around the mean value (for example 5"S means spectra from 10"S to theequator). We sum ail spectra from ail longitudes in columns 2 and 5. The émission angles are restricted to 50" everywhere and the retrieved average airmass is given in columns 3 and 6. The S/N indicated in columns 4 and 7 is the average signalto-noise ratio for ail spectra combined. It varies with wavelength across the spectrum depending on the CIKS data NESR (noise équivalent spectral radiance). Our analysis here is based essentially on the high-resolution data acquired by CIRS, with the medium-resolution spectra used as a check for consistency.

Modeling

Our method in deriving abundances of the trace constituents on Titan has been largely described in previous articles (see Coustenis et. al., 2007a and references therein). In brief, we use a line-by-line radiative transfer code to simulate the Titan spectrum and through itérative processes we dérivé the abundances of the various con stituents observed in Titan's atmosphère. We begin with tempéra ture profiles obtained by inversion of the émission observed in the V4 methane band at 1305 cm"1, assuming 1.4% of CH4 in the strato sphère as measured by tlie Huygens probe and compatible with the CIRS inferences from FP1 (Flasar et al., 2005;Niemann et al., 2005).

We use these profiles to then solve the radiative transfer équation in the FP3 part of the spectrum for the mixing ratios of the various components.

Our model includes collision-induced absorption and a twocloud System to represent the haze opacity (see Coustenis 4x10"7 700 800

Wavenumber (cm'1)

Fig. 1. Spectral averages performed over TB-T44 Titan fiybys as described in Table 1 for different latitudes binned in 10°: (a) at medium resolution ( 2 Even though ail the neutral species are contained in our model, we lack spectroscopic data and hâve not included several of the isotopic species that hâve been detected to date on Titan, nor ail the propane bands (Nixon et al., 2009).

For ail the molécules analyzed here, with the exception of C2H2, C2H6 and HCN, we adopt only constant-with-height profiles above the condensation level, because there are not enough strong Unes in their émission bands to allow for the retrieval of alritude-dependent information. Only C2H2 -with a resolved band at 730 cm"1 that probes multiple altitudes via strong and weak (optically thick and thin) Unes, aliows for an unambiguous discrimination between different vertical distributions. However, the C2H6 and HCN verti cal distributions extracted from GCM models or from limb calcula tions and tested here also produce a slight improvement of the spectral fit with respect to the constant-with-height profiles, as will be discussed hereafter.

Analysis of the data

We use the information contained in the v4 methane band at 1305 cm"1 to retrieve température profiles, which are then injected in the radiative transfer équation for the other molécules to solve for the gaseous opacifies. By comparison with the observa tions a best fit is obtained through an itérative process. For each sélection (see Table 1) the continuum level between émission sig natures is matched by varying the haze/cloud opacity. With the exception of C2H2, and in some cases -but not systematicallyof HCN and C2H6, the vertical distributions of the trace constituents were assumed to be constant-with-height above the condensation level and set to follow their saturation vapor pressure curve below.

î. The FP4 région

Température profiles were retrieved from zonally-averaged spectra over ail of the température maps up through T40, for a methane abundance of 1.4%. The spectra were averaged in 10°lat itude bins, based on the latitude of surface intercept, with the émission angle restricted to <50". For the procedure of this retrie val of Titan's stratospheric températures by CIRS see Achterberg et al. (2008a). The analyzed nadir mapping data presented in that paper sampled latitudes between 90°S and 60°N, providing tem pératures for pressure levels roughly in the 0.2-5-mbar range (in some cases, the contribution functions can reach 10 mbar). The lat itude range for the température profiles has been extended in this work to 90°N and they are inferred using the Huygens Atmospheric Structure (HAS1) profile (Fulchignoni et al., 2005) as the initial guess, so that below the altitudes where CIRS can retrieve tempér ature information, the profiles are set to the HASI data. Some of the température profiles retrieved from the inversion of the émission observed in the methane v4 band at 1305 cm"1 are shown in Fig. 2.

Using these température profiles, we then performed radiative transfer simulations of Titan's spectrum starting with the FP4 (1100-1500 cm"1) spectral région where the émission from 1200 to 1400 cm"1 is mainly due to methane opacity. However, CH3D, C3H8, C2H6 and other molécules contribute at shorter and longer wavelengths. CH3D has a strong émission band at 1156 cm"1, diag nostic of régions around 0.3 mbar, Fig. .3. The fit in this band gives us access to the D/H ratio in Titan, as discussed extensively in Coustenis et al. (2007Coustenis et al. ( , 2008)).

The best fits at 15°S in the FP4 région are shown in Fig. 4 with residuals for other latitudes. The v4 CH4 band yields information at different altitude levels as testified by the contribu tion functions of the Q-branch and its wings (Fig. 3). The 3 -o standard déviation due to noise atone in the FP4 région is about 2 k 10"10W cm"2 sr'/cm"1. When other error sources are in-Temperoture(K) Fig. 2. Température profiles at different latitudes for medium résolution. From 70"S to about 30"N, the profiles are very similar and close in shape. In the troposphère, the profiles are set to match the HASI température data, while the HASI profile is also used as an initial guess in the stratosphère for our température retrievals. the total uncertainty is on the order of 2 x 10"9Wcm"2sr"'/ cm"1. Some of the remaining misfits in this région can be explained by the presence of molecular bands or isotopes not included in our model.

Thus, the misfit of the spectrum between 1370 and 1500 cm"1 could be due to an uncertain représentation of the C2H6 bands, the Vg at 1376 and the v8 and 1468 cm"1, as well as to the several predicted propane bands (v18 at 1378, v7 at 1464 and V24 at 1472 are expected), and spectral aliasing beyond 1430 cm"1 due to numerical filtering in the instrument (see Nixon et al., 2009). Also note the weak émission of 13CH3D at 1148 cm"1. The cause for the returning discrepancy between the model and the observed radi ance between 1180 and 1250 cm"1 has not been dearly identified, but can be due to uncertainties in the CH4 spectroscopic data.

Trace gas abundances from the FP3 région

The 600-1000 cm"1 spectral région in Titan's spectrum contains the signature of many hydrocarbons, two nitriles and an oxygen compound (CjH,, C2HD, C2H4. C2HB, C3H4, C3H8, C4H2, C6H6. HCN, 1100 1200 1300 1400 1500

Wavenumber (cm'1) Wavenumber (cm ') Fig. 4. The best fit to the FP4 région induding the CH, band at 1305 cm 1 and thc Cl li ü band at 1156 cm V The methane abundance in the stratosphère is laken to be 1.4%.

: b -» _ I Model - || _ CI RS FP4 15°S - 1_ Data-Model 50°N 1 1 ' ' ' ' 1 ' ' * » 1 1 ' ' 1 1 1 * 1 ;4x À , * . -= 2x
Upper panel: medium resolution; lower panel: high résolution. When ail error sources are inctuded for this région, the total ur.certainty is on the order of 2 x 10 'Wcni 2sr '/cm '. See texc for more details.

HC3N and CO2). Ail of these molécules are included in our mode).

Using our radiative transfer code, we perform simulations of the 600-1000 cm"' région and compare them to the spectral averages described in Table 1 for ail the different latitudina) bins. The con tribution functions of a sélection of molécules for two latitudes (5°S and 50°N) are shown in Fig. 5.

Fig. 6 shows examples of fits in the FP3 medium and high-resolution data as a whole for five different latitudes. The fits are generally quite good, within the 3 -a standard déviation error, except for a few localized areas which can be associated in tlieir majority either to insufficient spectroscopic data (in particular for C3H8 or isotopes, at frequencies identified hereafter) or to noise interfér ence features in the Titan spectrum. already known to exist on Titan, C4H* C3H4, HC3N, C02 and C6H6

hâve their band centers at 628,633,663,667 and 674 cm-1 respectively. Ail of these molécules hâve been included in our gaseous mode! (Coustenis et al., 2007). In addition, C2HD was identified at 678 cm'1 and included in this model (Coustenis et al., 2008).

Furthermore, the left wing of the C2H2 band, centered at 728 cm-1, significantly contributes here, as does, to a lesser amount the R-branch of the HCN band (centered at 713 cm-1). In order to simulate this région, we first obtain the best possible fit for the C2H2 and HCN bands (as explained in the following section). Fig. 7 shows the fits and the residuals in this région (610-690 cm"1). The 3 -a standard déviation due to noise alone in the FP3 région for spectra taken at mid-latitudes is about 1.2 x 10'10Wcm'3sr1/cnr' for the high-resolution data (negligible with respect to other sources of errors). This is also the case for the medium resolution data. When other error sources are in cluded (see Section 4.6) the uncertainty Ievel (at 3-cr) on the émis sion observed increases and is given on the figure. Sonie of the remaining misfits in this région can be explained by the presence of molecular bands or isotopes not included in our model, such as the isotopes of C02, which hâve been identified at 648.5 and 662.5 cm-1 [START_REF] Nixon | Isotopic Ratios in Titan's Atmosphère from Cassini CIRS Limb Sounding: C02 at Low and Midlatitudes[END_REF], as well as those of HC3N at 658.7 and 663.0 cm-1 [START_REF] Jennings | Isotopic Ratios in Titan's Atmosphère from Cassini CIRS Limb Sounding: HC3N in the North[END_REF].

C,H2 and HCN

We hâve tried to simulate as well as possible the émission ob served in the whole C2H2 band between 680 and 760 cm-1. We hâve used both constant-with-height mixing ratios and vertical distributions inferred from the study of CIRS limb data or provided by the General Circulation Model (hereafter GCM) for Titan of Rannou et al. (2005).

Note that there are several systematic instrumental interfér ence features in the FP3 spectrum thaï hâve been identified. One such noise feature is near the center of the C2H2 Q-branch, at 729.25 cm-1, and could be the cause of the imperfect fit. Also, another strong interférence affects the spectrum near 765 cm-1. If any event the simulations obtained when using the constant verti cal profiles (Fig. 8) do not allow for a good fit of the whole C2H2 band (Fig. 9a-c). When the center of the band, at 729 cm-1, is fitted, the simulated émission in the wings is in excess of the obser vations, especially in the left wing (R-branch). From the contribution functions of C21I2 shown in Fig. 8, it is clear that these régions probe different atmospheric levels and hence require a ver tical distribution with a slope which gives lower abundances at lower altitudes, like the ones inferred from limb CIRS data and from the GCM simulations. For the purposes of our fit then when using constant profiles, and in order to dérivé the abundances of the weaker-band species (like HC3N, C02, C2HD and C6Hc) embedded in that part of the spectrum, we need to reduce the C2H2 abundance until we fit the left wing. The correspondingC2H2 abundance is significantly different from the one that fits the center of the Qbranch (by up to 80Ï).

Therefore, after having obtained the best possible fits with the constant-with-height profiles for ail latitudes, we then turned to testing vertical distributions for acetylene, HCN and C2H6. We obtained vertical distributions for these three constiruents from the GCM calculations (Crespin et al., 2008) for Ls = 300. In addition, for C2H2 we tested the vertical profiles produced by CIRS limb inferences (Vinatier et al., 2007).

Fig. 8 then also shows the C2H2 GCM vertical distributions tested in this work for latitudes of 50°S, 15 "S and 50°N. The distri butions for the mid and Southern latitudes are similar in shape and in mixing ratio (Fig. 8) in the 0.5-20 mbar range probed by the C2H2 band center and wings as shown by the contribution func tions (Fig. 8). On the contrary, the 50°N GCM-inferred distribution is quite different, both in shape and in abundance.

The GCM distributions that we tested worked well in satisfying both the Q-branch and the wings of the acetylene band for C2H2 abundance values close to the ones inferred by fitting the Q-branch ; 2x10'

-?xl( 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1c r B£ _ Model CIRS FP3 6OS - _ nt 60s _ -g T -,-, - Fil 50S _ - .
If-"-y- Wavenumber (cm"1) Fig. 6. Best global fits to the FP3 région at medium (a) and high (b) résolutions, using everywhere constant-with-height molecular profiles.

with constant-with-height profiles at inid-latitudes. They also work well for the shape reproduction at higher northern and Southern latitudes, but they need in eacli case to be adjusted by significant multiplying factors (about 0.3 for the North and 0.5 for the South). This means that these distributions need to be significantly reduced in abundance to match the data at high northern and Southern latitudes. But in general, their use allows for a better fit to the data than with constant-with-height profiles. This is the case also for the vertical profiles inferred ffom CIRS limb data (Vinatier et al., 2007), so that the vertical distributions always pro vide a better fit to the whole acetylene band simultaneously. The abundances retrieved for the minor species embedded within the C2Hj. wings (and in particular HCN and C3H8) are the same, since we always first adjust the C2H2 mixing ratio to match the observed émission in the left and right wings of the acetylene band before further exploring the abundances of other species.

The vertical profiles used for HCN hâve allowed for a slightly improved fit to the band at 713 cm'1. We can thus say that the na dir data hold proof that a vertical distribution with a marked slope (as inferred in the models and by the CIRS limb data analysis ( Teanby et al., 2008a,b;Vinatier et al., 2007[START_REF] Vinatier | Analysis of Cassini/CIRS limb spectra of Titan acquircd during the nominal mission. 1: Hydrocarbons, nitrilcs, and C02 vertical mixing ratio profiles[END_REF]) is required. But no further information can be inferred as a function of altitude from the nadir spectra.

The ethane and ethylene abundances

Ethane and ethylene dominate the émission observed in the 750-1000 cm"1 spectral région, with band centers at 821 and 949 cm"' for the v9 and v7 bands respectively (Fig. 10). The ethyl ene retrieval is difficult due to the strong interférence from a noise feature at 956 cm-1.

New spectroscopic data was made available for C2H6 from Vander Auwera et al. (2007),which greatly improves the fit in the C2H6 band ar 821 cm"' and allows us to include the v6 and v8 bands in the 1350-1500 cm"1 région. The région also contains weak pro pane bands at 869 and 922 cm-1 (Nixon et al., 2009). Also a possi ble gradient in the haze simulation could be affecting the continuum.

Contribution functions ot 50°N Contribution functions at 50°S

Fig. 8. Upper panel: vertical distributions for C3H2 testcd here: the GCM model (Crespin et al, 2008) and che limb-inferred profiles (Vinatier et al., 2007) are shown, along with the constant-with-height best-fit distributions. In order to fit the center and the wings o! the 729 cm 1 band in the high-resolution data at 50"N, 15°S and 50°S the vertical distributions need to be adjusted by certain multiplying factors. Also shown (lower panel) are the contribution functions as calculated with the GCM profiles for 50°N and 5CriS; the latrer is also représentative of équatorial and mid-latitudes.

The new fit obtained in the ethane band now yields much smaller abondances (by 45-60%) with respect to what was previonsly found (Coustenis et aL, 2007) due to the higher band strength (by 1.44 with respect to the previous GEISA version). When a ver tical distribution is used, we get a slightly better match to the data, but within uncertainties, we cannot distinguish between the con stant-with-height and the vertical profiles.

The uncertainties

The uncertainties on the abundances inferred here are calcu lated as explained in Coustenis et al. (2007), Section 4.1. We take into account « the instrument noise on the location of the émission peak (negligible with respect to other sources of error given the large number of spectra we generally sum up in a sample, with the exception of higher northern latitudes for medium resolution),

• the uncertainty on the position of the continuum level above which the émission rises. and the uncertainty on the température profile and its impact on the line formation région.

In addition, systematic errors due to the uncertainty on the CH., abundance, on the spectroscopic data, calibration problems and inconsistencies in the samples or the datasets used, were also taken into account when the absolute error bars are given.

The précision on the température profile inferences described in Section 4.1 is discussed in Achterberg et al. (2008a) and in Couste nis et al. (2007). The uncertainty from the instrument noise is small (0.1 K or so) so the main errors are systematic, mainly originating in the uncertainty of the CR, abundance: about 1 K around 1 mbar, increasing to ~4 K about 5 mbar and from the sensitivity to the ini tial guess (~2 K or so).

Due to the large number of spectra considered here, as well as to the improved spectroscopic parameters for C2H6. CjH,. HC:jN Fig. 10. Fit» the 750-1000cm ' région with theémissions mainly of ethane and ethylene, with constant in altitude profiles, for high-resolution spectra and at different latitudes. The fit » the ethane band with new spectroscopic parameters has significantly improved the fit » the émission in the band, but there is still room for improvement in the 800-840 cm 1 région, although a vertical distribution for C2Hs slightly improves the fit. The fit to the ethylene band is also shown. There is a strong bad pixel feature interfering near 957 cm \ 3 -a error bars are indicated.

) v «v»ï* y v v h v yj » V v v V w » v y' V v jj Doto-Model with constont C2Hi (3.1 • -06) V V v v w Ÿ v V v V ' y v V ï v V « i v Y V v J
and some isotopes, we hâve been able to considerably reduce the error bars on our inferences with respect to previous papers.

Méridional variations of the trace gaseous constituents

We hâve inferred the méridional variations for the different constituents présent in the CIRS spectra from 600 to 1500 cm-1.

They are shown in Fig. lia andb, along with the 3 -a error bars including ali uncertainties. For some latitudes, the mixing ratios in ferred from high-resolution data, are given in Table 2. Benzene and HC3N are difficult to detect at low Southern latitudes. However, while the HC3N émission is clearly observed, even at low latitudes, C6H6 éludés firm détection. Thus, for latitudes <40°S in Fig. 11 b, the values for C6H6 are upper limits.

As in Coustenis et al. (2007), we find thar the observed enrichment reported here for some species at high northern latitudes is less at this epoch (mid-winter at the northern pôle) than at the time of the Voyager encounter just after spring equinox (Coustenis and Bézard, 1995). The gases showing a strong enhancement in the North are the complex hydrocarbons (C3H4, C4H2, and C6H6) and the nitriles (HCN and HC3N). For most of these gases, the increase strongly manifests itself starting at around 40"N, with the excep tion of HCN which exhibits a regular increase from South to North.

Other constituents with a smaller increase in abundance (also from about 40°N and northwards) are: C2H2, C2HD (whose détection and abundance results are given in Coustenis et al. (2008)) and C2H4.

Propane (C3HS) is now found to be constant or slightly increasing the equator in the lower stratosphère. Our results here are in agreement with these inferences, although the altitudes of the limb data are not the same as the ones probed by the nadir data analyzed here. Furthermore, the abundance minimum in our case cannot be confïrmed given the error bars and there is nowhere an indication of anything similar for 50CS.

Interprétation of the results and discussion

Our findings on the trace gas abundances (Fig. Tl and Table 2)

reported here mostly support the results in the previous paper (Coustenis et aL, 2007) with the advantage that our new uncertain ties are smaller and we hâve secured more reliable spectroscopic data for some molécules (HCN, HC3N, C2H6), as well as information on new species (C2HD, see Coustenis et al., 2008). This allows us to improve, in particular, the inferences at high North and South lati tudes. Indeed, différences in Titan's composition between this pa per (Table 2) and the previous one pertain mainly to the ethane abundance and to mixing ratios found at higher latitudes for some species (such as C2H4 and C3H8 for instance, compare We also compared the findings in this paper with gas abundances in Titan's stratosphère reported in other papers by the CIRS team members. We are in good agreement with isotopic ratios and values reported by [START_REF] Bézard | Détection of "CHjD on Titan[END_REF] and Nixon et al. (2008a,b).

We hâve tested the vertical distributions for C2H2, C2116 and I1CN as recovered from limb data analyses by Vinatier et al. (2007[START_REF] Vinatier | Analysis of Cassini/CIRS limb spectra of Titan acquircd during the nominal mission. 1: Hydrocarbons, nitrilcs, and C02 vertical mixing ratio profiles[END_REF] and we find them to be compatible with our nadir data for nine dif ferent latitudes between 50°S and 80^. The abundances inferred from the limb data for other molécules as well are plotted in Figs. 12a and12b. These values are compatible with ours at mid and low latitudes and slighrly higher towards the North (as was also found in Coustenis et al. (2007)).

Hereafter, we compare our results with theoretical work and current models for Titan's atmosphère. Attempts to model the seasonally varying distribution of organic compounds in the middle atmosphère hâve used 2-D simulations with parameterizations of eddy and wave momentum flux convergences.

The Titan model proposed by Lawas et al. (2008a,b) is a 1-D simulation coupling between photochemistry, radiation transfer and aérosol microphysics. The production of ail main neutral gas species is followed from the initial phorolysis of N2 and CFLi, while the production of aérosols is described through spécifie Chemical Table 2 Mixing ratios for the Titan trace gaseous constituents in the stratosphère for a selected set of latitudes from nadir high-resolution data using constant-with-height distributions. Vinatier et al., 2007;red fines, 30°N, Vinatier étal, 2009) and the model simulations (black curves). The points at 1100 km in altitude are theINMS results (Waiie et al., 2007;[START_REF] Vuitton | Formation and distribution of benzene on Titan[END_REF]Cui et al., 2009). In the green boxes, the horizontal dimension is the uncertainty on the mixing ratio, while the vertical side gives the extent of the contribution function at FWHM. The model correctiy reproduces the results from nadir and limb C1RS data for these molécules within error bars. See text and Vinatier et al.

(2009) for more details. One set of these retrievals is our own work at 33°N, the others are the vertical profiles by Vinatier et al. (2007) at 15°S and at 30°N [START_REF] Vinatier | Analysis of Cassini/CIRS limb spectra of Titan acquircd during the nominal mission. 1: Hydrocarbons, nitrilcs, and C02 vertical mixing ratio profiles[END_REF]. In addition, the retrieved abundances in the thermosphère by 1NMS [START_REF] Waite | Ion Neutral Mass Spectrometer Results from the First Flyby of Titan[END_REF]Vuitton et al., 2007[START_REF] Vuitton | Formation and distribution of benzene on Titan[END_REF]Cui et al., 2009) are compared. Finally, water vapor was detected by ISO (Coustenis et al., 1998) and the abundance re trieved is given in Fig. 12c. Water vapor was also detected by CIRS and preliminary investigations show that the ISO abundance fits the data [START_REF] Nixon | Water in Titan's stratosphère from Cassini CIRS observations[END_REF].

Mixing Ratio

There is a general good agreement between tlie observations and the model although there are spécifie cases for which the agreement is less good. Tlie main hydrocarbon and nitrile species are well described by the model with the stratospheric abundances being in good agreement with our retrieved ones, except for the cases of C2H4 and HC3N. For the former the simulated profile decreases further than the stratospheric value towards the surface, although the observations suggest a reversai in the mixing ratio profile, while for latter the simulated abundance is larger than the observed one at the spécifie latitude. These apparent discrepancies could be a manifestation of advection processes not captured by the 1-1) character of the simulation. Another hint suggesting the presence of other processes taking place in the atmosphère is the case of C3H8 for which although the strato spheric abundance is in agreement with the CIRS nadir-retrieved value, the vertical profile from the limb observations suggests a different behavior than the simulated one. Furthermore, the HCN and C2H2 profiles at 3CTN appear to decrease above 250 km relative to the équatorial profile that présent a more monotonie altitude behavior, thus suggesting the contribution of advection phenomena. It is important also to note that the current model includes only the neutral chemistry contribution in the abundances of the simulated species. Although this is a good approach for most of the observed species, the contribution of ion-neutral chemistry reaction can become significantly important for some cases such as that of benzene (see Vuitton et al., 2009). Nonetheless, most of the CIRS-inferences reported here and from the limb data are well simulated by this 1-D mode, which proves to be a good tool to simulate vertical mid-latitude distributions of Titan's trace constituents, including the INMS results, while pointing to some issues with the currently adopted Chemical pathways for Titan's atmo sphère and the contribution of dynamical effects.

The results presented in this work were also compared with prédictions by recent 2-D dynamical models such as the General Circulation Model (GCM) developed for Titan at the 1PSL (Crespin et al., 2008 and references within). These models [START_REF] Lebonnois | Seasonal variations of Titan's atmospheric composition[END_REF][START_REF] Lebonnois | Impact of the seasonal variations of composition on the température field of Titan's stratosphère[END_REF][START_REF] Hourdin | The GE1SA spectroscopic database: Current and future archive for Earth's planetary atmosphère studies[END_REF]Crespin et al., 2008) méridional circulations and the mixing by barotropic eddies. These seasonal models predict that the circulation in the middle atmo sphère is dominated by a pole-to-pole cell with ascent over the summer pôle, followed by a circulation reversai during a short period after the equinox, with a brief transition period in which the ascending branch is at low latitudes and the circulation is more hemispherically symmetric. According to these models, and as described in Flasar et al. (2005), the subsidence over the winter pôle accounts for the enhancement in the concentrations of the organic compounds that are observed (see also [START_REF] Nixon | Isotopic Ratios in Titan's Atmosphère from Cassini CIRS Limb Sounding: C02 at Low and Midlatitudes[END_REF]Teanby et al., , 2009)).

We also find here a generally good agreement with the équato rial and the northern (winter) pôle distributions. In particular, the observed enrichment reported here at high northern latitudes, slightly stronger aller spring equinox (Voyager) than for mid-winter (Cassini) is found in the GCM. The vertical distribution predicted for C3H2 by the GCM for mid-latitudes is also compatible with our observations with minimal adjustment. Note, however, that at higher northern and Southern latitudes, the GCM acetylene profiles require a strong reducing adjustment factor while also ver tical profiles inferred by limb CIRS data are not reproduced by the GCM (Crespin et aL, 2008). An additional disagreement can be found in the model prédictions for the southem hemisphere (currently summer pôle), where the model produces a secondary cell in the low stratosphère, which maintains an enrichment over this pôle and a depleted région at 30°S, both of which features are not observed (see Fig. 10 of Crespin et al. (2008)). However, spatial mapping provided by CIRS, both in the nadir-and limb-viewing modes, shows a structure that is quite complex, particularly near the North (winter) polar vortex, that the 2-D models cannot adequately explain. [START_REF] Nixon | Isotopic Ratios in Titan's Atmosphère from Cassini CIRS Limb Sounding: C02 at Low and Midlatitudes[END_REF] reported that the isoclines of several compounds, including HCN, HC3N, C3H4, and C.,H2 hâve an altitude-latitude structure that is tilted in the middle atmo sphère, and these compounds exhibit a déplétion zone near the vortex in the mesosphere. Clearly, the 2-D parameterizations need to include additional transport processes to account for behavior like this. Understanding the seasonal change of the polar vortexits buildup in the late fait and early winter, its mature phase into early spring, and its subséquent break-up-and the attendant changes in the organic concentrations will be key in elucidating these processes.

In spite of this important breakthrough in modeling Titan's cir culation [START_REF] Nixon | Isotopic Ratios in Titan's Atmosphère from Cassini CIRS Limb Sounding: C02 at Low and Midlatitudes[END_REF], 3-D models may be required before the complexity of Titan's atmospheric composition and dynamics can be satisfactorily reproduced in a modeL try on Titan (e.g. [START_REF] Lara | Vertical distribution of Titan s atmospheric neutral constituents[END_REF]Wilson and Atreya, 2004;[START_REF] Hôrst | Origin of oxygen species in Titan's atmosphère[END_REF]).

Titan's known oxygen compounds to date are carbon monoxide (CO, ~47 ppm), carbon dioxide (C02, ~15 ppb) and water vapor (H20), where the abundance? are quoted for the low-Iatitude stratosphère (de Kok et al, 2007a). C02 was first detected by Voy ager 1 [START_REF] Samuelson | CO2 on Titan[END_REF], while CO was first seen by groundbased observations in the near-IR (Lutz et al, 1983). Subséquent observations in the sub-millimeter led to controversy as to whether CO was wcll-mixed or not (Uidayat et al, 1998;Gurwell, 2004). CO émission fines were later observed by the Cassini Com posite Infrared Spectrometer (CIRS), thus improving the previous abundance estimate ( de Kok et al, 2007a;Teanby et al, 2009).

Water was first detected in Titan's atmosphère by the Infrared Space Observatory (ISO) in 1997 (Coustenis et al, 1998). Two fines near 40-pm observed by the Short Wavelength Spectrometer (SWS) were modeled using a uniform mixing ratio above the con densation level and a value of 0.4 ppb was detected. An early attempt to measure H20 with Cassini CIRS was unsuccessful due to poor signal-to-noise (S/N) ratios in early versions of the calibration pipeline spectra and a limited number of available spectra. There fore, only an upper limit of 0.9 ppb could be retrieved (de Kok et al., 2007a). Since then, water émission in the CIRS data has been definitely observed, albeit without deriving any further information on its abondance and distribution [START_REF] Bjoraker | American Astronomical Society. DPS Meeting #40[END_REF]. The Cassini Ion and Neutral Mass Spectrometer (INMS) detected H20 in the upper atmosphère -between 950 and 1200 km -with a mixing ra tio in the range of ~ (0.4-3.4) x 10"5 (Cui et al., 2009).

While the presence of these oxygen compounds is now wellestablished, some details about their origin remain to be determined. Early photochemical models assumed that CO originated froin episodic outgassing from Titan's interior along with nitrogen (N2) or ammonia (NH3) and methane (CH.,), whereas water molé cules entered the top of the atmosphère and were photochemically dissociated to produce hydroxyl radicals (OH) [START_REF] Wong | Evolution of CO on Titan[END_REF]Wilson and Atreya, 2004). The combination of OH and CO led to the production of C02. However, [START_REF] Hôrst | Origin of oxygen species in Titan's atmosphère[END_REF] hâve recently challenged this model, arguing instead that both CO and C02 are the resuit of upper-atmospheric chemistry that occurs between in-falling oxygen species reacting with carbon produced by CH* photodissociation. In this hypothesis, water enters Titan's atmo sphère either in the form of H20 or OH (since the latter is quickly converted to H20 within the atmosphère) together with oxygen ( 0and O*). These forms of oxygen are thought to be deposited on Ti tan at two different altitudes. O* ions hâve been observed flowing into Titan's atmosphère (Hartle et al.. 2006a,b) and they are thought to be deposited in the upper atmosphère around 1100 km [START_REF] Hôrst | Origin of oxygen species in Titan's atmosphère[END_REF] where their interaction with methyl (CH3) radicals leads to the formation of CO. Water is instead depos ited at 750 km due to micrometeoritic ablation [START_REF] English | Ablation and chemistry of meteoric materials in the atmosphère of Titan[END_REF] where it is photolyzed to OH. The latter finally combines with CO to form C02 and possibly other complex species.

Satum's rings and the icy satellites that surround the giant planets, and also interplanetary dust, are probable sources of the water (oxygen) in Titan's atmosphère and recent results from INMS indicate that the plumes of Enceladus are the dominant source (e.g. Dougherty et al., 2006). Observations and models of the neutral H20, OH, and O toms formed from the Enceladus plume show that material from Enceladus extends well beyond Titan's orbit [START_REF] Melin | The distribution of atomic hydrogen and oxygen in the magnetosphere of Saturn[END_REF]Cassidy andjohnson, 2010;Eleshman et al., 2012). Based considérable increase of the number of data collected by CIRS and significant improvements to their calibratioa The increased signal to noise (S/N) ratio not only permits a definitive détection of H20 from the analysis of CIRS far infrared spectra, but it allows us to constrain its vertical and latitudinal profile.

2. Selected dataset CIRS (Flasar et al., 2004) is comprised of three Focal Planes observingin the spectral range 10-1400 cm"' with spectral resolu tions from 0.5 to 15.5 cm"1. Focal Plane 1 detector (FP1) is characterized by a circular field of view (EOW) of 3.9 mrad. It records data in the far infrared spectral range (10-600 cm"') with a spectral resolution of 0.5 cm"1, allowing us to observe the water vapor sig nature, and by modeling, to retrieve its abundance. Water présents its rotational fines in the CIRS FP1 spectral région up to 400 cm"1, with the strongest and most visible fines in the range positioned between 90 and 260 cm"1. We focus hère on the range from 150 to 260 cm"1 for the water détection, as this is the range of maxi mum responsivity of FP1. At lower wavenumbers the on-board electronics of CIRS create a moving interférence spike that can af fect the spectrum up to 150 cm"1. Therefore, we exclude wave numbers shorter of 150 cm"1. We use data from two different types of observations to obtain independent measurements: the far infrared on-disk intégrations (FIRNADCMP) and the far infrared limb intégrations (FIRLMBINT). Water is a trace species with relatively weak fines and therefore it cannot be observed in an individual spectrum. An average of a few thousand spectra of on-disk observations and a few hundred spectra of limb observations is necessary to achieve sufficient signal-to-noise (Fig. 1).

Limb 2010), therefore we model this average using a température profile retrieved for 15°S. We exclude from this analysis the higher northern latitudes where the stratospheric température profile changes significantly. We hâve also se lected data acquired by the spacecraft at a range less than 45,000 km in order to limit the size of the projected detector foot print on the limb to less than 150 km. iceof [START_REF] Murphy | Review of the vapour pressures of ice and supercooled water for atmospheric applications[END_REF] and assuming the température profiles retrieved at 58-S, 15°S, 15"N we find a condensation altitude of about 104,97 and 93 km respectively.

stratospheric condensâtes, and (iv) the ro-vibrational émission lines of atmospheric species seen by CIRS at the latitudes considered in our study: CH4, CO, H20, G,H2.

These quantities were used as input to the NEMESIS retrieval code (Irwin et al., 2008) to perform a combination of forward model computation and retrieval scheme based on the method of opti mal estimation (Rodgers, 2000). The computation of the forward model spectrum used the correlated-k approach of Lacis and Oinas (1991) and included a Hamming apodization of Full Width at Half Maximum (FWHM) of 0.5 cm-1 to reproduce the instrumental line shape. The retrieval scheme was used to optimize the fits and détermine the model free parameters including the water vapor abundance. This method was successfully applied to model the FP1 spectrum in Cottini et al. (2012) to reriieve surface tempéra ture. We solve the radiative transfer équation for 147 spherical atmospheric layers, using as a source function the thermal émis sion of the surface, for which a unit surface emissivity is assumed, and that of the atmospheric layers. The retrieval algorithm then iteratively computes a synthetic spectrum, compares it to the data and after applying a cost function, détermines the best estimate for the physical parameters in the model -the stratospheric aérosol profile and any necessary adjustments to the température profile and the mole fraction of included atmospheric gases. The cost func tion indudes two components: one that measures the quality of the fit to the spectra (similar to a x2 test) and another that mea sures the déviation of the retrieved parameters from a set of a pri ori quantities.

The continuum due to the CIA was calculated according to [START_REF] Borysow | Modelling of collision-induced infrared-absorption spectra of Hj-H2 pairs in the fondamental band at températures from 20 K to300 K[END_REF]Frommhold (1986a,b,c, 1987), [START_REF] Borysow | Modelling of collision-induced infrared-absorption spectra of Hj-H2 pairs in the fondamental band at températures from 20 K to300 K[END_REF][START_REF] Borysow | Modelling of collision-induced infrared-absorption spectra of Hj-H2 pairs in the fondamental band at températures from 20 K to300 K[END_REF]Tang (1993). For the N2-CH4 pair, we used CIA coefficient values increased by 50% as required to fit the continuum of the Cassini Descent Imager Spectral Radiometer (DISR) data (Tomasko et al., 2008) and the CIRS spectra [START_REF] De Kok | Far-infrared opacity sources in Titan's troposphère reconsidered[END_REF].

We hâve modeled the haze emission/absorption using the extinction cross sections of the hazes included in de Kok et al.

(2007b). Since scattering is negligible at these wavelengths for partides smaller than few microns, we hâve omitted it from our computations.

We hâve adopted the atmospheric vertical temperature-pressure profiles retrieved from CIRS data for three latitudes (58°S, 15°S, 15°N) from [START_REF] Anderson | Titan's aérosol and stratospheric ice opacities between 18 and 500 pm: Vertical and spectral characteristics from Cassini CIRS[END_REF] from the surface to 3.3 x 10"7 bar, corresponding to an altitude range of 0-600 km (Fig. 3). Spectroscopic information for the gas rotational lines in the far-infrared range was extracted from the H1TRAN 2004 database (Rothman et al., 2005). For CH, we hâve adopted the revised mole fraction of 1.48% in the stratosphère (Niemann et al., 2010) acquired by the Gas Chromatograph Mass Spectrometer (GCMS) on the Iluygens probe during in its descent to Titan's surface. In the stratosphère for H2 we assumed a uniform volume mixing ratio of 0.1% (Courtin et al., 2008). The geometry of the observations was also included in the computations.

An accurate model of the instrumental FOV is required to suc cessfully reproduce CIRS spectra and measure water abundance.

The FP1 FOV is circular and has a sensitivity with a quasi-exponential decrease from the center to the edge and a FWHM of 2.4 mrad.

For on-disk spectra the homogeneity of the field of view usually permits simple modeling with a single ray calculated for the detector center. For limb spectra, we bave to take into account the rapid decrease in atmospheric density with height and the variations of température and gas volume mixing ratio profile with altitude. In such cases, the FOV is not assumed to be uniform and a multiple ray model is required to fit the data We modeled the FOV using the minimum number of rays for which the synthetic spectrum and the water retrieval computation results became stable; this corresponds to nine rays with a step in altitude of 25 km. We also recomputed some of the results using 39 rays (step of 5 km) in order to show a smoother limb contribution function. The spectral radiance measured by the FP1 detector is modeled by a convolution of the emerging radiance at each point in the FOV (as described in Nixon et al. (2009) and in Teanby and Irwin (2007)), weighted by a response function for CIRS FP1 detector. This beam profile was determined for CIRS FP1 (Flasar et al., 2004) using Jupi ter as a point source. As 95% of the integrated response is contained in a radius of 1.95 mrad from the FOV center. the detector observes a maximum altitude range of about 70 km.

Independent line-by-line calculations to simulate the same ondisk FP1 sélections were also made using the Atmospheric Radia tive Transfer (ART) code that has recently been applied to CIRS data in Coustenis et al. (2010Coustenis et al. ( ,2007)). The code uses the most recent aér osol extinction dependence inferred from Vinatier et al ( 2012) and température profiles derived by fitting the v4 methane band at 1304cm"' in FP4 averages taken at similar conditions as the FP1 spectra. The spectroscopic parameters for ali the observed molé cules and isotopes are from GEISA 2009 (Jacquinet- Husson et al., 2011) and HITRAN 2008(Rothman et al., 2009). The results from these two different codes are very similar and their différence is smaller than the error bars on the data.

In

Results

We hâve retrieved the water vapor abundance from both ondisk and limb data assuming different vertical profiles. These were a constant water mixing ratio profile and three vertically increasing profiles from recent photochemical models: (a) Hôrst et al.

( 2008), (b) Wilson and Atreya (2004) and (c) [START_REF] Lara | Vertical distribution of Titan s atmospheric neutral constituents[END_REF].

The latter profile was adopted in Coustenis et al. ( 1998) for the first water détection on Titan by ISO. The constant mixing ratio profile (a priori assumed to be 0.1 ppb) was decreased to follow the satu ration vapor pressure curve in the lower stratosphère below the altitude where the relative humidity reaches 100% (see Fig. 3).

Using the saturation vapor pressure équation of water over ice of [START_REF] Murphy | Review of the vapour pressures of ice and supercooled water for atmospheric applications[END_REF] and assuming the température profile re trieved at 15"N we find a condensation altitude of 93 km. For tem pérature profiles retrieved at 15°S and 58°S we find respectively condensation altitudes of 97 km and 104 km. We hâve also computed the contribution functions -normalized inversion kernels -showing the sensitivity of each atmospheric layer to a variation of the HiO mixing ratio. These contribution functions were computed for each profile and for ail of the most intense water fines in order to provide an altitude range of validity of the retrieved val ues (Figs. 2 and4). Fig. 2 shows the contribution functions for ondisk observations computed at four different wave numbers: at 254cnr' (one of the two fines used for the ISO water retrieval) the upper shoulder of the contribution function is wider and sensi tive to higher altitudes compared to the other wave numbers used for the water retrieval in this work. In our case the fit of the 254 cm-1 line improves when using a profile increasing with alti tude rather than a constant profile. In Fig. 4 we show only the con tribution functions computed at wavenumber 202.75 cm-1, where the most intense water fine in the GRS spectrum occurs (discounting the fine at 150.5 cm-1 that is unusable due to an instrumental interférence).

For water retrievals obtained using a scaled constant water profile we show the retrieved mixing ratio values at the altitude where the water contribution function peaks for the assumed profile (Table 1). For the vertical error we use the Fuit Width Half Maximum (FWHM) of the contribution function for the corresponding water profile. We also retrieve a scaling factor to the water profile from each of the photochemical models considered in this work (Table 1).

On-disk water retrieval

To measure the water abundance from the on-disk average (0-30'N) data we first use a constant water profile. We retrieve a vol ume mixing ratio of 0.14 ± 0.05 ppb at 97 km (FWHM 93-130 km).

This value corresponds to a surface-normalized H20 total column density of 3.7 ± 1.3 x 101'1 molecules/cm2.

The largest source of error for the on-disk observations is due to the fact that the maximum of the contribution function occurs in the région where the water abundance is rapidly changing due to condensation. Other sources of error include small variations of the température profile in the stratosphère, random noise from the detectors, and a small dependence in altitude sensitivity with wavenumber.

We also fit the water fines for the three other water vertical dis tribution profiles [START_REF] Lara | Vertical distribution of Titan s atmospheric neutral constituents[END_REF][START_REF] Wilson | Mantle plumes and plate motions[END_REF]Atreya. 2004: Hôrst et al., 2008) and obtained the necessary scale factors to fit the data, which are shown in Table 1. These values, ranging between 0.11 and 0.63 times the considered profiles, show the retrieved water mole fraction to be less than predicted from these previous models.

We hâve analyzed two additional latitudinal bins from 45CS to 10°S and from 80°S to 45'S, centered on the latitudes corresponding to the température profiles previously retrieved from CIRS at 15°S and 58gS respectively [START_REF] Anderson | Titan's aérosol and stratospheric ice opacities between 18 and 500 pm: Vertical and spectral characteristics from Cassini CIRS[END_REF]. See Water vapor abundance results.

Water vapor rcttieved mole fractions On-disk average (0-30°)N The results obtained simultaneously for on-disk data using the independenr line-by-line ART code simulations of the same FP1 sélections confirm the retrieved water vapor value reported above within the error bars.

Limb water retrieval

The measurement of water vapor obtained by modeling the limb 1 spectrum under the assumption of a constant mixing ratio profile is equal to 0.13 ±0.04 ppb. According to the position of the peak of the corresponding contribution function, the radiance mostly originates from a région centered at 115 km (FWHM 95-165 km). Modeling the limb 2 spectrum we retrieved a water mix ing ratio using a constant water profile of 0.45 ± 0.15 ppb at an alti tude of 230 km . These values indicate an increase of the water mole fraction with altitude in the strato sphère from 115 km to 230 km of about three times.

The scaling factors to the model water profiles obtained from the two types of limb retrievals are shown in Table 1 and again illustrate the smaller amount of stratospheric water vapor detected by CIRS with respect to the ones predicted by the models considered in this study.

Conclusions

In this work we modeled CIRS data using a constant-withheight water vapor profile and assigned the retrieved mixing ratio to the altitude where the contribution function peaks.

By combining on-disk and limb observations we are able to constrain the vertical profile of water in the stratosphère from 12 mbar to 10-3 mbar. corresponding to altitudes between 93 and 280 km (considering the widths of the contribution fonctions).

In Fig. 5 we compare our water vapor retrieved values with the models of [START_REF] Lara | Vertical distribution of Titan s atmospheric neutral constituents[END_REF]Wilson and Atreya, 2004;[START_REF] Hôrst | Origin of oxygen species in Titan's atmosphère[END_REF].

The measurement of the stratospheric vertical profile of water adds useful constraints to the photochemical models of Titan's atmosphère. Qualitatively, the increase of the water mixing ratio with altitude is in agreement with an extemal source of oxygen and a lower altitude sink due to condensation. However, quantitatively, our retrieved abundance seems to be less (from ~0.1 to ~0.6) than predicted from the models considered in this work (see Table 1). We also observe that since the scaling factors to the photochemical models (a)-(c) in Table 1 are slightly different for the two limb altitudes, it implies that these models might hâve a slope for H20 not quite consistent with CIRS data. However, due to the rapid variation of the water vapor mole fraction with alti tude in the atmospheric région where water freezes, and where CIRS is actually observing (on-disk and limb 1 spectra), we should be cautious in assigning a slope to the water profile. Coustenis et al. (1998) fitted the ISO data with a scaling factor to the [START_REF] Lara | Vertical distribution of Titan s atmospheric neutral constituents[END_REF] water profile of 0.4^. We fitted the CIRS data assuming the same profile multiplied by a scaling factor of 0.48 + 0.07 (corresponding to a water column density of 3.8 ± 1.0 x 10'4 molécules cm-2). This resuit agréés with the scal ing factor from the ISO analysis although this agreement may be fortuitous because most of the ISO émission originates from above 300 km, with the contribution functions peaking around 400 km (Fig. 3 of Coustenis et al. (1998)). Since the beam size of ISO was much larger than Titan there is a strong émission from the limb occurring at high altitudes where 0.4 times the [START_REF] Lara | Vertical distribution of Titan s atmospheric neutral constituents[END_REF] profile is used. Due to its higher spatial resolution, the contribution functions for the CIRS nadir sélection using the [START_REF] Lara | Vertical distribution of Titan s atmospheric neutral constituents[END_REF] profile cover the range 95-145 km (at half maximum).

ISO also retrieved a water vapor abundance of 0.4 ppb assuming a constant mole fraction above the condensation level (Coustenis et al" 1998); we assign the same relative error bars as those de rived by ISO for the scaled Lira et al. (1996) profile: +0.3 and -0.2. From CIRS on-disk observations we retrieved a volume mix ing ratio of 0.14 ± 0.05 ppb around 97 km for latitudes 0-30°N. Our retrieval is only marginally consistent with the ISO détermination of 0.4t[!2 ppb above the water vapor condensation altitude. How ever it is necessary to make the assumption that the ISO results pertain to the same condensation altitude that we retrieve because unfortunately it is not reported in the ISO paper.

We now discuss how our retrieved water abundances compare with current models of oxygen photochemistry on Titan. The re trieved scaling factors are ail less than one, implying that there are still uncertainties in our understanding of oxygen processes on Titan.

In photochemical models of Titan prior to 2000 (e.g. [START_REF] Lara | Vertical distribution of Titan s atmospheric neutral constituents[END_REF], in orderto allow atmospheric production of CO, it was postulated that CO could be produced through a Chemical réaction be tween OH (available from H20 influx into the upper atmosphère) and CH3. However, it was emphasized by [START_REF] Wong | Evolution of CO on Titan[END_REF] that this reaction does not produce CO as previously assumed but instead it produces H20 [START_REF] Pereira | Température and pressure dependcnce of the multichannel rate coefficients for the CH3 OH System[END_REF].

Hence an influx of H20 or OH does not produce any significant abundance of CO and therefore C02 can be produced by an H20 in flux only with CO already présent (OH + CO -» C02 + H). For this reason these models were unable to reproduce the observed CO abundance and were substituted by other tnodels that suggest the existence of primordial CO in the atmosphère (Wilson and Atreya, 2004), or instead CO produced in the atmosphère using an extemal influx of O* rather than H20 or OH [START_REF] Hôrst | Origin of oxygen species in Titan's atmosphère[END_REF].

In the pre-Cassini model of Wilson and Atreya (2004) water is photolyzed to OH, which combines with CO to form C02 and other complex species. In this model, CO is assumed to be primordial on since the stratospheric abondances of photochemically produced species are highly dépendent on this parameter. As shown in Fig. 5, the water abundance retrieved in our study is best fit by the water profile with the lowest eddy diffusion coefficient value considered in their model (Ko =100 cm2 s-1). This is lower than the va lue (Ko = 400 cm2 s"1) they identified as best reproducing CIRS observations of hydrocarbon species and adopted in this work for comparison with our retrievals. However our results show that even the Hôrst model with K0 = 100 cm2 s-1 still has excessive water at the altitudes of our measurements.

Therefore, our work clearly points towards further refinement of oxygen chemistry in photochemical models of Titan's atmosphère.

Radar imaging

Filtering

Segmentation

Cassini synthetic aperture radar (SAR) images of Titan, the largest satellite of Saturn, reveal surface features with shapes ranging from quasi-circular to more complex ones, interpreted as liquid hydrocarbon deposits assembled in the form of lakes or seas. One of the major problems hampering the dérivation of meaningful texture information from SAR imagery is thespeckle noise, it overlays real structures and causes gray value variations even in homogeneous parts of the image. We propose a filtering technique which can be applied to obtain restored SAR images. Our technique is based on probabilistic methods and regards an image as a random élément drawn from a prespecified set of possible images. The despeckle filter can be used as an intermediate step for the extraction of régions of interest, corresponding to structured units in a given area or distinct objects of interest, such as lake-like features on Titaa This tool can therefore be used, among other, to study seasonal surficial changes of Titan's polar régions. In this study we also présent a segmentation technique that allows us to separate the lakes from the local background. Cassini's RADAR instrument opérâtes in the Ku-band (13.78 GHz, A = 2.17 cm), at both high and low resolution, viewingTitan's surface in four modes: imaging, altimetry, scatterometry and radioinetry (Elachi et al, 2004).

The SAR mode is used at altitudes lower than 4000 km.

resulting in spatial resolution ranging from about 350 m to more than 1 km. Images are acquired either left or right of nadir using 1-5 looks while a swath, 120-450 km in width, is created from five antenna beams. SAR coverage is dépendent on spacecraft range and orbital geometry (Elachi et al., 2004) and the polarization vector), near-surface roughness, or nearsurface dielectric properties.

The images obtained using SAR revealed that Titan hosts a very complex surface formed with features such as lakes, mountains, fluvial river networks, possible volcanic-like features and dunes Qaumann et al., 2008;Lopes et al., 2010 and references therein) which resemble Earth-like geomorphological structures (Coustenis and Hirtzig, 2009). However, both the material and the environmental conditions shaping their respective surfaces are considerably different. Despite these différences, the exogenic mechanisms forming the surficial expressions may be similar.

Indeed, one dominating terrestrial surface-affecting procedure is the water cycle while on Titan an active methane cycle is at play (Atreya et al., 2006). The impact of the endogenic processes in the surface construction on Titan is still under investigation through geophysical interior models (e.g. Tobie et al., 2005). Radar images provided evidence of drainage and branching channel networks with subdivided channels (Perron et al., 2006: Soderblom et al., 2007;Lorenz et al., 2003;Burr et al., 2009), as well as distinct fluvial erosional patterns (Jaumann et al., 2008) that indicate dynamic and evolved surface, processes. Additionally, SAR ima gery observed lake-like features in several swaths (Stofan et al" 2007;Turlle et al., 2009;Hayes et al., 2010;Wye et al., 2010).

Other features, identified by radar, are the potentially volcanic in origin structures (Elachi et al.. 2005;Lopes et al., 2007a;Soderblom et aL, 2009;NeLson et al., 2009;[START_REF] Wall | Cassini RADAR images at Hotei Arcus and western Xanadu, Titan: evidence for geologically récent cryovolcanic activity[END_REF], the impact cratcrs [START_REF] Wood | Impact craters on Titan[END_REF], the mountain chains (e.g. Radebaugh et al., 2007) and the linear dunes (Elachi et al., 2006;Radebaugh et al., 2008) 2. Lacustine features of Titan's surface

The thermodynamical conditions dominating Titan as well as its Chemical profile allow the hydrocarbon liquid phase to exist on its surface [START_REF] Kouvaris | Phase cquilibrium of méthane and nitrogen at low températures: Application to Titan[END_REF]Thompson et al., 1992;Atreya et al" 2006). Indeed, the température of its icy surface has been measured in situ (10.34°S, 192.34'W) [START_REF] Tomasko | Rain, winds and haze during the Huygens probe's descent to Titan's surface[END_REF] by Huygens at 93.7 K (Fulchignoni et al., 2005), while the Cassini composite infrared spectrometer (CIRS) found by remote observa tions an average température of 92 K between 2004 and 2008 [START_REF] Jennings | Titan's Surface Brightness Températures[END_REF] Cassini RADAR instrumentation confirmed the presence of lakelike features on the surface of Titan. Several models hâve suggested organic précipitation on Titan (Toon et al.. 1988;[START_REF] Lorenz | The life, death and afterlife of a raindrop on Titan[END_REF][START_REF] Graves | Rain and hail can reach the surface of Titan[END_REF] either violently through torrential storms (Hueso and Sanchez-Lavega, 2006) or smoothly through drizzle at the lower atmosphère [START_REF] Grard | Electric properties and related physical characteristics of the atmosphère and surface of Titan[END_REF] and rainfalls originated from occasional short-Iived clouds (Griffith et al, 2000[START_REF] Griffith | The evolubon of Titan's mid-labtude douds[END_REF][START_REF] Lorenz | Titan's surface before Cassini[END_REF], such phenomena hâve not been recorded yeL

The first recording of organic liquid pools was accomplished during the Tl 6 flyby at the northern polar locations, where more than 75 lake candidate features were identifïed (Sotin, 2007;Stofan et al., 2007). In 2008, the Cassini/RADAR swaths counted in both north and south polar régions, above 50"N and 50CS, respectively, more than 655 lake-like features (Hayes et al., 2008). SAR imaging shows lake-like features separated into three classes; dark lakes, granular lakes, and bright lakes (Hayes et al., 2008). Dark lakes are interpreted as liquid filled, while bright lakes are interpreted as empty basins and granular lakes are inferred as transitional between dark and bright lake features.

From a geomorphological aspect the lakes on Titan span over the range of observed morphologies on Earth (Stofan et al., 2007;Mitri et aL, 2007;Hayes et al" 2008). They are rimmed features, from circular to irregular and some with distinct edges, steep margins and smooth surfaces (Stofan et al., 2007) and they show very low microwave backscatter. Some of them are surrounded by a drainage network of dark channels, which may supply them with liquid (Stofan et al" 2007), while some are not.

It is suggested that this liquid is a methane/ethane mixture, with smaller concentrations of nitrogen and higher order hydrocarbons/nitriles (Lunine et al., 1983;Mitri et al., 2007;Brown et al., 2008;Raulin, 2008).

The lakes' radiometric brightnesses which appear to be warmer than the surrounding région (Janssen et aL, 2009) are consistent with the high emissivity expected for a smooth surface with the real part of the low dielectric constant between 1.7 and 1.9 of liquid ethane-methane solutions [START_REF] Bames | Near-infrared spectral mapping of Titan's mountains and channels[END_REF] with most possible value the 1.9. However, the imaginary part is still under investigation [START_REF] Notamicola | Cassini radar data: estimation of Titan's take features by means of a Bayesian inversion algorithm[END_REF] This study provides a qualitative method of récognition of lake-like features from Cassini SAR images. Tire intended goal is to label régions in an image into three classes (dark lakes, granular lakes and the local background). First, a filtering technique is applied to obtain the restored image. Then, a method of supervised segmentation is used. The segmentation method based on the minimum Euclidean distance is used here.

Filtering

Strip mapping SAR consists of a large antenna which is synthesized from many small antennas and rentains fixed with respect to the radar platform, so that the large antenna illumi nâtes a strip of the ground. This technique is used to improve the azimuthal resolution. As the platform moves, a sequence of closely spaced puises is emitted and the returned waveforins are recorded. An image is computed after the cohérent sum of reflected monochromatic microwaves. The image is distorted by a strong granulation, called speckle. Speckle noise exists in ail types of cohérent imaging Systems and its presence reduces the resolution of the image and the detectability of the target Speckle noise is not only signal dépendent but is also spatially correlated and reduces the effectiveness of image réduction. The Cassini RADAR measures the normalized backscatter cross-section (<xo) of Titan's surface (Ulaby et al., 1982). The SAR textures are generally affected by multiplicative speckle noise. In order to reduce the speckle noise, the multilook technique, based on incoherently averaging the independent neighboring pixels, is used to estimate the characteristics of the same ground area.

The statistical characteristics of multilook data départ considerably from those of single-look data. Multilook data tend to inix some physical and statistical properties of the terrain. The terrain appears more homogeneous and the multilook averaging will tend to be close to the Gaussian statistics (Chitroub et al., 2002).

Speckle noise overlays real structures and causes gray value variations even in homogeneous parts of the image and also makes automatic segmentation of such images difficulL We propose here a filtering technique that can be applied to obtain the SAR restored images. After filtering the structured parts of the image can be much better separated. The total sum preserving regularization (TSPR) fïlter, is based on a membrane model Markov random fïeld approximation optimized by a synchronous local itérative method (Bratsolis and Sigelle, 2003). The final form of despeckling gives a sum-preserving regularization for the pixel values of the image.

Image formation is the process of computing (or refining) an image both from raw sensor data that is related to that image and from prior information about that image. Information about the image is contained in the raw sensor data, and the task of image formation is to extract this information so as to compute the image.

Image space is the set F of model images that represent the true, underlying physical distributions that are measured by the sensors.

Working in a probabilistic imaging problem, an a priori knowledge about an image is most naturally incorporated through the use of a prior probability distribution (or prior) P(J) on the image space. Our objective here is to reconstruct the original (or true) image from its degraded version.

A random field is an appropriate model for image values.

Random variables characterized by conditional priors that account for local interactions are often used as natural and convenient priors in imaging problems. These priors are placed directly on the image space. However, the fundamental prob ability distribution on the field is the joint probability distribution P(f), and this is difficult or impossible to specify directly. One needs to verify that the chosen spécification of conditional distributions is sufficient and consistent in the sense that a unique joint probability distribution corresponds to this set of conditional probability distributions.

Assume that an image is formed on a finite rectangular lattice are équivalent under the positivity condition [START_REF] Besag | Spatial interaction and the statistical analysis of lattice Systems (with discussion)[END_REF].

Let g be the degraded observed image and / the restored image.

The general problem is now to estimate/from a set of input data g.

The Bayes theorem allows us to obtain the a posteriori probability for the field / given the data g:

P(g\T)P(f) P(.g) (1)
where P(f) is the a priori probability for the field /, P(g]f) is the conditional probability for the data g given / and l\g) is the probability of g, wliich is independent of/.

Considering a multilook image, we assume for the conditional probability P(g\f) a Caussian distribution (Cliitroub et al., 2002) and we hâve

P(g\f) = exp- fé-fc)2 là1 C, (2)
where c is the standard déviation of the Caussian distribution and C, a normalization constant.

The probability P[f) is given by exp-0E(/s-/r)2

Cr.s> (3) where /? is a smoothness factor, C2 a normalization constant, and the sum runs on ail pairs of neighboring sites. From ( 1) we obtain

P(f\g)= ^exp-U(f) (4)
where Z, also called the a posteriori partition function, is a normalization constant and U{f} is the total a posteriori potential function. which is written as m = +/>£<&(

The related local conditional potential functions are given by v,(f) = -ÿ + P £ (fi -A)2 C6)

reV,

The a posteriori usual maximum (MAP) estimate is assigned to the value / which maximizes (4), i.e. which minimizes the a posteriori energy U(f) (Bratsolis and Sigelle, 2003). We accept as a fast method of optimization a synchronous minimization of the local potential energies which we named synchronous local itérative method.

Our final TSPR filter is an itérative filter which uses the équation given by = Vnkl (?

)
where * is the 2-D convolution operation. Titus, at each step the current pixel value will dépend in a regularizing manner on its neighboring ones, according to the magnitude of parameter X. lt is also obvious that positivity is preserved when OsA^l.Kisa normalized matrix given by According to the notation of our method, it is évident that <r0 = gs and <(To> = <&> = </s>-The TSPR method préserves the mean values of local hoinogeneous régions and decreases the standard déviation up to six times (Bratsolis and Sigelle, 2003).

The segmentation method

The purpose of segmentation is to divide the image into spécifie régions that correspond to structural units in the scene or distinguish objects of interest. These régions are characterized by spatially connected, non-overlapping sets of pixels sharing common properties.

In

where /q with i = l 3 présents the mean values (<oo>) corresponding to three sampled régions of interest. The algorithm is applied for every site s and for every value fs of the restored (filtered) image. The value fi, which minimize the distance de gives the characteristic label (black, dark gray or light gray) to the segmented image of Ftg. 4, where fi, corresponds to -20.27 db, fi2 to -18.64 db and fi3 to -13.78 db.

We begin our analysis with the Cassini SAR image (PIA08630, NASA/JPL) acquired during the T16 flyby on July 22, 2006, at high latitudes near the north pôle. This image is centered near 80"N, 92°W. Our image has dimensions 750 x 3100 in pixel size and the pixel scale is set to 175.558 m per pixel ( ~256'), but the actual SAR resolution is around 350 m per pixel and the image has been subjected to some interpolation. Cassini SAR images across this région contain numerous very dark splotches with sharp-edged boundaries, which may be filled with hydrocarbon liquid.

After using the despeckling filter TSPR we apply the segmentation method. Figs. 1 and 2 illustrate the initial image derived form the Cassini SAR while Fig. 3 is the filtering resuit of the previous image.

The despeckling TSPR filter is used to smooth out the multiplicative Image ratio between the initial (Fig. 2) and the despeckled (Fig. 3) image.

Table 1 Characteristics of the sampled régions of interest. and the région 3 (local background) we take a value for the ratio std{o0)l (aQ > equal to 0.4 and for the région 1 (dark lakes) equal to 0.05 which means that the région 1 is flatter than we were waiting.

Régions of interest

The characteristics of the sampled régions of interest are listed in Table 1.

We can see the segmentation results in Hg. 5, where the black area corresponds to the dark lakes, the dark gray to the granular lakes and the light gray to the local background.

Discussion

Titan is indeed an active planetary body and both endogenic and exogenic processes hâve most possibly left their marks on the surface. Hence, one should focus on the surface geomorphology in order to identify any local tectonic field and estimate the importance as well as influence of each forming mechanism.

The TSPR filter, in combination with the minimum Euclidean distance method of supervised segmentation, can be used to extract régions of interest on the surface of Titan, such as lakes or seas using Cassini SAR images. Such despeckle filter can be applied in studying other surface features on Titan like the drainage networks, the équatorial dunes and the impact craters, where different textures appear.

Our approach allows to isolate each distinct surface feature from its surroundings and to study their distribution through out the surface. Then, in combination with radiative transfer modeling using Cassini Visual and infrared mapping spectrometer (V1MS) data, we can infer about the relation between surface composition and morphotectonic structures. When we détermine in a more accurate way the shapes of several surface structures,

we will be able to study their global distribution and perfonn effectively classifications.

Since the Cassini mission extended its operational duration, new swaths, overlapping the liquid areas, hâve shown that the lakes hâve been evolving during the past years. Recent studies, focused on lake Ontario, the largest lake of the Southern hemisphere, showed a significant recession of its shoreline (Turtle et al., 2009;Hayes et al., 2010;Wall et al., 2010) and support the hypothesis that these liquid deposits are not stable structures but evolve with time by expanding in winter and shrinking in the summer (Sotin, 2007). Applying the segmentation method on the same area at various periods of time, we can identify possible enhancements or réductions of the liquid coverage of the région. In particular, the temporal variation of the dark spots can provide information on the évolution of the lake System and consequently help us to better understand the methane cycle on Titan and therefore the mechanisms linked with the lake surface features, their origin and fate, through a global temporal and spatial coverage.

The passage from qualitative to quantitative results, requires to apply the aforementioned method in the same régions of interest with the same obsetvational characteristics at different time periods in order to measure the surface each time.

Using this temporal dataset will help us evaluate the volume variation through time and estimate the hydrocarbon loss rate, Cassini evolves in the Satumian System since 2004 and will continue returning data until at least 2017 in the extended mission. Althougli Cassini SAR recordings hâve unveiled new characteristics of Titan's surface, more advanced instrumentation witli higher resolution is necessary to give a full picture of this complex environment.

A new long-term mission to Titan which will perform efficient in situ experiments in order to study widely the nature and the dynamics of its environment would be required.

Several proposais hâve been studied to this end. Recently, the 2008 study of a flagship mission, Titan Saturn System Mission (TSSM), focused on Titan and Enceladus exploration [START_REF] Reh | TSSM final report on the NASA contribution to a joint mission with ESA[END_REF]Coustenis et al., 2009). TSSM consists of an orbiter, a lake lander and a balloon (montgolfière) and aims to a complété investigation scrutinizing thoroughly the. whole satellite front the exosphere to its interior. The montgolfière would contain advanced instnimentation such as the Titan radar sounder ( > 150 MHz), which could perform a detailed surface investiga tion. In addition to topography, the radars on botli the orbiter and the montgolfière would provide extended information regarding the lakes coverage. The proposed filtering and segmentation rnethod in this paper would be a helpful tool in enhancing the return of the analysis of ail SAR data acquired on Titan and other objects as well as in the exploitation of future missions to Titan. 428 g g s s g-» §g-g |.i al.

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i?E? ? c1?!? tectonism [START_REF] Rosenau | Fore-arc deformation Controls frequency-size distribution of megathrust earthquakes in subduction zones[END_REF]Scheidegger, 2004;[START_REF] Lidmar-Bergstrôm | long term morphotectonic évolution in Sweden[END_REF] by studying landform évolution and dégradation, since tectonic features are subsequently subjected to exogenous processes. Major morphotectonic features on Earth are represented by mountains, ridges, faults and escarpments, as well as by signifïcant types of linear features such as rifts, grabens and other linear terrestrial terrains that are subjected to érosion subsequently to deformational events (e.g., Scheidegger, 2004). However. geology on Earth is dominated by active plate tectonics where rigid lithospheric plates float and move on a plastic asthenosphere.

Although the other planetary bodies in our Solar System possess different surface and internai conditions, bodies like Titan, Europa and Enceladus may possess a Iiquid water layer underneath their icy crust. If confïrmed, then similarly to rocky plate tectonics on Earth, rigid ice plates may rupture and collide, floating over such a Iiquid substrate layer, resulting in surficial features, which may be reminiscent of terrestrial édifices. It is therefore possible that other pianets and moons in the Solar System harbor "tectonic activity" in varying degrees and even exhibit morphotectonic features on their surfaces, which are subsequently modified by exogenous processes.

Venus appears to hâve no plate tectonics due to a high surface température and a higher density of its lithosphère compared to that of the mantle, which prevents a subduction régime, despite the fact that the mantle is convecting (Niinmo and McKenzie, 1998). However, the planet shows deformation and morphotec tonic features such as faults, mountain crests and rifts, which probably originated from lithospheric movements in association with volcanism (Jull and Arkani-Hamed, 1995;[START_REF] Nimmo | Volcanism and tectonics on Venus[END_REF] association with active diapirism from the mantle [START_REF] Mège | A plume tectonics model for the Tharsis province[END_REF] and the Elysium région, which was a result of volcanic activity [START_REF] Hall | Elysium région Mars: tests of lithospheric loading models for the formation of tectonic features[END_REF]. lo, Jupiter's moon, présents morphotectonic features with no apparent association to plate tectonic activity. The mountains of lo are formed by stresses at the bottom of the lithospheric layer and subséquent uplift through thrust faulting System [START_REF] Schenk | Origin of mountains on lo by thrust faulting and large-scale mass movements[END_REF].

A good candidate for the study of morphotectonic features in the Solar System appears to be Titan. With a diameter of 5150 km

(1/3 that of the Earth), Titan is the largest satellite of the Saturnian System and the second in the Solar System, after Ganymede the moon of Jupiter. The température and pressure conditions at the surface near the equator are 93.65 ± 0.25 K and 1.467 +1 hPa, as measured by the Huygens probe Atmospheric Structure Instrument (HASI) (Fulchignoni et al., 2005). Titan is recognized as a world bearing several resemblances to our own planet, with respect to its atmosphère and to its surface morphology. Titan's dense atmosphère consists mainly of nitrogen (~97%), methane (1.4%) and hydrogen ( --0.2%) with traces of hydrocarbons, nitriles, oxygen compounds and argon (see table 6.4 in Coustenis and Taylor, 2008). This complex atmo sphère renders the surface difficult to access and analyze, apart from within a few methane spectral Windows in the near-infrared where the methane absorption is weak (Griffith et al., 1991).

Thirty-two years after the Voyager encounter in 1980, Cassini is today able to probe Titan's surface with a spatial resolution reaching a few hundred meters per pixel (RADAR), while the Huygens probe achieved the first in situ measurements in 2005

(for instruments and resolutions see Section 2). Even though Titan's surface morphology resembles that of the Earth, it is made of materials and subjected to surface conditions very distinct from the terrestrial ones. Indeed, morphotectonic features such as mountains (e.g. Radebaugh et al., 2007;Lopes et al., 2010), ridges [START_REF] Soderblom | Topography and geomorphology of the Huygens landing site on Titan[END_REF]Mitri et al., 2010), faults (e.g., [START_REF] Radebaugh | Régional geomorphology and history of Titan's Xanadu province[END_REF], rectangular drainage patterns and cryovolcanic structures are most likely controlled, at least in part, by tectonism (Burr et al, 2009).

Atmospheric processes, like cloud formation and précipitation create extensive fluvial features on the surface, as observed by Huygens near its landing site and constitute the visible part of an active methane cycle (Atreya et al., 2006;Coustenis and Taylor, 2008;Lorenz and Mitron, 2008;Raulin, 2008;Brown et al., 2009; Coustenis and llirtzig, 2009;Lebreton et al., 2009). The préserva tion limit of 100 Myr for this atmospheric methane requires a réservoir that would replenish occasionally the atmosphère [START_REF] Lunine | The methane cycle on Titan[END_REF]. One of the most prédominant théories suggests that methane sources exist in Titan's interior (e.g., Tobie et al., 2006;Fortes et al., 2007). Since volcanism is a major process associated with the terrestrial carbon release [START_REF] Bolin | Carbon Cycle Modelling[END_REF], cryovolcanism may play a similar rôle in the methane supply (Sotin et al., 2005), as well as significantly influence Titan's surface morphology.

Geophysical models suggest that Titan's partially differentiated interior consists of a silicate core (~1800km thick), a high-pressure ice mantle (~400km), a liquid layer of aqueous ammonium sulfate (50-150 km thick), and an external icy shell 100-170 km thick that possibly contains clathrate hydrates (Tobie et al., 2005;Fortes et al., 2007;[START_REF] Grindrod | The long-term stability of a possible aqueous ammonium sulfate océan inside[END_REF].

Castillo-Rogez and Lunine (2010) suggested possible déhydration of the core's hydrated silicates, which impacts the geophysical structure of the satellite as well as the possible intentai océan.

Regarding the icy shell, the methane stored as clathrate hydrates within the ice shell is a plausible methane réservoir that can replenish the atmosphère via cryovolcanism (Sotin et aL, 2005).

Indeed, surface discontinuités such as faults and fractures, which are probably the resuit of tectonic and volcanic-like processes, could provide the pathways of internai methane release to the atmosphère. The morphotectonic structures on Titan's surface provide good evidence of such a mechanism, in the same way as, over extensive zones of geological weaknesses, magma and volatiles are released on the Earth's surface.

In the last eight years, despile continuous observations by

Cassini and the development of models and interprétations based on them, we still lack long-term in situ measurements and geophysical data of Titan's interior, in order to be in a position

to accurately evaluate its endogenetic potential and how it affects morphotectonic features. However, in this work we attempt to use similarities between the surficial morphotectonic features on Titan and on Earth as the key for deciphering Titan's endogenetic processes, in spite of the fact that our understanding of Earth's endogenetic processes is rather recent [START_REF] Wilson | Mantle plumes and plate motions[END_REF].

Titan surface observations

From the interprétation of Voyager 1 recordings, a global océan of dissolved ethane and nitrogen, several kilometers deep, was first assumed to cover the entire surface of Titan (Flasar. 1983;Lunine et al., 1983). However, ground-and space-based observations refuted this assumption by unveiling, within the methane "Windows" of weaker methane absorption (centered at 0.94, 1.08, 1.28, 1.59, 2.03, 2.8 and 5 pm), a solid surface with heterogeneous bright and dark features (Muhleman et al., 1990;Griffith, 1993;Smith et al., 1996;Gibbard et al., 1999;Meier et al., 2000;Coustenis et al., 2001). The Cassini orbiter arrived at the Saturnian System in 2004 equipped with two spectro-imagers capable to probe down to the surface via several of the nearinfrared Windows: the Visual and Infrared Mapping Spectrometer (V1MS-with a typical resolution of 10-20 kin/pixel) and the Intaging Science Subsystem (ISS-with a typical résolution of 1 km/pixel). In the scope of this paper we also make use of the RADAR data from Cassini with a spatial resolution from 300 m to Endogenous, as well as exogenous dynarrüc processes hâve created diverse terrains with extensive ridges and grooves, impact units, icy flows, caldera-like structures, layered plains and stable liquid lakes (Mitri et al., 2007;Stofan et al., 2007). In addition, Cassini's radar has partially revealed the topography of Titan's surface, indicating several types of surficial expressions, which are non-tectonic. Features like dunes, lakes and drainage network are attributed solely to fluvial, aeolian and impact processes (Fig. 1). Thus, their formation is the resuit of exogenous processes with no influence of internai activity.

Morphotectonic features

Cassini's remote instrumentation and the Huygens lander brought evidence of many features on Titan's surface, which were probably formed by extension or compression of parts of the planetary solid crust due to endogenetic geological and geophy sical processes (Radebaugh et al., 2007;[START_REF] Soderblom | Topography and geomorphology of the Huygens landing site on Titan[END_REF] Please cite this article as; Solomonidou, A., et al" Morphotectonic features on Titan and their possible origin. Major mountain and ridge features on Titan. Huygens/January 2005 DLSR (Tomasko et al., 2005; Soderblom et aL, 2007b) 30«S, 315<W Proposed mechanisms for the formation of mountains and ridges of Titan. 2a.

Table 2a points at two intriguing aspects: mountain-like édifices exist at almost ail latitudes on Titan; however, they are concentrated in the équatorial région at latitudes between 30"S and 30"N (topes et al., 2010). Their height is significantly lower than that of the terrestrial mountains ranging from 100 to 2000 m (Barnes et al., 2007;Radebaugh et al., 2007;Soderblom et at, 2007b;Sotin et at, 2007;Mitri et at, 2010). This inay be partly materials that surround these structures (Radebaugh et al., 2007).

Alternative hypothèses include the construction of Titan's moun tains with materials with properties preventing height growth (Radebaugh et al., 2008) and the effects of high température gradients on the ice shell which, according to the calculations of Mitri et al. (2010), resuit in mountains from 100 to 2700 m high.

Similarly to Earth where terrain topography is defined by the interaction of tectonism and érosion (Montgomery and Brandon, 2002), we suggest here that there is a strong connection between slope morphology and erosional rates on Titan due to its extreme conditions of hydrocarbon rainfall and/or winds.

Fig. 2 présents three portions of the T8 and T43 RADAR swaths that provide the most reliable evidence so far for the existence of mountains and ridges on Titan.

Related mechanisms and terrestrial analogs

The mechanisms for mountain formation on Titan are summarized in Table 2b and include pure extension (Tobie et al., 2006;[START_REF] Sotin | Cassini/VIMS observatoire of Titan during the T20 flyby[END_REF]Radebaugh et al., 2007-scénario 2), pure compression (Radebaugh et al., 2007-scénario 1;Mitri et al., 2010) and transitions between compressional and extensional stresses (Mitri and Showman, 2008). Extensional deformation is observed on ali icy moons, but if orogenesis on Titan can be attributed to compressional forces this will render Titan unique.

Orogenesis in the terrestrial analog is predominantly a com pressional event due to the coming together of the lithospheric plates floating on the plastic asthenosphere. features sucli as linear mountain ranges and valleys. Fluvial, aeolian and otlier exogenous processes subsequently modified these features.

However, convective stresses on silicate bodies tend to be larger and tlieir rheological length scales are typically greater (Table 1).

In Titan therefore it is essential to note that local/regional ratlier than global stress mechanisms are commonly suggested in the models for mountain building. Régional or local stress mechanisms invoked in these models include (a) convection, which dépends in the ice shell thickness, (b) local gravity, and (c) ice viscosity, which dépends on the température and mostly on the grain-size [START_REF] Barr | Onset of convection in ice I with composite Newtonian and non Newtonian rheology: application to the icy Galilean satellites[END_REF]Collins et aL, 2009). Indeed, on large icy satellites, layers of high-and low-pressure ice may convect separately [START_REF] Mckinnon | Geodynamics oflcy Satellites[END_REF].

Among the local stress mechanisms presented above, one can include latéral pressure gradients that may hâve as a conséquence the latéral flow of floating ice shells on their low viscosity base.

Rigid ice floes rupturing and colliding are reminiscent of plate tectonics albeit in a randont fashion. This could lead to the création of blocks of high-standing topography that would sub sequently be subjected to erosional processes. Such elevated morphotectonic features on Titan, as mountains, ridges, hills and ranges, indicate a formation preference around the équatorial zone of the ntoon (Barnes et al., 2007;Radebaugh et al., 2007;Mitri et al., 2010).

Morphotectonics of faults and transverse processes

In this section, we describe Titan's morphotectonic features as a combination of a formation process and one or more superimposed "transverse processes" that occurred at the same time or subsequently, modifying the initial shape of the feature. élévation [START_REF] Schumm | Active Tectonics and Alluvial Rivers[END_REF]. On Titan the observed faults, fractures, canyons and ground linéaments are most likely the results of crustal movements due to tectonic and/or volcanic processes as well as structures associated with fluvial networks controlled by tectonism (Table 3). An investigation of canyon fonnation can augment our understanding of Titan's geology. A study in 2009 by Burr et aL (2009) provided evidence on the influence of subsurface tectonic activity on drainage patterns as observed by SAR data.

Based on that, we can infer that these fluvial networks are morphotectonic features since there are indications that both tectonic and fluvial processes operate for their formation. Titan studies based on observations and mapping, hâve suggested the presence of fault formations. As expected from the study of other icy satellites, the majority of formation mechanisms suggest extensional style tectonism. In the. areas where fluvial networks seem to be controlled by tectonic patterns, tire surface material seems to hâve the proper elasticity to create linear fractures. At the Huygens landing site, as proposed by [START_REF] Soderblom | Topography and geomorphology of the Huygens landing site on Titan[END_REF], the linear structures can function as the idéal path for the hydrocarbon liquids to escape towards the surface. Furthennore, the observed radial fault System around the possible calderas of Hotei Regio (Soderblom et al., 2009), argue in favor of cryovolcanism since the identifica tion of radial faults around caldera formations also on Earth, are indicators of ground élévation due to volcanic activity. Fig. 5 displays a complex feature on Titan (71'S, 240°W), which -even if it is not confirmed by radar data Processing yet -it appears to be a canyon-like morphotectonic feature since it consists of a sinuous dark, rather narrow feature with tributarylike off shoots and it is limited on ail sides by high albedo i.e.

elevated terrains (e.g. Radebaugh et aL, 2007). The morphology of the bright and the dark shape of this région resembles the terrestTial analog, which is the Grand Canyon in the United States.

This terrestrial feature is adjacent to the Basin and Range Province that was mentioned earlier. The formation of the Grand Canyon on Earth is the end resuit of the extensional tectonics that formed the Basin and Range Province and of continuous rifting and érosion [START_REF] Sears | Ceological structure of the Grand Canyon supergroup[END_REF].

From Table 3 we can distinguish a tendency for preferential formation of these features between 10°S and 26°S, that is, within the zone where mountains are formed. Such observations imply crustal movements are more frequent within this zone than around the pôles, including compressional and extensional stres ses. However, Titan's liinited-coverage observations (less than 50% of the surface), with instruments incapable of precisely uiweiling its geology, make this aspect a subject for future exploration.

Cryovolcanism and association with morphotectonics

Plate tectonics and volcanism are strongly associated on Earth [START_REF] Mcdonald | Mid-ocean ridges: fine scale tectonic, volcanic and hydro thermal processes within the plate boundary zone[END_REF] and this can also be the case on Titan in the presence of tectonic features overlying a liquid water océan.

These can function as leading 'pathways' for the ammoniawater cryomagma to reach the surface and for the release of methane. Liquid pockets with inethane clathrates and with a high ammonia mass concentration in a water solution can dissociate in the ice Shell and eventually exsolve on the surface and in the atmosphère (Tobie et al., 2006;Mitri et al., 2008). Hence, cryovolcanism can also act as the dynamic force that deforms tectonic features. Cryovolcanism is believed to represent an important geological process in the history of several icy Saturnian satellites and other icy satellites, such as Triton [START_REF] Kargel | Cryovolcanism on the icy satellites[END_REF][START_REF] Fagents | Considérations for etTusive cryovolcanism on Europa: the post-Galileo perspective[END_REF]Lopes et al,, 2007a). Caidos ( 2001) stated that tectonic extension could trigger cryovolcanic éruptions by reducing the minimum normal stress in an aquifer to a value ii. Fluvial sapping of bedrock that enlarges tectonic zones of weakness.

TO/July 2004 TA/October below the pore pressure. On Titan, a few topographie features are candidates for large volcanic édifices (Table 4).

The mechanism of cryovolcanism resembles terrestrial type volcanism, however the cryomagma dilTers from the terrestrial magma in composition, texture and température. Indeed, the cryomagma i.e. aqueous solutions of ammonia, methane, salts, etc., can be found at températures well below the freezing point of pure water and degassing replaces the traditional silicate volcanism. Cryovolcanic structures are openings, or ruptures, on a pianetary surface or crust, allowing for various internai products like water, other Chemical components, gases and cryoclastic ash to escape from the planet's interior (Fortes et al., 2007).

On Titan, internai heating due to radiogenic decay and tidal forces along with pressure fluctuations may trigger cryomagma éruptions. The cryolava déposition would then happen at tem pératures much lower than the terrestrial ones (Davies et al., 2010). In general, the températures of most terrestrial magma types range from 1150 to 1470 K while a plausible range of the Titan cryomagma températures is between 177 and 239 K (Mitri et al., 2008). Besides prédictions front theoretical modeling we also hâve surface features indicating possible cryovolcanism on Titan. Indeed, lobate and fan shaped features seen in radar images hâve been interpreted as cryovolcanic in origin, as for instance, the lobate circular structure of Ganesa Macula (5CPN, 87"W) (Lopes et al., 2007a). The structure contains bright rounded features, interpreted as cryovolcanic flows, while the curved or linear shapes are linéaments that could be caused by élévation of the crust due to cryovolcanic activity (Lopes et al., 2007a). Other such features like Tui Regio, Hotei Regio and Sotra Facula are more extensively discussed hereafter. The identification of cryo volcanic structures is rendered difficult mainly for two reasons:

firstly, the masking of cryovolcanic features by the interaction with major exogenetic processes, e.g. by fluvial or aeolian deposits (Lopes et al., 2010). As an example of the latter case, Lopes et al. ( 2010) report in the Winia Fluctus (dS^N, 30°W) a dune field, which has partially covered a cryovolcanic édifice. Secondly, the Cassini instrumentation with relevance to cryovolcanic investiga tions (SAR and VIMS) is not adéquate on terms of spatial and spectral resolution (Elachi et al., 2004;[START_REF] Brown | The Cassini Visual And Infrared Mapping Spectrometer (Vims) Investigation[END_REF].

However, we hâve currently some cryovolcanic candidate features. The most probable ones are three areas located close to the equator. Tui Regio, as well as Hotei Regio, lie within the bright région of Xanadu (100°N, 15"S), a large, reflective équator ial area. Tui Regio présents relatively high 5-p.m reflectivity and its size is 1500 km long and 150 km wide [START_REF] Mccord | Composition of Titan's surface from Cassini VIMS[END_REF]. lt is a massive flow-like terrain, which resembles flow fields in volcanic areas on Earth. In 2006, Barnes et al. (2006) noted tliat there are two bright and long areas within Tui Regio that are fïlled with a material darker than the surrounding terrain, forining a trench. These areas, located within the northwestem portion of Tui Regio's flow and including the trending dark linear marks (Fig. 6), may hâve a régional tectonic origin (Barnes et al., 2006).

Additionally, the latter authors pointed out a linear dark feature with similar spectral behavior which surrounds the Southern bright région, suggesting chat it might also be formed by régional tectonics. A recent and very reliable candidate for cryovolcanic activity is Sotra Facula; an area 235 km in diameter including a 1 km high mountainous peak next to a 1.5 km wide crater-like feature from which lobate flows seem to originale [START_REF] Kirk | La Sotra y las otras: topographie evidence for (and against) cryovolcanism on Titan AGU[END_REF]Lopes et al., 2010). Indeed, this area displays varying topography with adjacent uplifts and pit features suggesting probable tectonic control. Hotei Regio is also another candidate for the presence of tectonic features within a probable cryovol canic région [START_REF] Wall | Cassini RADAR images at Hotei Arcus and western Xanadu, Titan: evidence for geologically récent cryovolcanic activity[END_REF]. Hotei Regio extends over 700 km and includes a 1 km wide topographie dépression, characterized as a basin fïlled with flow-like features, a ridge-like mountainous terrain that surrounds the basin, dendritic channels, two calderalike structures, dark blue patches (as seen in VIMS infrared images; a color that suggests enrichment in water ice), and possible alluvial deposits (Soderblom et al., 2009). Hotei Arcus, a bright arc in the Southern margin of Hotei Regio, may represent a heavily eroded crater [START_REF] Barnes | A 5-micron-bright spot on Titan: evidence for surface diversity[END_REF]. This assumption reinforces the hypothesis of interplay among the different types of processes. Soderblom et al. (2009) correlated VIMS and RADAR images in order to unveil the geological history of this area. Their interprétation suggests that a wide range of processes occurredor are still occurring -in this varied terrain, including tectonism.

Furthermore, they suggested that impact-induced faults created Candidates of major cryovolcanic features on Titan and their association with volcanotectonic processes. The two dark morphotectonic features (Fig. 6) nortli of Hotei Regio, are interpreted as cryovolcanic calderas by Soderblom et al. (2009).

A terrestrial région that appears to resemble the évolution of Hotei Regio is Harrat Khaybar, located in north of Médina in Saudi Arabia. It is a 14,000 km2 volcanic field that was fornted by éruptions along a 100 km N-S linear vent System. The area contains multiple volcanic rock types, lava flows, and volcanic structures such as calderas and dômes (e,g., [START_REF] Baker | Jebel al Abyad: a recent alkalic volcanic complex in western Saudi Arabia[END_REF].

The internai mechanism that most likely formed the terrain is a mantle plume causing diffused lithospheric extension (Chang and Van der [START_REF] Chang | Mantle plumes and associated flow beneath Arabia and East Africa Earth and[END_REF]. The association of the volcanic centers that lie over a linear zone of weakness with the Red Sea transforin fault i.e. conservative plate boundary-where plates slide past each other along transfomi faults, characterizes them as geological structures that are presently active [START_REF] Rehman | Saudi Arabian geothermal energy resources an update[END_REF]. This suggests that local movements of parts of the crust, probably affect areas of great extent like Hotei Regio and Harrat Khaybar, even if they are not located precisely in the cerner of the active area. In the case of Harrat Khaybar, the adjacent Red Sea fault continues to propagate: its rifting causes seafloor spreading, triggering the volcanic centers of the région [START_REF] Camp | The Arabian continental atkali basait province: Part D. Evolution of Harrats Khaybar, Ithnayn, and Kura[END_REF]. Such a process illustrâtes the relation between the terrestrial volcanic terrains and tecton ism, a relation that seems possible for Titan's case as well.

Tui Regio, Hotei Regio and Sotra Facula are ail located in the 15°S-30°S latitudinal zone which is close to the Southern margin of Xanadu (Soderblom et al., 2009), implying that the région might be an extensive zone of crustal weakness. The existence of possible volcanic and tectonic features within a spécifie area seem to be manifestations of the most active région of Titan like the boundaries of tectonic plates on Earth. Although still under investigation in Titan's case, the definite identification and understanding of morphotectonic features in these régions is crucial in order to détermine the presence and origin of zones of crustal weakness, which will in turn impose additional constraints on cryovolcanism on Titan, lndeed, on Earth, as well as on other pianetary bodies, the interplay of volcanism and tectonism causes the formation of extensive and distinct geological terrains.

The moons of Jupiter, Europa and Canymede, possibly also display similarities with the morphotectonics of Earth (volcanism-tectonism)

and of the Saturnian icy satellites (cxyovolcanism-tectonism) Fig. 7a shows Europa's ridges, junctions and dômes, as seen by the Galileo instrument Solid State Imager (SSI) which are typical geological expressions on this moon. The ridges are most likely formed by cryovolcanic processes probably causing déposition of subsurface materials over surface units, accompanied by tectonic movements, that formed the linéaments (Figueredo and Greeley, 2004) This extensionai volcano-tectonic mechanism is similar to the terrestrial mid-oceanic rifting (Prockter et al., 2002). On Ganymede, volcanic processes may hâve occurred in the past, but current evidence suggests the presence of tectonic processes (e.g, [START_REF] Head | Evidence for Europa-like tectonic resurfacing styles on Ganymede[END_REF].

The linked tectonic-cryovolcanic hypothesis suggests graben rupture (depressed block of ice bordered by parallel faults) due to lithospheric extension and cryovolcanic déposition caused by flooding of the internai material (cryovolcanic resurfacing) (e.g., [START_REF] Schenk | Dark halo craters and the thickness of grooved terrain on Ganymede[END_REF][START_REF] Murchie | Terrain types and localscale stratigraphy of grooved terrain on Ganymede[END_REF]. Also fault blocks operating as zones of weakness can act as pathways for the bright material onto the surface while the caistal movements could produce the bright ridges (e.g, [START_REF] Head | Tectonic resurfacing on Ganymede and its rôle in the formation of the grooved terrain[END_REF][START_REF] Head | Evidence for Europa-like tectonic resurfacing styles on Ganymede[END_REF] (Fig. 7b).

Ganymede, and partially Europa, are the targets of a future mission proposai to thejovian System, which will, among others, could explore the icy moons' interior heat potential, as well as, the surface motion and morphology, with improved and enhanced instrumentation in order to better understand their surface composition, internai structure, dynamics and the mor photectonics. Further investigation and comparison of the mor photectonic features of many of the icy satellites will shed light on the potential different internai mechanisms that operate in the Solar System.

Discussion

The morphotectonic structures presented here are related to the most elevated, as well as the most fractured features observed on Titan. Major mountainous régions are concentrated in mid-latitudes between 30°S and 3CFN and probable cryovolcanic areas are located within the same zone (20"S-30CS) Linear features are displayed ako within the same région (10'S-26°S). Fig. 8 represents a location map of the major morphotectonic features presented in this study.

We hâve argtied here that ail these features are related to surface stress (ïelds. In analogy with terrestrial morphotectonic structures, the shape, size and morphology of Titan's observed mountains, ridges, hills and linear features such as faults, major fractures and canyons probably originale through some form If in future missions a number of Titan's surface features are definitely identified as a resuit of compressional processes, as it has been proposed here, and despite the fact tliat large stresses are required to form compressional features (Pappalardo and Davis, 2007), then this will render Titan unique among the rest of icy satellites where extensional features are dominant [START_REF] Jaumann | Icy satellites: geological évolution and surface processes[END_REF] and will ratify the thesis of Mitri et al.

(2010) that some of Titan's mountains represent folds and/or thrusts.

On the other hand, in regard to the extensional features on Titan as compared to their terrestrial analogs one may question if there are some essential différences in the development and propagation of fractures in the icy crusts vis-à-vis the silicate crusts.

However, before addressing these questions, one should initially examine which factors control tectonism on a planetary body; essential stress mechanisms that can be either global, régional and local. Global stress mechanisms include tides, nonsynchronous planetary rotation, polar wander, despinning, orbital recession and radiogenic decay. One major stress mechanism is provided by convection. Another array of mechanisms is due to volume changes up to the large density contrast between ice-l and water and is applicable to icy satellites (Collins et al., 2009).

Finally, "impacts" represent another local stress mechanism. To global, régional or local stresses, Solar System bodies and particularly their crusts may react in a brittle, ductile or more rarely elastic fashion, producing corresponding landforms. In this respect it is essential to compare and contrast the mechanical properties of icy and silicate crusts, as well as, the order in which stress mechanisms occur, to lind corrélations between the morphotectonic features on Earth and those on Titan (Radebaugh et al., 2007;Collins et al., 2009).

Thus, a question arises: is there some essential différence in the development and propagation of faults in the icy crust with respect to a silicate cnist? As indicated in Table 1, ice and silicates mainly share a similar crystal structure, differ in melting tem pérature but when ice involves water and methane, an additional similarity with the silicates is found in the three physical States of the material (solid, liquid, gas).

Earth and Titan share ail the global stress mechanisms and sources of internai heat such as radiogenic decay, heating by applied tidal forces and primordial heat.

However the analogy probably stops here since Earth has preserved a capital of its primordial heat. The outer core of the Earth with a thickness of 2890-5150 km has a température range of 4400-6100 °C, similar to the photosphère of the Sun. Furthermore the total heat flux from Earth's interior ranges from 0.08 to 0.4 W m~2 [START_REF] Pollack | Heat flow from the Earth's interior: analysis of the global data set[END_REF]Davies and Davies, 2010). Due to the State (liquid) and composition (Fe-Ni) of Earth's outer core, thermal tvnaways occur at the core-mantle interface and resur face as hot spots, causing local mantle convection and induce plate motion which implies the breaking of continents (Africa) or supercontinents (Pangea).

Hot spot volcanism is found in the middle of lithospheric plates and on the margins of extensional plates. Volcanism also occurs in compressional plate margins followed by orogenesis (mountain building). Titan's primordial heat flux front the satellite's interior is of the order of 0.02-0.06 W m"2 and Mitri et al.

(2010) proposed that it can similarly resuit in crustal fold processes and mountain building.

On Titan, stress mechanisms even if global, as in the case of tides, can hâve very local effects: for instance Saturnian tidal forces resuit in a concentration of morphotectonic structures mainly around Titan's equator. It is not a random event that ail the three candidates for cryovolcanic areas are concentrated in Titan's équatorial zone. Actually, it will be surprising if future detailed observations réfuté the case for cryovolcanism on Titan in this area, given positive indications from current studies (Lopes et al., 2010;Solomonidou et al., in préparation). Morphotectonic against each otlier (compression) forming mountain-or ridge-like features. Subséquent extension of the so-formed fractures will provide the pathways for cryovolcanism, which will be the younger event.

Primordial heat, as well as, heat produced by radiogenic decay and heat induced by tidal forces can be dissipated both by conduction and/or convection. It has been shown that transitions from a conductive to a convective State for an ice shell, overlying a pure liquid océan (Tobie et aL, 2005), can hâve major effects on surface morphoteetonics (Mitri and Showman, 2008). The same authors hâve shown that thermal convection can occur under a range of conditions in the ice-I shells of Titan and two possible scénarios can follow. A thin ice with a low viscosity base ( 1) and a thick ice with a high viscosity base (11). Ihus, Mitri and Showman (2008) [START_REF] Fowler | Thermal runaways in the Earth's mantle[END_REF][START_REF] Baumgardner | Runaway subduction as the driving mechanism for the Genesis Flood[END_REF][START_REF] Baumgardner | Catastrophic plate tectonics: the physics behind the Genesis Flood[END_REF][START_REF] Benn | Archean geodynamics and environments[END_REF].

Thermal stresses are responsible for updoming, weakening and subséquent fracturing of the crust of a planetary body. On Earth, such a paradigm is provided by the continent of Africa where hot spot activity has resulted in updoming, fracturing and volcanism (East African Rift, Ahaggar Mountains, Table 2a). On Titan, zones of tectonic weakness hâve probably formed in an analogous manner i.e. as a conséquence of thermal stresses and weakening of the crust with concomitant formation of open fissures, which act as pathways for the éjection of material from the interior (cryovolcanism) as it has already been proposed for Enceladus and Europa [START_REF] Manga | Pressurized océans and the éruption of liquid water on Europa and Enceladus[END_REF].

Localized compression and crustal thickening can also lead to mountain building on Titan. Linear mountains in T3 and T8 swath, seen in radar data, probably hâve formed this way (Fig. 3a andb; Table 2b). A terrestrial analog can be provided by the Laramide Orogeny. The Laramide event that affected Western-North Amer ica during Late Cretaceous and Early Paleogene time, involved compressive forces, conductive heating and crustal thickening that eventually led to mountain building (orogeny) [START_REF] Dickinson | Plate Tectonics of the Laramide Orogeny[END_REF]. A similar exampie is provided by the Acadian Orogeny during rniddle Paleozoic, which has affected the Eastem-North America as the resuit of collision of Baltica plate with the Laurentia.

The proposed formation mechanism for Titan's mountains, mainly the ones observed during T8, T30 and T43 fiybys (Fig. 2c), concerns folding of the upper crust due to high heat flows from the interior and high température gradients in the past (Table 2b). A terrestrial analog, the Rocky Mountains chain in Western-North America (Fig 3), was formed over an area where high values of mantle heat fiow occur [START_REF] Bird | Formation of the Rocky Mountains. Western United States: a continuum computer model[END_REF].

For the high-standing ranges of the T20 radar swath a forma tion mechanism has been proposed by [START_REF] Sotin | Cassini/VIMS observatoire of Titan during the T20 flyby[END_REF] involving the génération of stresses by tectonic extension that penetrate the icy shell and create pathways for the internai material. This finds a terrestrial analog in the formation of the Mid Atlantic Ridge in Atlantic Océan where divergent tectonic plates are associated with magma upwelling due to partial melting of the upper mantle in the interior.

Other surface features on Titan that seem to be the resuit of intense forces of tectonic extension are the probable blocks and grabens seen in T8 radar swath (Tobie et al., 2006;Radebaugh et al., 2007). On Earth, the Basin and Range Province in USA (Fig 4;

Table 2b) is such a large terrain subjected to forces of tectonic extension As it can be seen from the rim create hilly features [START_REF] Ramsey | Ejecta distribution patterns at Meteor Crater, Arizona: on the applicability of lithologie end-member deconvolution for spaceborne thermal infrared data of Earth and Mars[END_REF]. Furthermore, two terrestrial analogs presented and described in this study that seem to hâve a major resemblance with canyon-like and cryovolcanic features seen on Titan, are the Grand Canyon (Fig 5) and the Flarrat Khaybar (Fig 6). Both terrestrial terrains consist of a number of geological features that hâve been subject to multiple processes.

Titan's surface seems to be as complex as the Earth's and may be the end resuit of the action of multiple processes as well.

The fact that the mountains and ridges seem to be concentrated around Titan's équatorial band is a very strong argument for tidal forces shaping these morphoteetonics features. However, Lopes et al. (2010) showed that mountains exist at other latitudes, a possible sign for a global source of stress as well. It is plausible that radiogenic decay would heat uniformly the entire satellite,

similarly to what occurs on Earth (Dickin, 1995), and lead to morphotectonic structures evenly distributed ali over the body's crust. We argue, nevertheless, that mountain ranges are more extensive at the equator of the satellite as a resuit of tidal forces since the very geometry of tides causes local stresses larger at the equator than anywhere else, inducing convection and therefore enhancing mountain formation. This mountain formation can be contrasted with that on Earth where mountains occur globally along collisional plate boundaries. Furthermore, as noted hereabove, radiogenic heat production on small bodies like Titan or even Mars has significantly decreased over a long period of time (e.g., [START_REF] Schubert | Thermal historiés, compositions and internai structures of the moons of the solar System[END_REF]Grasset andSotin, 1996: Grott and[START_REF] Croit | Constraints on the radiogenic heat production rate in the Martian interior from viscous relaxation of trustai thickness variations[END_REF]. It should be noted here that Saturn's tides are still heating up Enceladus' interior (Uurford et al., 2007).

The existence of tectonism on Titan can provide significant insights on the internai stmeture of the Saturnian satellite. In reference to the terrestrial paradigm, where rigid lithospheric plates 'float on a weaker asthenosphere, it could provide indirect evidence for the existence of a subsurface océan on Titan. The importance of deciphering morphotectonic features on Titan that can be linked to tectonism has also conséquences in elucidating the methane cycle on Titan in analogy with the link between terrestrial tectonics and the global terrestrial carbon cycle [START_REF] Bolin | Carbon Cycle Modelling[END_REF][START_REF] Ruddiman | Tectonic Uplift and Climate Change[END_REF]. More specifically, Titan's tectonic activity can probably be directly related to the replenishment of the atmosphère in methane (Sotin et aln 2005;Atreya et aL, 2006;Tobie et aL, 2006;Coustenis and Taylor, 2008;[START_REF] Lunine | The methane cycle on Titan[END_REF]Mitri et al., 2008).

Finally, the origin of some morphological features can be attributed to exogenous processes such as meteorite impacts (Moore and Pappalardo,201 1), aeolian processes (Slofan et al., 2006;[START_REF] Radebaugh | Linear dunes on Titan and earth: Ini tial remote sensing comparisons[END_REF] and Jluvial érosion (Forco et al., 2005;[START_REF] Tomasko | Rain, winds and haze during the Huygens probe's descent to Titan's surface[END_REF], especially monsoonal rainfall causing flooding mainly in the équatorial région [START_REF] Tokano | Précipitation dimatology on Titaa[END_REF].

The Cassini-Huygens mission has significantly improved our understanding of Titan and the coupling of its atmosphère and surface. Since 2004 and for another five years, Titan is investigated by Cassini's fiybys. However, only limiied surface coverage will be achieved at high spatial resolution. Therefore, the composition and évolution of its diverse surface features will still demand extensive future investigation.

INTRODUCTION

Icy moons are small celestial bodies whose surfaces are par tially, if not principally, covered by ice, inostly water ice (Johnson, 2004). Tlte most remarkable icy moons arouttd the giant planets are Jupiter's Ganymede, Neptune's Triton, Uranus's Miranda and Satum's Titan and Enceladus in a variety of sizes, composition and températures. It was thought that due to the abundance of water ice, the large distance front the Sun, the absence of internai energy sources and of an at mosphère in most cases, the geology of these bodies would he simple, or rather simpler thaii the geology of Earth; however, subséquent images hâve shown complex surfaces with several notable morphological formations. Furthermore, the composition, as well as, the structure of tlte surfaces of the icy moons dépends on geological and geophysical factors (Johnson, 2004). far oui of the habitable zone, it is one of the most likely worlds in our solar systein of astrobiological interest (Raulin, 2008). Except for the new atmospheric discoverles such as the organic chemistry in the ionosphère, new components in tire neutral atmosphère and the properties of the troposphère, Cassini-Huygens' most surprising discovery was Titan's complex and Earth-likc geology (Cûustenxs & Taylor, 2008). As far as the surface is concerned, one of the moon'sexceptional characteristics is the existence of surface liquid bodies that resemble tcrrestrial lakes (Stofan et al., 2007). Otlier surface formations, were captured both by the Cassini orbiter's remore settsing instrumentation such as tire Syntlretic Aperture Radar (SAR) (Elachi et al., 2005); the Visual and Infrared Mapping Spectrometer (V1MS) [START_REF] Brown | The Cassini Visual And Infrared Mapping Spectrometer (Vims) Investigation[END_REF] and the Imaging Science System (ISS) (Porco et al., 200*1;[START_REF] Mccord | Composition of Titan's surface from Cassini VIMS[END_REF], as well as by the Huygens probe's in situ instruments i.e.: the Surface Science Package (SSP) [START_REF] Zarnecki | A soft solid surface on Titan as revealed by the Huygens Surface Science Package[END_REF], the Descent Imager and Spectral Radionretcr (D1SR) [START_REF] Tomasko | Rain, winds and haze during the Huygens probe's descent to Titan's surface[END_REF] and the Gas Chromatograph Mass Spectrometer (GCMS) (Fulchignoni et al., 2005). The surface discoveries include extensive mountains, ridges, dendritic networks, dunes, lakes, channels, canyons and riverbeds. Of even higlrer im portance is the possible existence of active zones on the satel lite due to past or recent cryovolcanic and tectonic activity (e.g. Soderblom et al., 2007;Lorenz et al., 2008;Solomonidou et al., 2010). Caldera-like édifices characterized by radial faults, features resembling lava flows and other possible volcanic structures and deposits, witlrin large areas of volcanic-like terrains, in addition to spectral data indica tions, suggest that Titan is a world tirât once suffered cryo volcanic activity which could possibly still be active. The suggestion of an active cryovolcanic interior that supplies the atmosphère with nrethane is compatible withthe current level of methane in Titan's atmosphère. Accordingto calculations, the lifetinre of atmospheric methane is limited to 10-100

Myrs (Wilson et al., 2004). If we assume tirât methane in the atmosphère should be replenished, tlren Titan needs a réservoir that would supply the atmosphère with enough methane to maintain the atmospheric abundance. The requirement of sufficient supplies of methane in combination with the volcanic-like expressions did trigger the theory of active cryovolcanism on Titan (Tobie et al., 2006).

Other than Titan, the Cassini mission unveiled another unique world among Saturn's icy nroons. Enceladus is a significantly smaller satellite than Titan (500 km in dianreter), it présents however, extremely interesting surface features including cratered as well as smooth terrains, extensive linear cracks, scarps, troughs, belts of grooves in addition to the spectacular phenomenon of volcanic geysers that Cassini in strumentation captured in 2005 in the south pôle (Porco et al., 2006). High-resolution data frora Cassini inagnetometer (MAG) (Doügherty et al., 2006), ISS (Porco et al., 2006) and the Ultraviolet Imaging Spectrograph (UVIS) [START_REF] Hansen | Enceladus' Water Vapor Plume[END_REF], reported cry ovolcanic activity in the form of jets in the Southern Polar région, at the geological surface ex pressions called "Tiger stripes''. Tire accumulation of multi ple jets resulted in the formation of a massive fountain tirât reached over 435 km in height (Porco et al,, 2006).

Our work provides; i/ an OverView of the geology of Titan and Enceladus, ii/ terrestrial analogues and iii/ the results of our data analysis regarding Titan 's potentially active régions.

This study implicates the presence of cryovolcanism on Titan's surface.

GEOLOGY AND CRYOVOLCANISM ON TITAN &

ENCELADUS

Cryovolcanism

Cryovolcanism Ls considered to be one of the principal geo logical processes that hâve shaped several of the icy nroons'

surfaces. This activity can be described as ice-rich volcanism, while the cryovolcanic ejecta are refereed to as cryomagma. The ctyonragma appears in tire form of icy cold liquid and, in sonre cases, as partially crystallised slurty (Kargue, 1994). The possibility of volcanic resurfacing on icy nroons was f irst noted by [START_REF] Lewis | Satellites of the Outer Planets: Their Physical and Chemical Nature[END_REF][START_REF] Lewis | Metal/silicate fractionation in the Solar System[END_REF] and sulrsequently addressed by Consolmagno & Ijîwis (1978), but it was not until after the Voyager flybys of Jupiter and Saturn tirât evidence for past and présent tectonic and volcanic ac tivity on inoons such as Europa, Ganymede, and Enceladus was brought to light. In our Solar System the only observed recent éruptions are limited toEartlr and three otlrer locations:

1) lo, moon of Jupiter; 2) Triton, moon of Neptune; and, 3)

Enceladus, moon of Saturn. Titan is also major candidate for past and/or présent cryovolcanic activity awaiting for a de finitive evidence.

Subséquent to the Cassini-Huygens finditigs, the term 'cryovolcanism' has been associated with Titan more than any other Satumian moon (Soderblom et al., 2009). Even though, for the case of Enceladus the ciyovolcanic origin of the plume is now confirmed (Porco et al., 2006), the cryo volcanic activity on Titan présents a controversial scientific issue witlrin the scientific community. However, sonie facts are in favor of such processes like the theory of cryonragma being relevant to the formation of prebiotic compounds (e.g. Fortes, 2000).

The composition of the material called cryomagma on Titan's surface is still unknown, due to the lack of in situ measurements and in depth investigations, whichmay reveal its properties. Cryovolcanic features on Titan's surface are believed to be a significant source of tire methane présent in the atmosphère (Lorenz & Moton, 2008). Considering this, a model has been suggested regarding the évolution of Titan, indicating that the methane supply may be trapped in a methane-rich ice and episodically released by cryovolcanic phenomena (Tobie et al., 2006). However, tire definite answer of the composition of Titan's cryomagma is still a subject of research.

According to Tobie et al. (2006) the methane could hâve where subséquent cooling and erystallization of the outer layer occurred. Fortes et al. (2007) suggested that ammo nium sulphate is the possible origin of cryovolcanism. Titan's interior is broadly described by these authors, from tlie core to the crust, in distinct layers: a serpentinite core, a high-pressure ice VI mantle, where ice VI has a differentiated crys talline structure ordering and density than typical water ice, a liquid layer of aqueous ammonium sulphate and an externally heterogeneous shell of methane clathrate with low-pressure ice Ih (similarly as ice VI) and solid ammonium sulphate (Fortes et al., 2007).

The geological map below (Fig. 1) has been derived from albedo and texture variations and indicates that the circular feature shows signs of several sériés of flows, as shown by the black lines (Sotin et al., 2005). Tlie black circle indicates a caldera, similar to vents that appear above réservoirs of molten material associated with volcanoes on Earth. The colours of the map are représentative of the brightness of the features where yellow-green to light brown are the bright patches; blue are the dark patches, red the mottled material and finally tlie yellow area marks the location of the volcano (Sotin et al., 2005).

Enceladus

Cassini's observations on Enceladus did reveal distinct geo logical features. The surface of Enceladus is covered by smooth and cratered terrains, ridges, grooves, escarpments and extensive linear fractures (Johnson, 2004). The most interesting and youthful terrain seen on this moon called "Tiger

Stripes" and présents a very complex structure and évolu tion. The Tiger Stripes (Fig. 2) are tectonic structures consisting of four sub-parallel, linear dépréssions located in the south polar région (Porco et al., 2006). In 2005 Cassini's in strumentation and especially the ISS experiment provided evidence of active cryovolcanism (Fig. 3), emanating from a sériés of jets located within the Tiger Stripes (Porco et al., 2006).

The jets of water ice front the fractures of Tiger Stripes produce a plume of gas and particles like N Ht, Na, K salts (Waite et al., 2009). These tectonic fractures, discharge ma terial by endogenic dynamic and most probably hydrothermal activity. The source of the jets is a controversial issue as ex tensive internai stratification as well as dynamic modeling, is required for the source to be identifîed. The recent discovcry of salts in Saturn's E-ring composition, which is fed front Enceladus' plumes (Postberg et al.. 2009), suggests that the source of jets is possibly a "chamber" of liquid water that lies undemeath the ice shell (Tobie et al.. 2008). Altematively, the material could dérivé front originally wann ice that is heated and explodes by the dissociation of clathrate hy-Fig. 1. Geological map of Titan possible volcano (Sotin étal., 2005). Anezina Sohmonidvu, Géorgien Bampasidis, Konstanlinos Kyndkojmlos, Emmanuel Bratsvlis, Mathieu Hirtzig. Athéna Coustenis <£ Xcnvphon Moussas drates [START_REF] Porco | Cassini Observes the Active South Pôle of Enceladus[END_REF]. The clathrate hydrates are crys talline water-based ices where the host ntolecule is water and the trapped-guest ntolecule is typically a gas. The VIMS in strument detected simple organic conipounds in the Tiger Stripes. Sucli Chemical composition which consists of liquid water, ammonia, carbon dioxide,Na and K saJts, benzene and other hydrocarbons (Waite étal., 2009), has not been found in any other région on Enceladus (Brown et al., 2006). The presence of liquid water might also make it possible for Ence ladus to support life (Lammer et al., 2009).

Recent data front Cassini reported pockeLs of heat that appear along a fracture named Baghdad Sulcus (Fig. 15),one of the Tiger Stripes that erupt with jets of water vapor and ice particles [START_REF] Hurford | Gcological impli cations of a physical libration on Enceladus[END_REF]. Tlie température along Baghdad Sulcus exceeds 180 Kelvin (Waite et al., 2009). As is the case for Titan's Hotei Regio, Tiger Stripes on Ence ladus and in particular Baghdad Sulcus represent tectonic zones of weakness from which the internai materials find their way to the surface. The idea of a subsurface sea becontes ail the more contpelling since Enceladus' south polar région (Tiger Stripes area) is actually a half-kilometer deep basin distinguishing from the surrounding expressions (Collins & Goodman, 2007). Such figure, like the deep basin in Tiger stripes, resembles Titan's Hotei Regio which is a basin lying one kilometer deeper titan the surrounding area (SODERBLOM et al., 2009). This basin could be the surface ex pression of a subsurface sea (Collins & Goodman, 2007). Titan Titan's geology has been extensively studied using Cassini image data. In this researeh, we invesrigate and process data acquired from VIMS in order to identify areas of cryovolcanic déposition.

The most intriguing problem in regard to the decoding of Titan's surface is the atmospheric veil that covers the surface. et al., 2005). In general. Titan's surface appears to hâve smooth and rough areas of various altitudes which include extensive mountains and ridges (Fig. 4a) (Lofes et al., 2010), longitudinal dunes (Fig. 4c) [START_REF] Radebaugh | Linear dunes on Titan and earth: Ini tial remote sensing comparisons[END_REF], dendritic networks (Fig. 4b) (Lorenz et al., 2008), liquid lakes (Fig. 6) (Stofan et al., 2007) and impact craters that are intermittently filled by attnospheric précipitations (Elachi et al., 2005). Radebaugh et al. (2008) suggests that mounrains on Titan range front 200 m to 2000 m in height (Fig. 4a). Erosional processes that operate at the area where mountains lie, are probably the reason of the significatif short height of the mountains. However, there is also the assumption that the mountains are built by material with properties that prevent the altitudinal growth (Radebaugh étal., 2008). Radebaugh et al. (2007) mentioned tliat the notably SAR bright features on Titan's surface niost probably correspond to mountains and tectonic ridges which represent mountain chains (Fig. 4a). In particular, the tectonic ridges could hâve suffered atmospheric précipitation (i.e. hydrocarbon rain) acquiring a rough and fractured surface (Soderbi.om et al., 2007). Rivers are coramon on Titan, while in some cases a few craters are traversed by them [START_REF] Wood | and the Cassini Radar Team[END_REF]. The observation of river Systems with dendritic patterns (Fig. 4b) (Lorenz et al., 2008), in addition to the observation of storm clouds (Porco et al., 2005), suggest that rainfall may be a continuing erosional force erasing impact craters. Other surficial structures observed on Titan are impact craters. Of particular impor tance is the small number of impact craters which has been observed by the Cassini/Radar which suggests that the sur face of Titan is relatively young (i.e. [START_REF] Wall | Cassini RADAR images at Hotei Arcus and western Xanadu, Titan: evidence for geologically récent cryovolcanic activity[END_REF]. [START_REF] Radebaugh | Linear dunes on Titan and earth: Ini tial remote sensing comparisons[END_REF].

Moderately variable winds that either follow one mean direction or alternate between two different direc tions hâve formed the observed longitudinal dunes (Lorenz et al., 2006).

Our knowledge regarding Titan's surface deposits is limited to the data acquired from Huygens' landing site. The Huygens captured image was tliat of a dark plain covered in pebbles mainly composed of water ice (Fig. 5) [START_REF] Tomasko | Rain, winds and haze during the Huygens probe's descent to Titan's surface[END_REF]. The size of pebbles is estimated to be roughly 10-15 cm. There is evidence of érosion at the base of the icy rocks, indicating possible fluvial activity [START_REF] Tomasko | Rain, winds and haze during the Huygens probe's descent to Titan's surface[END_REF].

The surface is darker than originally expected, consisting of a mixture of water and hydrocarbon ice. It is believed that 22,200622, (Stofan et al., 2007a)).

Anczina Salomonidou, Geurgios Bampasidis, Konsumtinos Kyriakopoulos. Emmanuel Braisolis, Mathieu Hirtzig, Athéna Coustenis & Xetutphon Moussas the visible ground "powder" in the image is possibly précip itation from the hydrocarbon haze above [START_REF] Tomasko | Rain, winds and haze during the Huygens probe's descent to Titan's surface[END_REF].

One of tlie moon's exceptional characteristics is the exis tence of large liquid bodies described as lakes of surface liq uida (Fig. 6) (Stofan et al., 2007). These feutures resemble terrestrial lakes consitute a unique characteristic displayed by the icy moons. Based on data provided by the Cassini/Radar, the presence of hydrocarbon lakes on Titan's surface is now well established (Fig. 6) (Lofes et al., 2007a).

Candidates of cryovolcanic areas on Titan

Our study involves l\vo major areas on Titan that are the most significant, as well as, interesting cryovolcanic candidates.

These areas are Tui Regio and Hotei Regio (Fig. 7) lying Fig. 7. The location of Tui Regio and Hotei Regio on Titan's globe (Barnes et al., 2006).

Distinct areas on Tui Regio

Wavelength (nm) Plot 1. Spectral plot of five distinct areas at the seven almospheric spec tral Windows.

within the bright région ofXanadu (100°N, 15°S).Tui Regio is centered at 130°VV. 20°S and présents relatively high 5 pin reflectivity. Its size is 1,500 km long and 150 km wide. This bright area lias been identified as a surface feature and not as the image capture of fog, due to the area's spectral behavior at 2.7 mm (McCord et al., 2006). Tui Regio is a massive flow-like terrain, which resembles flow field volcanic areas on Earth. Another area whose spectrum matches that of Tui

Regio is Xanadu's Hotei Regio. Hotei Regio is centered at 78°W, 26°S and comprises a 700 km wide field that is probably volcanic in origin. VIMS images confirm the interpré tation that the area is a low basin surrounded by higher terrains with possible calderas, fault structures and extensive cryovolcanic flows (Soderblom et al., 2009).

Method and Data analysis

Both Tui Regio and Hotei Regio are suggested to be geologically young due to the fact that both présent anomalously bright and spectrally distinct areas that hâve not changed from seasonal précipitation (Barnes et al., 2006).

In order to investigate geologically tire régions of interest, it is essential to study their Chemical composition that lead to the aforementioned brightness as well as their rnorphology in order to dérivé the geological factors that led to their for mation.

We hâve processed spectral images acquired from VIMS, for both areas in the seven narrow spectral Windows centered at 0.93, 1.08, 1.27, 1.59, 2.03, 2.8 and 5 pm for which absorption by atmospheric methane is minimal.

The main goal is to identify the composition as well as the alterations of the components that compose the possible cryovolcanic structures. We hâve used the principal compo nents (PCs) of the Principal Component Analysis (PCA)

method. The PCA method involves a mathematical proce dure that transforms a number of possibly correlated vari ables into a smaller number of uncorrelated ones called principal components (Jollifff., 2002). The PCA is well adapted to our study. as our primary concem is to détermine the minimum number of factors that will account for the maximum variance of the data we use in this particular multivariate analysis. The main goal of PCA is to reduce the dimensionality of a data set consisting of a large number of interrelated variables, while retaining as much as possible of the variation présent in the data set. This is achieved by transforming, the principal components (PCs) into a new set of variables, which are uncorrelated, and which are ordered so that the first few retain most of the variation présent in ail of the original variables (Joluffe, 2002).

PCA images for both Tui Regio and Hotei Regio allowed us to isolate areas with distinct and diverse false coloring, which imply areas of distinct and diverse spectral and Chem ical composition (Fig,9;11). Such Chemical diversity suggests that endogenic and/or exogenic geodynamic processes hâve forrned these régions.

Hellenic Journal of Geascietices,vol. 45, Fig

RESULTS

Tui Regio

We hâve isolated five distinct areas (Fig. 9) within Tui Regio.

The Principal Component Analysis projections (Fig. 8c,d on the satellite normalized for Titan [START_REF] Thekaekara | The Extraterrestrial Solar Spectrum[END_REF]Barnes et al., 2007;[START_REF] Brown | The Cassini Visual And Infrared Mapping Spectrometer (Vims) Investigation[END_REF]).

The plot (Plot 1) of "Bright cryolava field'" terrain is dif ferent from the other plots, presenting higher I/F values. This suggests that, additionally, tliis area is extremely brighter than the resr of the région. The wavelengths at which this area présents obvious alterations are the 2 pm, 2.8 pm and 5 pm.

Hotei Regio

We hâve also isolated five distinct areas within Hotei Regio's probable volcanic fietd (Fig. 11 ). The PCA projections (Fig. Hotei Regio's spectral graph (Plot 2) indicates that the "Bright Cryovolcanic area" présents the highest I/F values and remains brighter than the other areas at ail wavelengths.

In addition, the "Dark area" remains darker with low values of I/F at ail wavelengths. Surprisingly, the spectra from caldera-like structures présent medium I/F values, lying almost in the average between the brighter ("volcanic area") and the darker ("primai surface") at most wavelengths. This is compatible with terrestrial caldera structures that consist partially of primary surficial components on which the volcano is being built, as well as new material coming from the [START_REF] Radebaugh | Linear dunes on Titan and earth: Ini tial remote sensing comparisons[END_REF].

Earth-like body (Coustenis & Taylor, 2008). In addition, many atmospheric aspects such as the cHmate and the meteorology, as Titan displays a 'methanologic al' weather cycle of clouds, rainfall and évaporation that parailels the Tiydroiogical' cycle of the Earth, as well as, its complex morphology, make Titan an extremely important astrobiological place. Specifîcally, cryovolcarrisniha s important a strobiological im plications, as it provides a mechanism to expose Titan's organics to liquid water, transfôrming hydrocarbûns and nitriles into more evolved and exidized prebiotic species [START_REF] Neish | The potential forprebiotic chemistry in the possible cryovolcanic dôme Ganesa Macula on Titan[END_REF]. Also it has been suggested that life could exist in the lakes of liquid methane on Titan (McKay & Smith, 2005). The existence of liquid bodies identified as lakes exposed on the surface (Stofan et al., 2007), the équatorial dunes (Fig. 4c), dendritic flows, potential tectûnics and volcanism, enhance Titan's resemblancetoour ownplanet. Prior to this discovery, such combination of surface features and dense nitrogen atmosphère had only been identifîed on Earth.

Ail the aforementioned aspects, which mainly are the nitrogenic atmosphère, the liquid lakes, as well as the Earthlike geological structures, suggest that Titan resembles Earth more than any other body in the Solar System; despite the httge différence s in température and other environmental con ditions . Thus, an holistic understanding of Titan's System will help us better under stand Earth's évolution starting with its primordial phase sinceearly Earth probably looked much like Titan looks today [START_REF] Owen | Planetary sdenoe: Huygens rediscovers Titan[END_REF]. In general, the activity of cryovolcanism might operate in analogy to terrestrial hydrûvolcanic éruptions (Solcmonidou et al., 2010). Fig. shows the supervolcano of Lake Toba in North Sumatra, Indonesia, which is 100 km long and 30 km wide, in comparison with the possible supervolcanic structure called Ganesa Macula (50°N, 87°W) [START_REF] Bames | Near-infrared spectral mapping of Titan's mountains and channels[END_REF]. The Volcanic Explosivity Index (VEI) for Lake Toba, which provides a rel ative measure of the explosiveness of volcanic éruptions high named Littlc Black Peak. The peak consists of three nested ciader cônes and a solidified lava pond. Tiic Titanian analogue (Fig. 14) is the possible cryovolcanic flow, Tui

Regio; one of the largest seen on Titan. It présents similar shape as Carrizozo flow fleld (Fig. 14).

As indicated above, Hotei Regio is probably a massive cryovolcanic terrain that consists of caldera-like figures and depositional areas fillcd with lava flows. One terre striai ter rain that rcscmbles Hotei Regio is the Harrat Khaybar vol canic field (Fig. 15), which is located at Nortli of Médina in Saudi Arabia. The western half of the Arabian Peninsula contains extensive lava fields known as haraat [START_REF] Pint | Prospects lorLava-Cave Studies in Harral Khaybar, Saudi Arabia. A5IC[END_REF]. One such fleld is the 14,000 km2 volcanic field that was forrned by éruptions along a 100 km N-S vent System over the past 5 million years. The area contains a wide range of volcanic rock types, spectacular landforms and several générations of dark fluid basait lava flows [START_REF] Pint | Prospects lorLava-Cave Studies in Harral Khaybar, Saudi Arabia. A5IC[END_REF]. Jabal Abyad, in the center of the image, was forrned from more viscous, silicarich lava classified as a rhyolite. While the 322 m high Jabal al Qidr exhibits stratovolcano, Jabal Abyad is a lava dôme. To the west. (image top center) is the impressive Jabal Bayda.

This symmetric structure is a tuff cône, forrned by éruption of lava in the présence of water. The combination produces wet, sticky pyroclastic deposits that can build a steep cône structure, particularly if the deposits consolidate quickly [START_REF] Pint | Prospects lorLava-Cave Studies in Harral Khaybar, Saudi Arabia. A5IC[END_REF]. In Fig. 15, we présent a comparison between Earth's, Titan 's and Enceladus' possible volcanic terrains and tectonic zones of weakness that consist the passage for in Following their strides, astrobiology, the study of evidence for life outside the Earth, is not only the research of the origin, distribution, and évolution of life in the whole universe but also that of structures and processes related to life and its destiny (Raulin, 2007). In general, astrobiology brings together different scientific disciplines such as astrophysics, geology, chemistry, geochemistry, biology, and more in order to shed light on the many aspects regarding the création of our solar System aa well as the initiation of life. Starting with the terrestrial paradigm, astrobiology focuses on extraterrestrial environments, posing the unansvvered ques tion on the origins of life on Earth and elsewhere, while investigating the more easily accessible organic compounds-and in particular the prebiotic chemistry-on other celestial bodies.

It is generally admitted today that life arose on Earth as the dérivative of a long Chemical évolution, implying three major raw ingrédients: liquid wat.er, carbonaceous matter, and energy, working together over time. Since life primitive structures should be able to emerge, evolve, and develop in suitable environments, the quest for possible habitats outside our planet is focused on places where these ingrédients are or hâve been présent.

With the discovery of planets beyond our solar System and the search for current living organisms or for favorable conditions for past and future life in Moreover, Titan orbits around Satum within 16 Earth-days almost synchronously; thus, its solid surface rotâtes very slowly. Instead, due to strong zonal winds (Birdet al., 2005;Lorenz et al., 2008b), its atmosphère is in super-rotation.

Due to Titan's distance from the Sun of about 9.5 astronomical units (AU), the satellite receives slightly more than 1% of the solar flux that the Earth registers at the top of its atmosphère at 1 AU (Fig. 1). In addition, Titan revolves far enough from the giant planet (about 20 Satumian radii), to avoid any critical interactions with the rings or the magnetosphere. Although sometimesTitan moves close enough to Satura to allow its atmosphère to interact with the energetic particles of the magnetosphere of Satura. Together with the solar photons, these interactions play a key rôle in Titan's Chemical évolution. Indeed, Titan possesses an extensive atmosphère made mostly of N2 with a column density ten times that of Earth's (Fig. 2). Furthermore, Titan's astrobiolcgical potential isenhanoedby the presence Both Titan and Earth expérience a greenhouse effect and complex photochemical reactions in their upper atmospheric layers, of a rich organic cheniistry which is produced in its atmosphère, thanks to the presenoe of its second most abundant gas, methane (about 1.4% in the strato sphère and 5% on the ground), and on the surface from the interactions among the various constituents.

Thus, Titan's unique and dense atmosphère harbors a whole host of organic trace gases: hydrocarbous andnitriles (e.g., Coustenis et al., 2007[START_REF] Coustenis | Titan trace gaseous composition from CIRS at the end of the Cassini-Huygens prime mission[END_REF]. Its sur face présents many morphological similaritieswith the Earth's, which surprisingly show similar structural diversity, but the raw materials are different from those on our own planet. With an environment very rich in organics, Titan, along with comets, is thus often considered as one of the best targets to search for prebiotic chemistry at a full planetary scale and a possible habitat for extraterrestrial life in ail probability different from the terrestrial one. More importantly, our understanding of the origin of life on Earth could greatly benefit from studying Titan, where the low solar influx, the composition of the atmosphère, and the possible presenoe of an internai water océan give us the opportunity to study the conditions prevailing on the primitive Earth.

After more than 8 years of close observations by remote sensing and in situ instruments on board the Cassini-Huygens mission, Titan is revealed as an evolving planet, geologically active, not only from erosional processes in the face of the lack of impact craters (Lopes et al., 2010) but also because of its possible cryovolcanism (Tobie et al., 2005;Barnes et al., 2006;Nelson et al., 2009a, b;Soderblom et al., 2009) andmorphotectonism (Solomonidou et al., 2012), aeolian and fluvial érosion (Fulchignoni et al., 2005;Israël et al., 2005;Niemann et al., 2005;[START_REF] Owen | Planetary sdenoe: Huygens rediscovers Titan[END_REF][START_REF] Tomasko | Rain, winds and haze during the Huygens probe's descent to Titan's surface[END_REF]Zaraecki et al., 2005;Jaumann et al., 2008;Lorenz and Mitton, 2008;Raulin. 2008: Lebreton et al., 2009), clouds and pré cipitations, and a methane cycle very similar to the water cycle on Earth (Bird et al., 2005;Atreya et al., 2006;Coustenis and Taylor, 2008;Brown et al., 2009a;Coustenis and Hirtzig, 2009). The organic chemistry products are found on Titan ail the way from the upper atmosphère to the surface and possibly in the interior, indicating close exchanges between the different éléments and planetary layers (interior-surface-atmosphere).

Enceladus, first observed by Sir Frederick William Herschel in 1789, is another intriguing moon of Saturn. Although it is quite small compared to Titan, with a mean radius of 252 km, the large plumes ejected from its south-polar région, as first discovered by the Cassini-Huygens mission magnetometer (Dougherty et al., 2006), make it very important for astrobiology. These geyser-like features mainly consist of water vapor and ice and include organic compounds (Dougherty et al., 2009;Waite et al., 2009). This strongly suggests the potential presence of a complex organic chemistry ongoing in the interior and the presence of liquid water, providing grounds for the search of a liquid océan at short distances under the surface.

Enceladus is an unambiguous example of a cryovolcanically active icy satellite identified in the outer solar System and can be used to understand active processes that. are thought. to hâve once played and/or are still possibly playing (e.g., Titan's case) a rôle in shaping the surfaces of other icy moons. These processes include tidal heating, possible internai convection, cryovolcanism, and ice tectonics, which ail can be studied as they currently happen on Enceladus. Moreover, the plume source région on Enceladus samples a warm, chemically rich, environment that may facilitate complex organic chemistry and biological processes.

What is the level of complexity reached by the organic chemistry in Titan and Enceladus? What is the corrélation between the interior, the surface, and the atmosphère, especially regarding the biological aspects? What are the habitability potentialities in the Kronian environment, in particular for Titan? Such are the questions that further exploration of Titan and Enceladus should answer in the coming years (Brown et al., 2009b). After ail, Darwin's evolutionary ideas could also be successfully applied sensu lato in the Saturnian neighborhood.

Habitability Issues for Outer Planetary Satellites

Liquid water is established as the necessary solvent in which life emerges and evolves.

Water, as an abundant compound in our galaxy, can be found in various places, from cold dense molecular clouds to hot stellar atmosphères (e.g., [START_REF] Cernicharo | Water in space: tlie Water World of ISO[END_REF][START_REF] Hanslmeier | Water in the universe[END_REF]. The thermodynamic behavior of water, which enables it to remain liquid in a large range of températures and pressures and to be a strong polar-non polar solvent, makcs it essential for maintaining stable biomolecular and cellular structures (Des [START_REF] Marais | Remote sensing of planetary properties and biosignatures on extrasolar terrestrial planets[END_REF].

A large number of organisms are capable of living in water. However, in a deposit of pure water, life will probably never spontanéously originate and evolve, since while there are many organisms living in water, none we know of is capable of living on water alone, because life requires other essential éléments such as nitrogen and phosphorus in addition to hydrogen and oxygen. Moreover, no known organism is made entirely of water. "Just water5' is therefore not an auspicious place for the emergence and évolution of life.

The concept of the habitable zone (Fig. 1) is, of course, based on our understanding of life on Earth and is related to the presence of liquid water on a body's surface. But the requirements for extraterrestrial life do not hâve to be the same, suggesting that life could exist outside the habitable zone [START_REF] Cohen | Evolving the Alien: The Science of Extraterrestri al Life[END_REF], in particular if liquid water exists underneath the surface. Furthermore, the changes that occurred on Earth's primordial atmosphère under the influence of early primitive plant life [START_REF] Wolstencroft | Photosynthesis: likelihood of occurrence and possibility of détection on Earth-like planets[END_REF] require that diachronie alter nation be taken into account regarding the habitability opportunities of a planetary body. In addition, internai processes such as volcanic activity, hydrothermal movements, and radioactive decay that possibly occur within satellites located outside the habitable zone could affect the radiation and thermal level of the body and thus change significantly the environmental conditions favoring life (Horneck, 2008).

Of the large satellites of the gas giants, there are those that may house underground water deposits in direct contact with heat. sources below their icy crust and those expected to hâve either liquid water layers encapsulated between two ice layers or liquids above ice. In the study for the emergence of life éléments on such satellites, the timescale is of essence. If it is long enough, the liquid water underground océan may host. life. Thus, the icy satellites of the outer planets of the solar System, as well as the recently discovered exoplanets, host unique condi tions which may inhibit the emergence of life precursors in isolating environments that can prevent the concentration of the ingrédients necessary for life or the proper Chemical inventory for the relevant biochemical reactions. Conversely, according to Trinks et al. (2005), a coupled sea/ice System could in theory provide the necessary conditions for life emergence in the primitive Earth. Additional laboratory experiments and in situ studies of deep subglacial isolated lakes in Antarctica (Kapitsa et al., 1996) would improve our understanding in t.his field, as the physical properties of deep subglacial lakes resemble those found on both Jupiter's moon Europa and Saturn's moon Enceladus (Bulat et al., 2009).

As a conséquence, the satellites of the giant planets like Europa, Enceladus, Ganymede and Titan are possible habitable environments and valid targets in the research for life with space missions and/or télescopés. However, life structures that do not influence the atmosphère of their host planet on a global scale will not be remotely détectable. In the solar system's neighborhood, such potential habitats can only be investigated with appropriate designed space missions, as in the case of Europa, where the Europa Jupiter Space Mission [START_REF] Clark | Europa Jupiter System mission. Joint summary report[END_REF] will attempt to look for the hypothesized internai liquid water océan. In the case of the Saturnian satellites (see hereafter), proposed future missions shall also address this question.

Titan: Organic Factory and Habitat

Retracing the processes, which allowed for the emergence of life on Earth around 4 billion years ago, is a difficult challenge. Our planet has drastically evolved since that time, and most of the traces of what were the initial environmental conditions hâve been erased as a conséquence of plate tectonics and érosion. It is thus crucial for astrobiologists to find extraterrestrial locales with similarities to our planet's early stages. This will provide a way to study in the présent time some of the pro cesses, which occurred on the primitive Earth, when prebiotic chemistry was in its young stages (Fig. 3). For instance, a subsurface océan in the interior of the satel lite of a gas giant (as on Europe, Enceladus, Titan, etc.) may be habitable for some kind of life form-even though not necessarily an Earth analogue-but also infor mation on the terrestrial-like atmospheric and surface conditions on any planetary body can provide valuable information.

Titan is a good candidate in this instance as its atmosphère exhibits more similarities with the Earth's today-and even more so in the past than any other solar System body (e.g., Coustenis and Taylor, 2008). Recent Cassini-Huygens findings hâve revolutionized our understanding of Titan's System and its potential for harboring the "ingrédients" necessary for life (Coustenis and Taylor, 2008; Lorenz and Mitton, 2008). Recent discoveries reveal that beyond its rich organic budget, and sufficient energy sources to drive Chemical évolution, Titan also probably contains a vast subsurface océan (Lorenz et al., 2008b).

Titan is indeed a very complex world much as our own planet. It is the only one other than Earth, that possesses a thick nitrogen-based atmosphère, four times denser than on our own atmosphère, with a rich organic chemistry (Fig. 4).

It also has a geologically complex and active surface including lake-like features filled with organic liquid (e.g., Stofan et al., 2007). The physical processes within this world invite further close-up investigation that will provide a better understanding of the terrestrial processes as well.

Current investigations hâve shown that Titan fulfills many of life's prerequisites for a carbonaoeous portfolio. Due to its nitrogen atmosphère, which is not in Chemical equilibrium but like a Chemical factory initiâtes the formation of complex positive and négative ions in the high thermosphère as a conséquence of induoed magnetospheric-ionospheric-atmospheric interactions involving EUV, UV radiation, energetic ions, and électrons (as recently demonstrated by the Cassini Ion andNeutral Mass Spectrometer, INMS). In this dynamic evolving environment, the second most abundant atmospheric constituent, methane, is dissociated irreversibly to produoe a variety of trace gases such as hydrocarbons (e.g., ethane, acetylene, and propane) and in combination with the nitrogen, nitriles (e.g., hydrogen, cyanide, acetonitrile, cyanoacetylene), which are detected in the stratosphère (betvveen 70 and 500 km in altitude) by the Composite Infrared Spectrometer (CIRS) onboard Cassini (e.g., Coustenis et al., 2007[START_REF] Coustenis | Titan trace gaseous composition from CIRS at the end of the Cassini-Huygens prime mission[END_REF].

Literally, some of these gases would be considered as signs of life if they were on our planet (HCN is considered a prebiotic molécule, a precursor of life). Hence, finding how they form on Titan could give us dues on how life began on Earth.

The Cassini-Huygens mission lias revealed the essential details of the organic and methane hydrologie cycles that we see today on Titan (Raulin et al., 2008;Brown et al., 2009b;Lebreton et al., 2009). Methane on Titan seems to play the rôle of waler on Earth, with a similar complex cycle as shown in Fig. 5 (Atreya et al., 2006). On Titan, where methane is photodissociated and forms ethane and other organic products in the atmosphère, it should hâve disappeared after 10-100 million years, with around 34 Myrs as a nominal period (Atreya et al., 2006).

The intriguing question then is of how methane gets replenished in the atmosphère. On Earth today, it is life itself that refreshes the methane supply since methane is a by-product of the metabolism of many organisms. Hence, could this mean there is life on Titan? However, the Huygens Gas Chromatograph Mass Spectrometer (GCMS) data hâve shown that methane is not of biogenic origin, because the isotopic ratios are compatible with inorganic values (Niemann et al., 2005(Niemann et al., , 2010)). Thus, the sinks of atmospheric methane on Titan are relatively well understood, but the major sources of replenishment are still very model dépend ent, as will be discussed later.

THE ATMOSPHERIC ORGANIC-RICH EN VIRONMENT OF TITAN

Although Titan's atmosphère is much colder than Earth's, it présents many direct similarities with our planet (Fig. 4), at different levels which hâve been pointed out since Voyager days. To begin with, both are made of the same main constituent, dinitrogen. A similar vertical structure from the troposphère to the ionosphère is also présent, as well as a surface pressure of 40% larger than that of the Earth's (Fulchignoni et al., 2005). This is the only case of an extraterrestrial planetary atmospheric pressure so similar to that of Earth. Furthermore, a very exciting and complex organic chemistry takes place in Titan's atmosphère.

The direct analysis of the ionosphère by the Cassini Ion and Neutral Mass Spectrometer (INMS) instrument during the low-altitude Cassini flybys of Titan shows the presence of many organic species at détectable levels, at very high alti tudes (1,100-1,300 km). Extrapolation of the INMS measurements (limited to mass up to 100 Da) and of Cassini Plasma Spectrometer (CAPS) data strongly suggests that high molecular weight species (up to several 1,000 Da) may be présent in the ionosphère (Waite et al., 2007). These observations open a fully new vision of the organic processes occurring in Titan's atmosphère, with a strong implication of the ionospheric chemistry in the formation of complex organic -80km. showed évidence for their formation around 1,000 km in altitude (Waite et al., 2007), and the Cassini/ CIRS instrument detects the neutral chemistry in the stratosphère (Flasar et al., 2005).

compounds in Titan's environment (Fig. 6), which was not envisaged before.

These compounds are détectable in solar and stellar UV occultations and initiate the process of haze formation starting at about 950 km (Waite et al., 2007) to finally condense out.

In the neutral atmosphère of Titan (between roughly 100 and 500 km), as we mentioned above, CH4 chemistry is coupled with N, chemistry producing the formation of many organics in gas and particulate phase: hydrocarbons, nitriles, and complex refractory organics. Several photochemical models describing the Chemical and physical pathways involved in the Chemical évolution of the atmosphère of Titan hâve been published for the last20 years (Yung et al., 1984;Toublanc et al., 1995;Wilson and Atreya, 2004;Lawas et al., 2008). These papers estimate the resulting vertical concentration profiles of the different molécules.

Based on these models, the cycling of volatile Chemicals starts with the dissocia tion of N2 and CH4 through électron, photon, and cosmic rays impacts in Titan's atmosphère. The primary processes allow for the formation of acetylene (C,H2)

and hydrogen cyanide (HCN) in the high atmosphère. These molécules play a key rôle in the general Chemical sclieme: once they are formed, they diffuse down to the lower levels where they allow the formation of higher hydrocarbons and nitriles and perhaps aromatic compounds. Additionally, methane dissociation probably also occurs in the low stratosphère through photocatalytic processes involving acetylene and polyynes. The end products of the Chemical évolution of methane in the atmosphère are complex refractory organic compounds and ethane.

As these aérosols and haze particles fall through the atmosphère and grow, they become détectable with imaging Systems su ch as the Cassini Imaging Science Subsystem (ISS) at about 500 km altitude and are ubiquitous throughout the stratosphère (Porco et al., 2005). They are strong absorbers of solar UV and visible radiation and play a fundamental rôle in heating Titan's stratosphère and driving wind Systems in the middle atmosphère, much as ozone does in the Earth's middle atmosphère.

Experiments that simulate the reactions taking place in Titan's atmosphère (such as the Miller-Urey experiment) produce refractory organics, usually named tholins (Sagan and Khare, 1979) ("tholins" are solid products in the laboratory mimicking complex refractory organics: Nguyen et al. (2007); and references therein). Tholins represent laboratory analogues of Titan's aérosols and are useful to interpret many obsejvational data which require information on the aéro sols. As being experimental analogues of Titan's atmospheric particles, tholins also allow the study of the aérosols behavior in Titan's conditions with the tools available in the laboratory. Several organic compounds hâve already been detected in Titan's upper and lower atmosphère (Waite et al., 2007;[START_REF] Coustenis | Titan trace gaseous composition from CIRS at the end of the Cassini-Huygens prime mission[END_REF].

The list includes hydrocarbons (both with saturated and unsaturated chains) and nitrogen-containing organic compounds, exelusively nitriles, as expected from laboratory simulation experiments (Fig. 6). Moreover, since the Cassini arrivai in the Saturnian system in 2004, the presence of water vapor and benzene has been unambiguously confirmed by the CIRS instrument.

TITAN'S PREBIOTIC RELEVANCE

Several of the organic processes that are occurring today on Titan form some of the organic compounds which are considered as key molécules in terrestrial prebiotic chemistry, such as hydrogen cyanide (HCN), cyanoacetylene (HC3N), and cyanogen (C2N2). In fact, with several percent of methane in dinitrogen, the atmosphère of Titan is one of the most favorable atmosphères for prebiotic synthesis, although (Coustenis et al., 2010a).

The Aérosol Collector Pyrolyser (ACP) experiment on Huygens provided the first direct in situ Chemical analysis of Titan's aérosols. It collected lraze particles from the stratosphère and the troposphère, heated them at different températures, and sent the produced gases for analysis to the GCMS instrument.

The obtained results indicated that the aérosols are ma de of a refractory nucléus, composed of H, C, and N atoms (Fig. 7) and producing NH3 and HCN after pyrolysis at 600°C (Israël et al., 2005). This strongly supports the tholin hypothesis. It also strongly suggests that Titan's aérosols may evolve once in contact with water ice on the surface and may produce a variety of organics of biological inter est, such as amino acids (Neish et al., 2010;[START_REF] Ramirez | The fate of aérosols on the surface of Titan[END_REF].

Analogies are thus obvious between the organic cliemistry activity currently occurring on Titan and the prebiotic chemistry which was once active on the primitive Earth, prior to the emergenoe of life (e.g., McKay and Smith, 2005).

Indeed, in spite of the absence of permanent bodies of liquid water on Titan's surface, both chemistries are similar. As noted earlier, many resemblances can also be made between the rôle of methane on Titan and that of water on the Earth, with a complex cycle that has yet to be fully understood. Indeed, on Titan, methane can exist as a gas, liquid, and solid, since the mean surface température is almost 94 K (Fulchignoni et al., 2005), approaching the triple point of methane.

The atmosphère we enjoy today on Earth is probably radically different from the primitive one dating back to 4.6 billion years ago when our planet was nothing more than a molten bail of rock surrounded by an atmosphère of hydrogen and hélium. In the absence of a magnetic field at those early stages, the intense solar wind from the young Sun blew this early atmosphère away. Then, as Earth cooled enough to form a solid crust, it was covered with active volcanoes, ejecting water vapor, carbon dioxide, and ammonia to form an early toxic atmosphère.

Eventually, light from the Sun broke down the ammonia molécules exsolved by the volcanoes, releasing nitrogen into the atmosphère. Over billions of years, the quantity of nitrogen built. up to the levels we see today. Although life formed just a few hundred million years later, it was not until the évolution of bacteria, 3.3 billion years ago, that the early Earth atmosphère changed into the one we know today.

During the period from 2.7 to 2.2 billion years ago, these early bacteria-known as cyanobacteria-used energy from the Sun for photosynthesis and released oxygen as a by-product. They also trapped carbon dioxide in organic molécules.

In just a few hundred million years, these bacteria completely changed the Earth's atmosphère composition, bringing us to the current mixture of 21% oxygen and 78%

nitrogen. Schaefer and Fegley (2007) predict that Earth's early atmosphère contained CH,, H2, H20, N2, and NH3, similar to the components used in the Miller-Urey synthesis of organic compounds, often related to Titan's and Enceladus's atmospheric inventory. Furthermore, Trainer et al. (2006) looked at the processes that formed haze on Titan and on early Earth and found many similarities for what.

could hâve served as a primary source of organic material to the surface.

Before the rise of the atmospheric oxygen in the terrestrial atmosphère 2.5 Gy ago, it considered possible that the abundance of methane gas was 10-20 times higher than the today's value of 1.6 * 10"6 (Pavlov et al., 2003). Hence, if the atmospheric C02/CH4 ratio had become equal to 10 at the mid-Achaean era, an organic haze could hâve formed on this early environment (Pavlov et al., 2000;DeWitt et al., 2009). This hydrocarbon haze produced the anti-greenhouse effect which reduced the température of the atmosphère (Kasting and Howard. 2006). Titan also hosts a thick methane-induced organic haze, similar to the one predicted for the early Earth and, consequently, expériences the same antigreenhouse effect (McKay et al., 1999). The absence of vast amounts of CO, on Titan is one of the major différences between the two atmospheric envelopes.

On the other hand, hydrogen cyanide and other prebiotic molécules are among the starting materials for biosynthesis. The existence of hydrocarbons, and in partieular acetylene and benzene, has really enlarged the borders of photochemical organic products.

Especially, the presence of benzene (C6H6) seems extremely interesting, as it is the only polycyclic aromatic hydrocarbon (PAHs) discovered on Titan today.

The presence of PAHs on Titan's atmosphère is very important as they could contribute to the synthesis of biological building blocks. Moreover, the combi nation of the liquid deposits on the surface of Titan and the low température could host the proper environment for this biosynthesis. Recent laboratory experiments showed that aromatic compounds could be plausibly produced on icy surfaces (Menor-Salvân et al., 2008). Benzene was first detected at 674 cm"1 based on Infrared Space Observatory (ISO/SWS) data (Coustenis et al., 2003) with the mixing ratio of 4 x 10"10 and was then also detected in the thermosphère (950-1,150 km) from the analysis of Cassini/INMS data (Waite et al., 2007) and firmly in the stratosphère (100-200 km) at ail latitudes by Cassini/CIRS ATHENA COUSTENfS ETAL measurements (Flasar et al., 2005;Coustenis et al., 2007[START_REF] Coustenis | Titan trace gaseous composition from CIRS at the end of the Cassini-Huygens prime mission[END_REF]. Moreover, benzene has been tentatively identified on Titan's surface by Huygens/CGMS measurements (Niemann et al., 2005).

As Titan lacks oxygen and sufficiently high températures, as did primitive Earth, different evolutionary pathways from Earth must hâve been followed on Titan, and polyphenyls may possibly be created (Delitsky and McKay, 2010).

The abundances of hydrocarbons are higher on Titan tlian those on Earth by a factor of about 10M04. Moreover, the temporal variations of the hydrocarbon traces on Titan expérience a full cycle during the Titan year (Coustenis et al., 2010a) and are probably inftuenced by local or régional sources and sinks. Photochemieal models are trying to reproduce these phenomena. Taking into account ail the above-described characteristics, the prebiotic potential of Titan is enormous, and a huge effort in astrobiological studies is focused on its environment. Eventually, Titan still seems to be the idéal planet-size laboratory for increasing our knowledge of the évolution of the Earth's atmosphère.

ORGANIC CHEMISTRY AND MORPIIOLOGY OF TITAN'S ACTIVE SURFACE

The Cassini-Huygens mission has significantly enhanced our understanding of Titan as the largest abiotic organic factory in the solar System. The abundance of methane and its organic products in the atmosphère, seas, and dunes exceeds the carbon inventory in the Earth's océan, biosphère, and fossil fuel réservoirs by more than an order of magnitude (Lorenz et al., 2008a). As discovered by the Cassini/INMS, in the high atmosphère, heavy ions are formed (Waite et al., 2007).

Measurements throughout the atmosphère, both remote and in situ, hâve indicated the presence of numerous hydrocarbon and nitrile gases, as well as a complex layering of organic aérosols that persists ail the way down to the surface of the moon [START_REF] Tomasko | Rain, winds and haze during the Huygens probe's descent to Titan's surface[END_REF]Coustenis et al., 2007;de Kok et al., 2007).

Radar observations suggest that the ultimate fate of this aérosol "rain" is the génération of expansive organic dunes that produce an équatorial belt around the surface. Condensation of these species on aérosol particles is a probable explanation for these atmospheric characteristics. These particles, for which no direct data on their Chemical composition were previously available, were analyzed by the Aérosol Collector and Pyrolyser instrument aboard Huvgens probe. ACP results show that the aérosol particles are made of refractory organics which release HCN and NH3 during pyrolysis.

This supports the tholin hypothesis (as described in section 3.1). From these new in situ measurements, it seems very likely that the aérosol particles are made of a refractory organic nucléus, covered with condensed volatile compounds (Israël et. al., 2005). However, Huygens/GCMS did not detect. a large variety of organic compounds in the low atmosphère (Niemann et al., 2005).

Figure 8. Linear dunes on Titan [START_REF] Radebaugh | Linear dunes on Titan and eartli: Initial remote sensing comparisons Ramirez SI[END_REF]). Titan's dunes are believed to be composed of ice and organics grains that possibly dérivé from a combination of the surface icc and the organic Chemicals that fall through in Titan's atmosphère.

Moreover, the nature and abundances of the condensâtes hâve not been measured. Even more importantly for astrobiology, neither the elemental composition nor the molecular structure of the refractory part of the aérosols has been determined.

Eventually, these complex organic molécules are deposited on Titan's surface in large quantities, where data from Cassini's instruments hint at their existence.

Hence, the upper thermosphère is linked intimately with the surface and the intervening atmosphère. In spite of the low surface température, the organics reaching the surface are probably evolving once in contact with water ice and may form organic molécules of biological interest.

Radar observations suggest that the ultimate fate of this aérosol "rain" is the génération of expansive organic dunes that produce an équatorial belt around the surface. Indeed, the surface of Titan shows also the presenoe of sedimentological and meteorological processes, as we see on Earth: There are many large dune areas (Lorenz et al., 2006;Radebaugh et al., 2008[START_REF] Radebaugh | Linear dunes on Titan and eartli: Initial remote sensing comparisons Ramirez SI[END_REF] (Fig. 8) where the terrestrial silica sand is probably replaced, once more, by water ice particles, mixed with the organic material of the aérosols.

Cassini's radar instrument finally unveiled what appears to be a land of lakes in Titan's northern polar régions (see Fig. 9) (Stofan et al., 2007). Cassini/ISS images also show a kidney-shaped dark feature about 200 km in lengtli, named Ontario Lacus, that is outside the area of radar coverage and has receutly been confinned by the Cassini Visual and Infrared Mapping Spectrometer (VIMS) not only to be a lake but also to be composed of ethane liquid (Brown et al., 2008). In the absence of a massive surface océan but with analogues to ail other terrestrial hydrological phenomena présent, Titan's methane cycle is very exotic.

The liquid bodies are one of the main astrobiological aspects of Titan.

Cassini's caméras (ISS) hâve allowed scientists to compile a nearly global surface map and tomonitor the surface and atmosphère for activity. Intriguingly, repeated south-polar imaging by ISS revealed différences consistent with ponding of hydrocarbon liquids on the surface due to précipitation from a large storm.

More recent ISS images of high northem latitudes illustrate the full extents (>500,000 km2) of hydrocarbon seas, also observed by Cassini's radar. These observations demonstrate dynamic processes at work on Titan and indicate that the pôles harbor liquid-hydrocarbon réservoirs, the extents of which differ from pôle to pôle and which may be coupled to seasonally varying circulation (Turtle et al., 2009).

The lakes and seas observed on Titan in the polar régions (Mitri et al., 2007;Stofan et al., 2007) make Titan the only body in the solar System having large liquid bodies on its surface. These very dark features at the high northem latitudes of Titan were finally shown to be liquid-filled (most probably with ethane-rich mixtures (Brown et al., 2008)) basins-classifying them as lakes.

The features range in size from less than 10 km2 to at least 100,000 km2.

They are limited to the région poleward of 55°N. Currently, Cassini's instruments hâve identified and mapped almost 655 geological structures referred as lakes and/or basins (Fig. 10) (Hayes et al., 2008). Titan is thus the only planetary body, other than the Earth, with long-standing bodies of liquid on its surface (although direct observational evidence of the longevity of Titan's surface liquids remains {rtght) represents the spatial distribution of mapping units. Lake Michigan is illustrated for scalc purposes (Hayes et al., 2008).

to be obtaiiied). Ail of this suggests that Titan maybe even more similar to primitive Earth than we thought. However, the degree of complexity which can be reached from such an organic chemistry in absence of permanent liquid water bodies on Titan's surface is still unknown, although it could be quite high.

McKay and Smith (2005) noted the astrobiological importance of these geological features that are filled with liquid hydrocarbons, since there is a possibility for a different form of life to exist in such environments. It has been hypothesized that such a methanogenic life form consumes H2 instead of 02 that could be measured in the lower atmosphère. Two papers, by Strobel (2010) and Clark et al. (2010), based on data from the Cassini orbiter focus on the complex Chemical activity on the surface of Titan. Strobel (2010) shows that hydrogen flows down through Titan's atmosphère and then somehow disappears on the surface. One of the most interesting phenomena occurring on Titan is that impor tant quantities of atmospheric hydrogen précipitâtes and disappears when reach the surface. Such process resembles the oxygen consumption as oocurring on

Earth although in Titan's case, the element is hydrogen (Strobel, 2010).

Even though this is not supportive to the Titan's terrestrial-type life theory, it represents a hypothetical second form of life independent from water-based life we know on Earth. Strobel (2010) describes densities of hydrogen in different parts of the atmosphère and the surface. Previous models had predicted that hydrogen molécules, a by-product of ultraviolet sunlight breaking apart acetylene and methane molécules in the upper atmosphère, should be distributed fairly evenly throughout the atmospheric layers. The authors found a disparity in the hydrogen densities that lead to a flow down to the surface; at a rate of about 10,000 trillion hydrogen molécules per second. This is about the same rate at which the molécules escape out of the upper atmosphère. Strobel (2010) States that it is not likely for hydrogen to be stored in a cave or underground space on Titan. Titan's surface is also so cold that a Chemical process that involved a catalyst would be needed to convert hydrogen molécules and acetylene back to methane, even though overall there would be a net release of energy. The energy barrier could be overcome if there were an unknown minerai acting as the catalyst on Titan's surface.

Another possible indicator for life on Titan is the lack of acetylene on the surface since there is no clear evidence of this compound in the received data to date, while it is expected to hâve been deposited through the atmosphère. It has been suggested that this could be due to the fact that sonie form of life on the surface is using acetylene as an energy source (Clark et al., 2010). This theory is largely debated and controversial among the scientific community, especially due to suggestions of nonbiological origin of this phenomenon; however, it has the merit to propose new interesting astrobiological théories. In detail, Clark et al.

(2010) map hydrocarbons on the surface from Cassini/VIMS data and find a lack of acetylene. McKay and Smith (2005) had identihed acetylene as the best. energy source for methane-based life on Titan. While nonbiological chemistry offers one possible explanation, these authors believe these Chemical signatures bolster the argument for a primitive, exotic form of life or precursor to life on Titan's surface. According to one theory put forth by astrobiologists, the signatures fulfill two important conditions necessary for a hypothesized "methane-based life", which would consume not only methane but also hydrogen. However, one other possibility is that sunlight or cosmic rays are transforming the acetylene in icy aérosols in the atmosphère into more complex molécules that would fall to the ground with no acetylene signature.

To date, methane-based life forms are only hypothetical. Scientists hâve not yet detected this form of life anywhere, though there are liquid-water-based microbes on Earth that thrive on methane or produce it as a waste product. At Titan's low températures, a methane-based organism would hâve to use a sub stance that is liquid as its medium for living processes, but not water itself. Water is frozen solid on Titan's surface and much too cold to support life as we know it.

The list. of liquid candidates includes liquid methane and related molécules like ethane. While liquid water is widely regarded as necessary for life, there has been extensive spéculation published in the scientific lit.erature that this is not a strict requirement. The new hydrogen findings are consistent with conditions that could produce an exotic, methane-based life form, but do not prove its existence.

INTERIOR MODELS FOR TITAN AND ITS POSSIBLE SUBSURFACE OCEAN

Regarding Titan's interior structure, since the only in situ data from its surface are the Huygens probe recordings, only modeling assumptions can be presented.

Structural models for planetary interiors suggest that Titan, like Europa, Ganymede, Figure 11. Illustration of Titan's internai structure with a liquid océan between two subsurface ice layers (From Tobie et al., 2005). Cassini-Huygcns rccorded cxtremcly low-frequency radio waves which supports the existence of this liquid subsurface layer (NASA/LPGN).

and Callisto, lias maintained internai liquid water réservoirs, probably mixed with sorae ammonia and more speculatively sulfur and possibly entrained methane clathrates (Fig. 11).

Another possible location to look for life ou Titan would be in an undersurface liquid water océan. The presence of such an internai océan is supported by Titan internai structure raodels (Tobie et al., 2005;Mitri et al., 2008), radar and gravityCassinimeasurements (Lorenzetal., 2008b), and the HuygensAtmospheric Structure Instrument (HASI) experiment. [START_REF] Beghin | New insights on Titan's plasma-driven Schumann résonance inferred from Huygens and Cassini data[END_REF] thus interpreted the extremely low-frequency electric signal recorded by HASI as a Schumann résonance between the ionosphère and a modestly conducting océan (siuce the ice is not conductive) roughly 30-50 km below the surface. Thermal évolution models suggest that Titan may hâve an ice crust between 50 and 150 km thick, lying atop a liquid water océan a couple of 100 km deep, with some amount (a few to 30%, most likely -10%) of ammonia dissolved in it, acting as an antifreeze material (Lorenz et al., 2008b). This correspond to a pH arouud 11.5. The pressure reaches~5 kbars at 200 km depth, and it could include hot spots reacliing -20 °C.

Such conditions are not incompatible with life as we know it on Earth (Fortes, 2000;Raulin, 2008;[START_REF] Raulin | Titan's astrobiology[END_REF]. Tobie et al. (2005) suggested a layered interior structure of Titan, consisting of a rocky core overlaid by high-pressure ice, a liquid layer overlaid by low pres sure ice, and finally a solid icy crust (Fig. 11). An earlier model by Fortes (2000) noted that underneath Titan's icy crust, at a depth of approximately 200 km, there lies an ammonia-water solution océan in which life could survive. Another, more recent model by Mitri et al. (2008) suggests pockets of methane clathrates trapped within an ammonia-water océan which could exsolute and produce overpressure and consequently the ammonia-water can erupt to the surface leading to cryovolcanic phenomena.

With regard to Titan's morphology and internai dynamic geology, it lias been suggested that there may be active cryovolcanoes on Titan (Sotin et al., 2005;[START_REF] Lopes | Much like Earth: distribution of géologie processes on Titan from Cassini RADAR data[END_REF] since traces of former flows hâve been found across parts of the surface. A wide variety of Cassini data support the presence of cryovolcanism on Titan. Indeed, at least two régions hâve been observed to change réflectance on Titan's surface (Tui Regio 20°S, 130°W; Hotei Regio 26°S, 78°W) (Barnes et al., 2006;Nelson et al., 2009a, b), and one of them (Hotei Regio) in addition to a recent discovery (Sotra Facula 15°S, 42°W) in radar images exhibit lobate "flow" forms (Soderblom et al., 2009;Lopes et al., 2010), consistent with the morphology of volcanic terrain, support.ing the hypothesis of cryovolcanic éruptions (Sotin et al., 2005).

However, on Titan, fluid water mobilized and made buoyant by ammonia and other materials could replace terrestrial melted silicates. Cryovolcanism suggests a dynamic process than involves the interior, the surface, and the atmos phère as well. It is a multi-complex activity that resembles terrestrial volcanic pro cesses as it follows a similar pattern although in extremely altered conditions and different initial and depositional products. Cryovolcanism on Titan is believed to be a significant source of the methane in the atmosphère (Tobie et al., 2006).

An underground liquid océan, several hundred kilomet.ers deep at. the surface of Titan, is suggested to be the source of cryomagma, hence outgassing methane into the atmosphère and thus replenishing the destroyed amounts.

There are indeed studies suggesting the presence of an internai ammoniawater océan (Grasset and Sotin, 1996;[START_REF] Grasset | On the internai structure and dynamics of Titan[END_REF]Tobie et al., 2005;Mitri et al., 2007) while another study has modeled and suggested an océan filled with methane clathrate pockets that lead to explosive cryovolcanism (Fortes et al., 2007).

The theory of trapped methane clathrates in the potential liquid océan is of a major astrobiological interest.. The presence of methane clathrates in an aqueous environment is attached to the "clathrate gun" hypothesis. This hypo thesis suggests that potential movement and rise of the température in an underground liquid réservoir could trigger the sudden release of methane from methane clathrate compounds buried in permafrost. or seabeds (Kennett et al., 2003) or an océan like on Titan's case. The initiation of such process leads to further température rise and further methane clathrate destabilization that could easily cause and trigger cryovolcanic éruptions (Kennett et al., 2003). For Titan, a dynamic process like the one suggested by the clathrate gun hypothesis could result. to increase of température values, creating an environment more favorable for life to exist.

With respect to the astrobiological interest., there is a controversy about the effectiveness of methane as a medium for life compared to water or ammonia.

Water has higher solubility than methane, enabling easier transport of substances in acell, while methane's lesser Chemical reactivity allows for the easier formation of large structures corresponding to proteins (Benner et al., 2004). In addition, the cryovolcanic activity suggests higher températures within the océan and the volcanic conduit where heat transfer between the interior and upper layers would be critical in sustaining any subsurface oceanic life (Grasset et al., 2000).

Thus, the possibility of life in this océan cannot be excluded. Moreover, models also predict Lhat during the first tens millions of years after Titan's forma tion, the océan was in direct contact with the atmosphère on one side and with the bedrock on the side. This could hâve provided conditions very favorable for an efficient prebiotic chemistry toward the emergence of life, with the possible involvement of hydrothermal vents. Thus, the internai océan of Titan not only is habitable but could be inhabited.

Hence, in spite of the low température. Titan is not a congealed Earth:

The Chemical System is not frozen. Titan is an evolving planetary body and so is its chemistry. Once deposited on Titan's surface, the aérosols and their complex organic content may follow a Chemical évolution of astrobiological interest.

Laboratory experiments show that, once in contact with liquid water, Titan tholins can release many compounds of biological interest, such as amino acids (Khare et al., 1986). Such processes could be particularly favorable in zones of Titan's surface where cryovolcanism may occur. The N,-CH4 by-products in Titan's atmosphère eventually end up as sédiments on the surface, where tliey accumulate presently at a rate of roughly 0.5 km in 4.5 Gyr.

Long-term Chemical évolution is impossible to study in the laboratory: In situ measurements of Titan's surface thus offer a unique opportunity to study bv a ground-truth approach some of the many processes which could be involved in prebiotic chemistry, including isotopic and enantiomeric fractionation (Nguyen étal., 2007).

There are suggestions that Titan is presently geologically active on the sur face [START_REF] Nelson | Photometric properties of Titan's surface frorn Cassini VIMS: relevance to titans hemispherical albedo dichotomy and surface stability[END_REF](Nelson et al., , 2009a, b) , b) and in its interior. If Titan is currently active, then these results raise the following questions: What is the full extent of current géologie activity? What are the ongoing processes? Are Titan's Chemical processes today supporting a prebiotic chemistry similar to that under which life evolved on Earth?

Although the Chemical reactions that lead to life on Earth take place in liquid water, the reactions themselves are almost entirely between organics.

The study of organic chemistry is therefore an important, and arguably richer, adjunct t.0 the pursuit of liquid water in the solar System. Titan's organic inventory is impressive, and carbon-bearing compounds are widespread across the surface in the form of lakes, seas, dunes, and, probably, sedimentary deltas at the mout.hs of channels. As concerns oxidant. agents on Enceladus, [START_REF] Gaidos | Biogeochemistry: life in ice-covered ooeans[END_REF] also noted their production through E-ring particles, charged from the Saturnian magnetospheric environment. Then, some suitable geological process is needed to mix the reducing and oxidizing compounds. In a model by Cooper et al. (2009), the astrobiological parameters that support life on Enceladus are evaluated as higher than for Europa due to a less extreme State of oxidation and greater residual abundance of organics.

Indeed, Enceladus, Saturn's most. active moon as observed from Cassini, présents a mystery in the studies of planetary science and, more specifically, in geophysics. The enigma evolves around how a small moon can possess suffîcient dynamical energy to drive a geyser plume rising 600 miles in space out. of the moon's south pôle, eventually feeding its material to the outer E-ring of Saturn (Postberg et al., 2009). The heat source for Enceladus is still an open question, as is the possibility for life to exist. on this small satellite if an underground liquid water océan or liquid water subsurface pockets exist to explain the plumes [START_REF] Kieffer | Enceladus -oasis or ice bail?[END_REF].

The geyser plumes arise from the warm surface surrounding and including four parallel faults located at the moon's south pôle, the "Tiger Stripes" fractures (Porco et al., 2006;Spencer et al., 2006), spewing a sériés of jets more than 600 km high (Fig. 12). The mass production rate of the plume gas has been estimated to be ~150 kg/s from occultation data (Tian et al., 2007). This value is surprisingly high, suffîcient to remove a significant fraction (~20%) of Enceladus's mass over the âge of the solar System (Kargel, 2006).

Data obtained by the VIMS instrument on board Cassini indicate C02 and organics as possible components (Brown et. al., 2006). Data processing showed components that Cassini's INMS identified H20 as the prédominant component, C02 as the second most abundant, methane, and trace quantifies of acetylene and propane (Waite et al., 2009). During the 9th of October 2008 flyby, Cassini dived into the south-polar plume, and INMS reported the presence of ammonia and other various organic compounds like deuterium and 40Ar, as well as complex organics like benzene and other probable species such as methanol and formaldé hyde (Waite et al., 2009). The Chemical composition of the plume and surface material of Enceladus suggests the presence of a heat source in its interior, hot et al., 2006).

enough to décomposé ammonia into N2 and to drive reactions with hydrocarbons, implying internai températures on the order of 500-800 K (Matson et al., 2007).

There is a plethora of competing théories regarding the triggering and exso lution of geysers from the Tiger Stripes. The main controversy lies on as to whether the plumes are formed by a massive underground océan (Tobie et al., 2008;Postberg et al., 2009) (Fig. 13a) or if the material generates from ice warmed, melted, or crushed by tectonic-like motions (Nimmo et al., 2007) (Fig. 13b).

Given the aforementioned observations and analysis, it is now known that Enoeladus contains enough heat to drive complex and dynamic géologie activity.

Interprétation of Cassini data with a view to explain the major internai réservoir that triggers and feeds this dynamic prooess points to a possible underground liquid océan beneath the icy crust Since the prerequisites for life to emeige are the simultaneous existence of energy, organic compounds, and liquid water that hâve been found on Enoeladus, it seems that it possesses ail the neoessary components to support life.

In general, as described above, the south-polar région of Enceladus présents extremely high températures while the exsoluted vapor has also been shown to contain simple organic compounds. Most théories regarding the origin of this active région suggest that it is very likely for a liquid water environment to exist beneath the Tiger Stripes. This hypothesis enables the parallelism between potential biological ecosystems on Enceladus and the already existing on Earth (Fig. 14). VIMS observations of several other satellites of Satura, in the near-infrared région, show that their surface is covered by dark materials. This is particularly the case with Dione but also with Phoebe, Iapetus, Hyperion, Epimetheus and even with the F-ring [START_REF] Clark | Europa Jupiter System mission. Joint summary report[END_REF]. This dark material could be made of cyanide compounds and could be of cometary origin [START_REF] Clark | Europa Jupiter System mission. Joint summary report[END_REF].

Spectral signatures of hydrocarbons hâve also been found from VIMS data on Iapetus and Phoebe suggesting the presence of organic compounds such as PAH's, kerogen-, or coal-like structure compounds (Cruishank et al., 2008).

Some of these satellites may also hâve internai liquid water pockets and thus may présent the requisite essential for the emergence and development of life: liquid water and organic compounds. The solstice mission may be able to discover such properties and thus extend drastically the list of planetary bodies of impor tant astrobiological interest in the Saturnian system.

5. Discussion: What Can Titan and Enceladus Tell Us About Life?

The complex mechauisms that hâve led to the emergence and development of life on Earth are still under investigation. Despite the great strides on biological sciences, the roots, the sources, and the initial conditions of life still remain unknown. Some answers can be found in extraterrestrial environments, hopefully in our solar system neighborhood or in exoplanets orbiting Sun-type stars. Without doubt, every planetary body, with a possible astrobiological potential. is a target for further investigations. Currently, after years of explorations, extended missions, and data analysis, it. appears that among the main candidates for finding signs of past or current life witliin our solar System are Mars, Enceladus, Europa. and Titan (not in any particular order of priority or importance).

In this context, Titan and Enceladus, both orbiting Saturn, seem idéal locations for further investigation. As said before, their astrobiological impor tance is obvious, because they propose uniquefy ail the necessary ingrédients for life emerging and évolution. Not surprisingly, scientists expect that their study will provide some important insights on the origin of life.

The surface of Titan appears, like the surface of Mars or Europa, as an unlikely location for extant life, at least for terrestrial-type life. Even though Titan présents terrestrial-type geology with complex structures formed mostly from dynamic processes, the absence of water on the surface makes it unlikely to support terrestrial-type life. Liquid water, if it exists, is not presently in contact with a silicate core, which is isolated from the subsurface océan by a layer of a high-pressure ice phase (Tobie et al., 2005). However, Fortes (2000) noted that Titan's internai water océan might support terrestrial-type life that had been introduced t.here previously or formed when liquid water was in contact with silicates early in Titan's history. According to McKay and Smith (2(305), photochemically derived sources of free energy on Titan's surface could maintain an exotic type of life, using liquid hydrocarbons as solvents. Similarly, Stoker et al. (1990) stressed that terrestrial bacteria can in fact satisfy their energy' and carbon needs by "eating" tholin. In this sense, a methane-rich atmosphère may act as a "poor planet's photosynthesis," providing a means to capture the free energy from ultraviolet light and make it available for metabolic reactions.

Consequently, it cannot be excluded that life may hâve emerged on or in Titan. In spite of the extreme conditions in this environment, life may hâve been able to adapt and to persist. Even though the possible current conditions (pH, température, pressure, sait concentrations) are not incompatible with life, as we know it on Earth (Tortes, 2000), the détection of a potential biological activit.y in the putative liquid mantle seems very challenging. Furthermore, as mentioned above, another possible location to look for life on Titan would be in a possible subsurface liquid water océan, and thus, it seems astrobiologically essential to confirm its presence.

Marine geologists and Marine biologists are nowadays close to confirm after many years of research and analysis that in the lower part of Earth's océan, below the thermocline, where the environmental conditions are extremely different than on the surface, life exists. In deep oceanic layers that suffer from low température and high pressure, tliere are two extreme environments where life is abundant.

Such environments are the cold seeps (vents) (e.g., Ritt et al., 2010) and the hydro thermal vents. Cold seeps are areas resembling brine pools, from which méthane and hydrogen sulfide and other hydrocarbon-rich fiuid seepage are released into the océan. In such conditions, types of life exist, feeding on single-cell Archaea and Eubacteria microbes that consume the methane and hydrogen sulfide from the seep. The environmental and living conditions described above resemble that of Titan's hy'pothesized internai liquid océan. Similarly, the température is low, the pressure is high, and possible bydrovolcanic events exist, releasing inethane and sulfides in the liquid océan. Thus, there are many aspects of Titan that might argue for the presence of sonie sort of very basic life on Titan. With Titan, we are observing an active hydrologie cycle subjected to seasonal, annual, and longer term changes, as on the Earth. Moreover, the future increase in the solar luminosity make it almost inévitable that eventually water on the Earth will no longer be trapped in our océan and troposphère but will escape rapidly in a process we see today for methane on Titan. The late stages of this évolution-an Earth with liquid water in the polar régions, in the crust, but no longer in an océan-may be echoed by the configuration we see today in Titan's methane hydrologie cycle.

Our understanding of the future of living beings on Earth (and hence the habitability in many ways of our planet) may then hâve something to gain from a thorough exploration of Titan's current state.

Enceladus, on the other hand, may hold the key to understanding an impor tant source of energy, plate tectonics, and volcanism. Indeed, many tectonic features on Enceladus may be analogous to features observed on other icy satellites such as Europa, Ganymede, and perhaps Titan. Thus, the study of the tectonics of Enceladus, which is currently active, can be used as a natural laboratory to investigate the response to stresses of the other surfaces of the outer solar System.

Moreover, Enceladus possesses a warm, chemically rich, environment that may facilitate complex organic chemistry and biological processes.

The Cassini-Huygens mission has enormously advanced our knowledge of the Saturnian System and the satellites within. As far as Titan and Enceladus are concerned, the wealth of data retrieved by the Cassini-Huygens mission will definitely be the reference point for future planetary investigation. However, the key contribution to planetary science of Cassini may be the questions raised rather than those answered. Some drawbacks of the mission, such as insufficient Figure 15. The TSSM basic concept. The orbiter (left), the balloon (center), and the lake lander (rigfit) [START_REF] Reh | TSSM final report on the NASA contribution to a joint mission with ESA[END_REF].

global coverage which has inhibited a full mapping of the atraospheric structure, composition, and temporal variations as well as the surface features of Titan and Enceladus, point to the need for further studies. Similarly, the part of the Titan's atmosphère between 400 and 950 km will remain unexplored (Coustenis et al., 2009). In addition, we will lack in situ measurements since the Cassini orbiter can only perfonn flybys of Titan and Enceladus, and the single vertical profile of the atmosphère taken by Huygens probe is limited to its landing site.

It is clear that Titan's organic chemistry and the possible subsurface océan (among other) remain to be investigated. In particular, joint measurements of large-scale and mesoscale topography and gravitational field anomalies on Titan from an orbiter and from an aerial platform would impose important constraints on the thickness of the "lithosphère", the presence of mass anomalies at depth, and any latéral variation of the ice mantle thickness. As discussed above, it is astrobiologically crucial to confirm the presenoe of such an internai océan, even though the water layer may not be in contact with the silicate core, like on Europa.

Lessons from Huygens will be used in the future to go back to Titan and explore in details its surface, in many locations. The intriguing discoveries of geological activity, excess warmth, and outgassing ou Enceladus (due perhaps to the éjection of water and organics from subsurface pockets bathed in heat or by some other meclianism) mandate a follow-up investigation of that tiny Satumian world that can only be achieved with high-resolution remote observations and detailed in situ investigations of the near-surface south-polar environment.

Among other options, a flagship (large) mission, TSSM (for Titan/Saturn System Mission), was proposed [START_REF] Reh | TSSM final report on the NASA contribution to a joint mission with ESA[END_REF]Coustenis et al., 2009), jointly studied by NASA and ESA. TSSM could explore extensively one of the Titan's lakes and study for several months Titan's atmospheric and surface environment (Fig. 15).

This concept would also investigate the astrobiological perspective of Titan and Enceladus. In particular, since hydrocarbon lakes on Titan's surface may harbor evidence for présent or past life, the proposed lake lander would play a crucial rôle. Similarly, the balloon platform would be responsible for understanding how volatile-rich worlds evolve and how organic chemistry and planetary évolu tion internet on large spatial and temporal scales.

The primary science goals of TSSM are to understand Titan's and Enceladus' atmosphères, surfaces, and interiors; to détermine the pre-and proto-biotic chemistry that may be occurring on both objects; and to dérivé constraints on the satellites' origin and évolution, both individually and in the context of the complex Saturnian System as a whole. Many internai processes play crucial rôles in the évolution of Titan and Enceladus. The formation and replenishing of Titan's atmosphère and the jet activity at Enceladus' south pôle are intimately linked to the satellite's interior structure and dynamics. Open issues are listed below:

1. To détermine their present-day structure and levels of activity 2. To détermine whether the satellites underwent significant tidal deformation and whether they possess intrinsic or induced magnetic fields and significant seismicity 3. To identify beat sources, internai réservoirs of volatiles (in particular met.hane and ammonia), and eruptive processes 4. To detect plausible evidence for life by analysis of hydrocarbons in the plume during close encount.ers A mission like TSSM would answer important astrobiological questions with précisé measurements as follows:

1. What degree of complexity is reached by Titan's organic chemistry in the dif ferent parts of the geological System?

2. Is Titan a habitable world? Does it hâve an undersurface liquid water océan or episodical liquid water bodies on the surface?

3. Is there currently, has there been or will there be biological activity on Titan?

To answer these questions and complément the Cassini-Huygens exploration of Titan, a dedicated orbiter and in situ éléments would help by providing data and analyses directly in the atmosphère, on the surface, and the interior of Titan. The exchanges among the different media and the external processes that affect Titan on time lapses of days, years, or seasons beg for further investigation, even beyond the solstice Cassini mission which will be operating until 2017. Besides TSSM, other concepts for future missions to return to Titan hâve been proposed

Such as a simple orbiter to perform close-up investigations of the surface and the atmosphère of Titan (JET, C. Sotin, PI). Further more, several in situ éléments hâve been proposed like:

The Aerial Vehicle for In-Situ and Airborne Titan Reconnaissance ( AVIATR), an alternative proposai to a Titan balloon mission. Since Titan expériences low gravity and a dense atmosphère, such a nuclear-powered airplane could fly more than 20 times easier than on Earth. It could saniple directly the atmosphère and cover huge swaths of Titan's landscape [START_REF] Barnes | AVI ATR : Aerial Vehicle for In-Situ and Airborne Titan Reconnaissance[END_REF][START_REF] Mckay | Titan's atmosphère and surface in 2026: the AVIATR Titan Airplane Mission, Titan Through Time[END_REF].

Titan Mare Explorer (TiME) is a proposed probe focusing on exploring Titan's lakelike features. This lake lander could study the Chemical composition and the geological characteristics of the hydrocarbon pools (Lorenz et al., 2009;[START_REF] Stofan | Titan Mare Explorer (time): a discovery mission to a Titan sea[END_REF].

Similarly, Titan lake probe is a lake lander which could be considered as part of the TSSM mission or as a stand-alone mission. The main objective of this proposai is to investigate the lake deposit and the physical properties of the liquids like the TiME concept (Waite et al., 2010).

Future exploration of the Saturnian neighborhood shall no doubt bring forth extremely important insights on our quest for the possibility of life and habitable sites elsewhere.

Europa, Ganymede, Titan, Enceladus Icy satellites of the gaseous giants, Jupiter and Satum, at orbits beyond the ice-line, constitute extremely interesting planetary bodies due to -among other -their unique geological [START_REF] Prockter | Characteristics of Icy Surfaces[END_REF]) and surface composition characteristics (Dalton, 2010;Dalton et al. 2010;[START_REF] Fortes | Phase Behaviour of Ices and Hydrates[END_REF], the possible internai water océan undemeath their icy crust and the habitability potential (Schubert et al. 2010;[START_REF] Tobie | Surface, Subsurface and Atmosphère Exchanges on the Satellites of the Outer Solar System[END_REF]Hussmann et al 2010;Sohl et al. 2010). Two of the largest Jovian satellites, Ganymede and Europa, as well as Satum's Enceladus and Titan, show not only surface features similar to the terrestrial planets and especially the Earth, but also internai heating and occasionally volcanism. Hereafter we summarize some of the characteristics of these moons.

In the neighborhood of Jupiter, two moons are good candidates for the internai liquid water océan and together form today the main targets of the Europa Jupiter System Mission (EJSM), a concept studied by ESA and NASA for a launch in 2020.

Europa hosts one of the smoothest surface topographies in the Solar System, composed essentially of water ice and other (Dalton, 2010;Dalton et al. 2010), and displays structures like linear chains (lineae) (e.g. Figueredo & Greeley, 2004), which most likely are crack developments caused primarily by diumal stresses [START_REF] Prockter | Characteristics of Icy Surfaces[END_REF]) and display the main characteristics of ridges. Other structures are dômes (pits that are surface dépréssions), dark spots and very few craters (Pappalardo et al. 1998a) as the Galileo mission showed. Internai stratigraphie modeling suggests that the satellite is primarily composed of silicate rock (olivine-dominated mineralogy) (e.g. [START_REF] Schubert | Interior composition, structure and dynamics of the Galilean satellites[END_REF]Sohl et al. 2010) and most probably an iron-rich core (Anderson et al. 1998). It is also suggested that the icy crust is decoupled from the moon's deeper interior due to the presence of a subsurface liquid océan (e.g. Schenk &Mckinnon, 1989;Zimmer et al. 2000;Schenk et al. 2004). The cracks, faults and lineae observed on the surface, hâve most probably been formed by endogenic processes that lead to tectonic movements and imply the presence of a subsurface océan that interact with the crust. Also, Europa suffers heavily bombardment by charged particles from the Jovian magnetosphere that lead to surface material décomposition and form the tenuous atmosphère composed mainly of oxygen [START_REF] Shematovich | Near-surface oxygen atmosphère at Europa[END_REF]Coustenis et al. 2010).

On the other hand, Ganymede, the largest satellite of the Solar System, is the only known moon to possess a magnetosphere (Kivelson et al. 2002). Most of Ganymede's surface coverage displays dark and brighter régions.

The former are filled with impact craters while the latter are covered by terrains curved by tectonic ridges and grooves (e.g. [START_REF] Pappalardo | Geology of Ganymede[END_REF]Patterson et al. 2010). Its internai structure possibly consists of a relatively small iron-rich core, overlain by silicate rocky material, which is covered by an icy crust. It is believed that a liquid océan exists within the mantle almost 200 km deep (e.g. McCord et al. 2001). The satellite possesses a thin oxygen atmosphère (Hall et al. 1998;Coustenis et al. 2010 and référencés therein).

The aforementioned information for Europa and Ganymede has been provided mostly from ground-based but also from in situ missions such as Pioneer 10 and 11, Voyager and more recently and efficiently Galileo. Taking into account the geological and structural éléments of both satellites, they readily become promising worlds of astrobiological potential. The fortheoming Europa Jupiter System Mission (EJSM) which is composed of two orbiters, one dedicated to Europa and one to Ganymede, will provide detailed investigation of these satellites (e.g. [START_REF] Blanc | LAPLACE: A mission to Europa and the Jupiter System for ESA's Cosmic Vision Programme[END_REF]).

Orbiting at circa 10 AU, the Kronian satellites Titan and Enceladus, stand as intriguing objects as the Jovian satellites. The Cassini-Huygens mission has unveiled the multivariable Earth-like geology of Titan since its arrivai in 2004. Furthermore, Titan is the only one planetary object except the Earth that possesses a unique nitrogen and full of organics atmosphère (e.g. Coustenis & Taylor, 2008). Recently, a subsurface liquid océan was suggested at Titan based on thermal and orbital calculations (e.g. Tobie et al. 2005) spin rate measurements and SAR reflectivity observations (Lorenz et al. 2008;[START_REF] Stiles | Erratum: "Determining Titan's Spin State from Cassini Radar Images[END_REF]Hussmann et al. 2010). The surface investigation brought to light many geological expressions such as extensive mountains, ridges, dendritic networks, dunes, lakes, channels, canyons and riverbeds (Lopes et al. 2010). Furthermore, there is the possible existence of active zones on the satellite due to past or recent cryovolcanic and tectonic activity (e.g. Soderblom et al. 2007;Lorenz et al. 2008;Nelson et al. 2009a;2009b).

Another Satumian satellite, Enceladus, has been shown by Cassini to be one of the most active objects in the Solar System despite its minor size. Indeed, jets from its Southern polar région consisted mainly of water ice particles that reach over 435 km in height, enrich the E-ring of Satum (Dougherty et al. 2006;Porco et al. 2006;[START_REF] Waite | Cassini Ion and Neutral Mass Spectrometer: Enceladus Plume Composition and Structure[END_REF]. The source of such activity could be a large internai liquid deposit, most likely an underground liquid océan (e.g. [START_REF] Porco | The restless world of Enceladus[END_REF]Collins et al. 2007;Postberg et al. 2009). The surface expressions observed on Enceladus include craters and smooth terrains that consist of extensive linear cracks, scarps, troughs and belts of grooves.

Due to their complex surfaces and intriguing interiors, both Satumian satellites, require new long-term missions with advanced instrumentation and possibly in situ éléments. Such a concept was proposed within the Titan Satum System Mission (TSSM), studied by ESA and NASA and could be launched after EJSM, around 2025.

The mission consists in an orbiter, a Titan montgolfière hot air balloon and a Titan lake lander for simultaneous in situ and remote exploration of Titan and Enceladus.

A number of évidences, either surficial expressions or geophysical factors, endorse the existence of internai océans beneath the four aforementioned satellites. The considération of the possible implications from geological features on the surface and their possible relation to the interior is presented in the review hereafter.

Geology: Surface features and their implications

The diversity ofgeologicalfeatures on Europa and Ganymede Without doubt, the surface expressions observed on a planetary body are the signatures of both extemal and internai processes as occurred with time. Hence, although the environmental conditions and the working material are different from the terrestrial case, in the presence of similarities in the surface features observed, the héritage of Earth science can be used as a tool for their study on other bodies.

Europa and Ganymede présent a major diversity in terms of appearance and surface geological structures and therefore in terms of the surface-shaping forces. Europa seems to be subject to active tectonism and cryovolcanism since it displays a young, smooth and active surface. On the contrary, Ganymede is heavily cratered on most of its surface and internai processes like cryovolcanism seem to hâve played only a minor rôle in the surface modification since there is little indication of resurfacing. In general, érosion as well as mass movement and landform dégradation seem to play an important rôle in resurfacing as it reduces the topographie relief by moving surface materials to a lower gravitational potential [START_REF] Moore | Erosion, Transportation, and Déposition on Outer Solar System Satellites: Landform Evolution Modeling Studies[END_REF].

The most distinct and characteristic morphotectonic features on Europa are the lineations that intersect the entire upper part of the satellite's crust (Fig. 1). These formations are probably the major resurfacing mechanism since their genesis is based on the intersection of any two parts of the surface by bedding or/and cleavage. Figure 2 displays a scheme of intersect lineae formation as seen on the Earth [START_REF] Park | Foundations of Structural Geology[END_REF]. The tectonic activity that forms these édifices along with the heating from the subsurface diapirs (mobile material that was forced into more brittle surrounding rocks and follows an upward direction) composes the dominant dynamic potential of the satellite. (Fig. 3). A distinct intersecting 1,500 km feature, Agenor Linea, is a candidate active région as found in photometric observations (Hoppa et al. 1998;[START_REF] Geissler | Evidence for non-synchronous rotation of Europa[END_REF].

Fig. 3. Four classes of linéaments on Europa [START_REF] Geissler | Evidence for non-synchronous rotation of Europa[END_REF].

Other geological expressions seen on Europa by Galileo are features that are called 'lenticulae' and 'chaos'. In terrestrial geology, a lenticulae feature is a depositional body that is thick in the middle and thin at the edges, resembling a convex lens in cross-section. Both features hâve probably formed by rising diapirs that produced partial melt. On Europa, the lenticulae are ovoidal features ranging from 5 to 20 km in diameter (Pappalardo et al. 1998) while the chaos structures, like Conamara Chaos (8°N, 274°W), are larger features presenting blocky material [START_REF] Carr | Evidence for a subsurface océan on Europa[END_REF][START_REF] Spaun | Conamara Chaos Région, Europa: Reconstruction of mobile polygonal ice blocks[END_REF]Sotin et al. 2002). Another hypothesis for the genesis of lenticulae features suggests a corrélation between them and the chaos formations. This hypothesis by [START_REF] Greenberg | Unmasking Europa In: The Search for Life on Jupiter's Océan Moon[END_REF] assumes that the lenticulae are chaos formations in smaller dimension that low-resolution Galileo imaging interpreted wrongfully as distinct structures. Chaos, area covered with big blocks of crust that are mixed and moved suggesting they floated on a liquid layer (Galilleo Project/NASA). The area of chaos terrain shows plate-like features (marked with A) with ridges and valleys, and régions lower than the plates (marked with B). Also, plate displacement is obvious with surface expressions like faults or ridges deviating from their linear structure (marked with C).

In contrast to Europa's flattened topography and homogeneous surface, Ganymede possesses two distinct types of terrain. The dominant terrain comprises the brighter régions marked with geological expressions such as extensive ridges, grooves and faults while its âge considers as the youngest in comparison with the rest of the surface [START_REF] Pappalardo | Grooved Terrain on Ganymede: First Results from Galileo High-Resolution Imaging[END_REF]) (Fig. 5). The other geological terrain-type is the dark terrain [START_REF] Prockter | Dark Terrain on Ganymede: Geological Mapping and Interprétation of Galileo Regio at High Resolution[END_REF]) (Fig. 6), which is the oldest one in terms of âge, heavily cratered with astrobiological interest due to the existence of organic materials (McCord et al. 1998). Considering the fact that the satellite is heavily cratered and the dark terrain comprises one of the oldest terrains of the Galilean satellites [START_REF] Zahnle | Cratering Rates on the Galilean Satellites[END_REF], this surface is an indicator of the system's cratering history [START_REF] Pappalardo | Grooved Terrain on Ganymede: First Results from Galileo High-Resolution Imaging[END_REF]. The régions that are particularly interesting in terms of geology are the transitional régions (Fig. 7). Such régions correspond to surface areas that constitute a transition from the dark terrain to the bright grooved terrain of Ganymede. Even though cratering is présent on both types of terrain, the dark ones seem heavily and more extensively bombarded [START_REF] Showman | The Galilean Satellites[END_REF] suggesting that they represent the oldest preserved (nonresurfaced) surfaces on Ganymede. [START_REF] Pappalardo | Geology of Ganymede[END_REF] suggested that the modification of the dark terrain's material due to tectonic and cryovolcanic resurfacing formed the primary bright terrain. Another significant région on Ganymede is the Galileo Regio (Fig. 8), which has likely been formed during an active geologically period [START_REF] Casacchia | Géologie évolution of Galileo Regio, Ganymede[END_REF]. The Galileo Regio is a heavily cratered area but not an impact crater. It seems to hâve been shaped under the influence of tectonic processes and young and bright material that arose from the interior [START_REF] Harland | Jupiter Odyssey: The Story of NASA's Galileo Mission[END_REF]. Active worlds: Titan and Enceladus and their habitability Even though Titan and Enceladus' surface expressions are very different in terms of composition, materials and size, they highly resemble the Earth's geomorphology. Titan's surface consists of structures like mountains, ridges, faults and canyons (Fig. 9), formed most probably by tectonic processes, as discovered by the Cassini-Huygens mission (Brown et al. 2009;Lopes et al. 2010;Mitri et al. 2010). Titan, other than its atmospheric uniqueness, is also the only among outer planet satellites where aeolian and fluvial processes operate to erode, transport, and deposit material [START_REF] Moore | Erosion, Transportation, and Déposition on Outer Solar System Satellites: Landform Evolution Modeling Studies[END_REF]. The phenomenon that most probably formed terrestrial-like volcanic structures like calderas, flows and dômes on Titan, is cryovolcanism. Currently, there are three possible cryovolcanic régions. These are Tui Regio (20°S, 130°W), Hotei Regio (26°S, 78°W) and Sotra Facula (15°S, 40°W). The latter is considered as the most promising one with obvious surface expressions of peaks, flows and one caldera structure (Fig. 10). Furthermore, aeolian and fluvial processes acting on Titan's surface create édifices such as lakes, seas, riverbeds, sand dunes, shorelines, and dendritic drainage networks (Fig. 11). Extemal impact phenomena like impact craters are rarely observed on the surface due to the resurface activity. Similarly to Titan, Enceladus présents major tectonic features and active cryovolcanism. The most fascinating phenomenon occurs on Enceladus currently due to its tremendous internai dynamic convective forces that cause Geyser-like fountains at its Southern pôle that could reach more than 400-km in height (Porco et al. 2005). The rest of Enceladus' surface is covered by smooth and cratered terrains, rifts, ridges, grooves, escarpments and extensive linear fractures (Johnson, 2004). The geology of this tectonized moon is a field of active scientific research awaiting for new observations. Up to date, the observations and analysis showed two types of tectonic terrains. The north pôle consists of heavily cratered landforms while the central région and Southern pôle of tectonic molded terrain with cryovolcanic features. 3. Interior models and liquid water subsurface océans in giant pianets' satellites Our current knowledge of the icy moons' internai stratification and their composition is being built on a combination of spacecrafts data, laboratory experiments, and theoretical geophysical modeling. Resembling

Earth's moon in terms of structure, icy moons consist of a core, a mantle, and a crust, with the specificity of the existence of a liquid océan lying within the icy mantle (Fig. 13).

Fig. 13. From left to right: Europa, Ganymede, Titan and Enceladus' internai stratigraphie models (NASA/JPL).

According to current models of internai structure, the existence of subsurface océans is expected for most of the icy moons of the Outer pianets (e.g. Sohl et al. 2010;Schubert et al. 2010 and references therein). Even if an océan is not currently hidden within the interior, it is suggested that a liquid layer was présent in the past but cooled to ice over time. An example of such case is Neptune's moon Triton.

Evidence for hydrated sulfate salts on the surfaces of Europa and Ganymede from spectroscopic data support the possible existence of subsurface océans [START_REF] Mccord | Non-waterice constituents in the surface material of the icy Galilean satellites from the Galileo near-infrared mapping spectrometer investigation[END_REF]2001;Cassidy et al. 2010;[START_REF] Fortes | Phase Behaviour of Ices and Hydrates[END_REF] and references therein) suggesting the déposition of minerais following internai hydrothermal events. In addition, Galileo's magnetometer (e.g. Khurana et al. 1998;Kivelson et al. 2002), detected induced magnetic fields at Europa and Ganymede that imply the presence of an electrically conductive subsurface layer (e.g. Sohl et al. 2010). Furthermore, the détection of a low viscosity layer undemeath the icy crust again endorses the presence of a subsurface liquid océan inducing recent geological activity [START_REF] Stem | Triton's surface âge and impactor population revisited (evidence for an internai océan)[END_REF][START_REF] Ruiz | Seas under ice: Stability of liquid-water océans within icy worlds[END_REF].

The common properties that need to be satisfied on ail bodies in order to sustain a liquid subsurface océan are:

(a) Heat production which mainly originates from radiogenic heating or other triggering mechanisms (e.g. , 1999;Tobie et al. 2005). In the absence of an internai liquid layer, the tidal deformation of the ice shell remains very small as there is no découplé of the core and the mantle, thus the resulting stress and conséquent heating is negligible [START_REF] Moore | The tidal response of Europa[END_REF]. Other possible heat sources are the dissipation of tidal energy due to orbital interaction between the satellites and their pianets, or the exothermal geochemical production of heat like hydration and crystallization of solids (e.g. Sohl et al. 2010;Hussmann et al. 2010).

McKinnon

(b) Efficiency of heat transfer, which is based on thermal diffusion and thermal convection (e.g. Hussmann et al.

2010).

(c) Components capable of decreasing the melting point of ice and supporting the ocean's liquid State (e.g. Sohl et al. 2010;[START_REF] Tobie | Surface, Subsurface and Atmosphère Exchanges on the Satellites of the Outer Solar System[END_REF]. The composition of such océans should offer an antifreeze constituent like a solution of water with ammonia, so that it remains in liquid State. The aqueous évolution of a possible internai océan dépends on the several Chemical interactions between liquid water and rocks, on the hydrodynamic processes within the océan (degassing), on the freezing of the expected plumes as well as on the presence of secondary organic and inorganic species (Sohl et al. 2010).

(d) Stability of the crust against convection, keeping the océan subsurficial as well as preventing stagnant lid convection (e.g. [START_REF] Mckinnon | Geodynamics oflcy Satellites[END_REF][START_REF] Rainey | Océans in ice-rock bodies: conditions for the existence of subsurface liquid water[END_REF].

The Voyager and Galileo missions data suggest that the internai stratification of Europa consists of an iron core (300 to 800 km) covered by a rocky mantle probably 100 km thick that it is overlain by a liquid water océan less than ten kilometers thick. The océan is covered by a surficial icy layer possibly 15 km thick (Fig. 13). The Galileo magnetometer's measurements that showed the crust had shifted by almost 80°, also indicated that the mantle is not attached to the crust, thus reinforcing the theory of an internai océan (Kivelson et al. 2000;[START_REF] Greenberg | Europa: The Océan Moon: Search for an Alien Biosphère[END_REF]. Nevertheless, the most convincing evidence for the liquid water océan is the existence of the controversial 'chaos terrain' (Fig. 4) that possibly formed by melt-through from below (O'brien et al. 2002). The controversy debates on the mechanisms that formed the terrain, whether through induced impacts or cryovolcanic processes, as well as on the thickness and State of the ice shell. On one hand, one model (Fig. 14a)

suggests that the océan is a warm convective ice layer located several kilometers below the icy crust and on the other hand the other model (Fig. 14b) suggests a liquid océan hidden more than 100 km below the crust. The radioactive decay cannot provide the amount of heat required to modify entirely the satellite surface, as observed on Europa. In this case, the surface température is limited to 110K at the equator and 50K at the pôles; such cold températures make the ice locally as hard as terrestrial igneous granité [START_REF] Mcfadden | The Encyclopedia of the Solar System[END_REF]). On the other hand, hydrothermal activity has the potential to reshape the surface crust. Recent studies suggested that the influence of Jupiter on Europa due to its small but non-zero obliquity probably generates large tidal waves that keep the océan warm (Tyler, 2008).

If the heat propagation and the buoyant oceanic currents are not intense then the ice shell will be thick and a warm ice layer will be formed at the bottom of the shell (Fig. 15b). During the hydrothermal processes this warmer ice will rise and slide like the terrestrial glaciers. Such movements can cause surface modification and produce structures like the 'chaos terrain'. Other features supporting the existence of a thick ice layer are the large impact craters surrounded by concentric rings filled with ice. Geophysical models, associating the mere presence of these structures to the amount of heat generated by the tides, suggest an icy crust 10-30 km thick with a warm ice layer at the bottom and a liquid océan probably 100 km thick (Schenk et al. 2004) Altematively to the previous model, if the heat flow and the plumes are intense then the ice shell is expected to be thin (Fig. 15a). The fragmentation of such a thin crust is most probably expressed with tectonic-like formations like the Conamara région. The debate regarding the thickness of Europa's crust gave birth to an alternative hypothesis regarding the thin ice model, which suggests a thin upper crust layer, almost 200-m thick, that behaves elastically and is in contact with the surface through zones of weakness like multiple linear ridges [START_REF] Billings | The great thickness debate: Ice shell thickness models for Europa and comparisons with estimâtes based on flexure at ridges[END_REF]. Europa's case supports the existence of a stagnant lid undemeath its crust (Fig. 16) [START_REF] Showman | Numerical simulations of convection in Europa's ice shell: Implications for surface features[END_REF].

The stagnant lid is a relatively cold and stiff conductive layer covering the warmer convective icy interior [START_REF] Schubert | Interior composition, structure and dynamics of the Galilean satellites[END_REF]). Since the expected activity within Europa's interior is upwelling thermal diapirism, the stagnant lid can prevent cold near-surface icy material from sinking towards the océan. Température Fig. 16. Simulation of convection within Europa's ice shell [START_REF] Showman | Numerical simulations of convection in Europa's ice shell: Implications for surface features[END_REF].

Compared to the Earth's océans, the composition of Europa's hidden océan should be significantly different. The Earth's océan major component is sodium chloride while Europa's should be magnésium sulfate as indicated by Galileo data (e.g. Fanale et al. 2001). Indeed, Europa's weak magnetic moment is induced by the varying part of the Jovian magnetosphere [START_REF] Schilling | Time varying interaction of Europa with the Jovian magnetosphere: Constraints on the conductivity of Europa's subsurface océan[END_REF]) and requires a highly conductive subsurface océan. Such conductive materials candidates are the magnésium sulfate (e.g. [START_REF] Mccord | Non-waterice constituents in the surface material of the icy Galilean satellites from the Galileo near-infrared mapping spectrometer investigation[END_REF] or sulfuric acid hydrate (Carlson et al. 2005).

Similarly, Ganymede consists of a four-layer interior, based on measurements regarding its gravity field, mass, density and size. Geophysical models suggest two kinds of internai stratigraphy, one of an undifferentiated mixture of rock and ice and one of a differentiated body consisting of a rocky core, an extended icy mantle and an icy crust. We think the latter model is the most plausible as it is compatible with the gravity field measurements made by Galileo (e.g. [START_REF] Sohl | Implications from Galileo Observations on the Interior Structure and Chemistry of the Galilean Satellites[END_REF]. Such intrinsic magnetic field supports the existence of an iron-rich core [START_REF] Hauk | Sulfur's impact on core évolution and magnetic field génération on Ganymede[END_REF]. The thickness of the layers in the interior dépends on both the fraction of olivine and pyroxene and the amount of sulfur in the core (e.g. [START_REF] Sohl | Implications from Galileo Observations on the Interior Structure and Chemistry of the Galilean Satellites[END_REF]) (Fig. 17). Thus, the 2,634-kmradius Ganymede consists of an 800 km iron sulfide core, an almost 900-km thick outer ice mantle and a silicaterich mantle 700 km thick (e.g. [START_REF] Kuskov | Internai structure of Europa and Callisto[END_REF]. Although the mechanisms that formed Ganymede's complex surface are still unidentified, multiple scénarios hâve been proposed. Most of the scénarios agréé on a general mechanism that uses tectonism as the main chisel that formed and shaped the existing structures, especially the grooved terrain [START_REF] Sohl | Implications from Galileo Observations on the Interior Structure and Chemistry of the Galilean Satellites[END_REF]. On the other hand, and in contrast with the other three aforementioned satellites, cryovolcanism does not seem to participate actively in the geodynamic processes. The radiogenic heating from within the satellite as well as tidal heating from past events are considered as the main forces that generate the stresses that lead to tectonic movements and eventually to tectonic structures.

Contrary to Ganymede, probably Titan but especially Enceladus provide evidence of past and current cryovolcanism that shapes their surfaces. The mechanisms that formed Titan's surface by endogenic factors are still unknown, although the central idea is focused on cryovolcanism and morphotectonism, with the latter being the short-and long-term surficial expressions of any tectonic activity originated from endogenic processes (Solomonidou et al. 2010). However, it is possible that Titan underwent a period of tectonism resembling those on Europa's and Ganymede's.

According to geophysical models, Titan's differentiated interior consists of a serpentinite core (-1,800 km), a high-pressure ice mantle (-400 km), a liquid layer of aqueous ammonium sulphate (50 to 150 km wide), and an extemally heterogeneous icy few kilometers wide (Tobie et al. 2005;Fortes et al. 2007).

The subsurface instability due to the interactions within an interior liquid océan causes the modification of extended features on Titan's surface, whether they dérivé from cryovolcanic or morphotectonic dynamic processes. Currently, ail the geophysical models that try to explain the geodynamics of Titan support the existence of an oceanic layer that découplés the mantle from the icy crust. Additionally, the identification of a small but significant asynchronicity in Titan's rotation from Cassini SAR data favors the aforementioned decoupling (Lorenz et al. 2008). Internai geodynamic activity can transport effusively the explosive material from the oceanic layer to the surface and form the cryovolcanic structures like the lobate flows in Sotra Facula. r \>yN. High pressure ices CZ3ke I Methare daîhrate Tobie et al. (2006) suggested a cryovolcanic model for Titan's case to solve the mystery of Titan's methane replenishment since it is should vanish in 100 Ma. These authors suggested that episodic methane outgassing events occurred through three distinct épisodes covering a chronological period from 2000 Ma ago until 500 Ma ago (Fig. 18). The internai océan provides the 'magma' chamber by means of material, while convective processes are the triggering mechanisms that initiate the dynamic activity. A convective model of a stagnant lid is capable to explain such activities [START_REF] Solomatov | Scaling of température-and stress-dependent viscosity convection[END_REF] as explained for Europa's case earlier. However, the ice shell for Titan's case is expected to be thicker. In opposition, several studies suggest that Titan is not currently convective (Mitri et al. 2010;[START_REF] Nimmo | Shell thickness variations and the long-wavelength topography of Titan[END_REF].

The uniqueness of Titan's tectonism -even though not yet confirmed -lies in that the tectonic processes are contractional rather than extensional, setting Titan out as the only planetary body in the Solar System other than the Earth, where contractional deformation occurs. Mitri et al. (2010) proposed an internai thermal model, focused on the changes in volume of a potential underground océan caused by heat flux variations during freezing or melting. The authors suggest that the continuing cooling of the moon can develop global volume contraction, as described by Tobie et al. (2005;2006).

Cassini's data proved that, despite its small size (about 505 km in diameter), Enceladus is an active planetary body that spews material through hydrothermal vents resembling terrestrial geysers. This moon is not in hydrostatic equilibrium (Schubert et al. 2007;Schubert et al. 2010), thus a simple and very general stratigraphie interior is being suggested: it consists of a 169 km rocky core overlain by an icy 82 km mantle [START_REF] Barr | Convection in Enceladus' ice shell: Conditions for initiation[END_REF]Fortes, 2007;Schubert et al. 2007). The subsurface océan is lying between the two layers, and supplies the fountains observed at the south polar région through cracks called 'Tiger Stripes' (e.g. [START_REF] Porco | The restless world of Enceladus[END_REF]Postberg et al. 2009). The moon's tidal dissipation is proposed as the triggering mechanism as well as the cause of the dynamic phenomena like tectonics (convection-conduction, expansion contraction) and cryovolcanism (e.g. Mitri and Showman, 2008;Mitri et al. 2010) that formed the surface expressions described in the previous section.

Discussion: Internai activity and surface modification

The possible existence of subsurface liquid océans undemeath the crusts of the icy moons of Jupiter (Europa and Ganymede) and Satum (Titan and Enceladus), places them in a potential group of planetary bodies where life could emerge and evolve. Other than data processing that provide evidence and information about the internai liquid layers, the surfïcial expressions that are related to the hydrothermal and dynamic processes occurring within these layers are the surface evidence that could lead to their identification. Specifically, the cryovolcanic and morphotectonic structures seen on the aforementioned satellites are the surface expressions of the internai activity while they are formed by modification of the crustal layer and déposition of material coming from the subsurface océan. Therefore, investigating these surface exposures and associating them to terrestrial features where water is involved could shed some light on the investigation of internai liquid water océans in the icy moons. Trying to model the triggers of internai active phenomena, basic geophysical models usually propose liquid water réservoirs while the geodynamic models point at both radioactive decay and tidal stresses, caused by the giant planets Jupiter and Satum. We shall herein try and reconcile both views.

The major and most significant structures on Europa's surface are the bunch lineae (e.g. [START_REF] Prockter | Characteristics of Icy Surfaces[END_REF] (Fig. 3). Data analysis showed that the geometry and the spectral properties vary depending on âge, which indicates an evolutionary sequence [START_REF] Geissler | Evidence for non-synchronous rotation of Europa[END_REF]Dalton, 2010). This means that persistent and drastic processes occur in order to form these features. Such processes seem to be different types of cryovolcanism in which multiple éruptions of ice (warmer than crustal ice) emerge through 'tectonic' crustal weaknesses (Figueredo & Greeley, 2004). This tectonic formation resembles the terrestrial mid-ocean ridges (MOR), which are the extensive opening seafloor terrains, and considered to be a global rather than a local phenomenon. MOR are dynamic and volcanically active structures that constantly provide and deposit new material from the mantle to create new oceanic crust. As seen in Figure 4, the red streaks, as well as the red spots called lenticulae, are most possibly evidence of upwelling warmer material emerging from the liquid layer while colder ice near the surface sinks downwards. After spectroscopic analysis, the red streaks are thought to be rich in magnésium sulfate, another hint in favor of their internai origin [START_REF] Mccord | Non-waterice constituents in the surface material of the icy Galilean satellites from the Galileo near-infrared mapping spectrometer investigation[END_REF]Dalton et al. 2010). Furthermore, the linéament formations like the dark and red streaks présent the basic resurfacing System that dominate on Europa's surface. Their formation is an indicator for current geological activity. Additionally, the thermal diapirism that most likely formed these structures, as well as the red spots, would imply that convective upwelling thermal plumes originate in the lower boundary of the convective System and gets in contact with the cold stagnant lid with the icy plumes [START_REF] Showman | Numerical simulations of convection in Europa's ice shell: Implications for surface features[END_REF]). Such température ranges indicate possible habitability at the upper part of the ice shell. Consequently, Europa's surface morphology is directly connected to the internai dynamic processes and provides evidence of a subsurface interactive océan.

Unlike Europa's conflicting oceanic hydrothermal System that originates in a rocky sea layer, Ganymede's océan lays between two icy layers that découplé it from the mantle. Barr et al. (2001;2004) suggested that large magmatic events due to convective plumes could occur at Ganymede's rock -ice boundary. The surface expressions that are most possibly connected to tectonism are the deep fault structures like the horst-and-graben (resembling terrestrial continental rifts) as well as many cracks that are observed in Ganymede's bright terrain [START_REF] Showman | The Galilean Satellites[END_REF]. Tidal heating events, either past or current, could cause dynamic forces that modify the icy lithosphère. The main mechanism that deforms tectonically the lithosphère is most likely warm liquid plumes that rise from the upper mantle to the surface, following a pattern similar to the plume-lithosphere interactions at the Hawaiian Swell which cause thinning and instabilities at the crust layer [START_REF] Moore | Three-Dimensional Simulations of Plume-Lithosphere Interaction at the Hawaiian Swell[END_REF].

Even though the plume theory indicates cryovolcanic processes, little evidence of such activity has been detected on Ganymede. Notably, the extreme morphological différence between the bright and the dark terrain as described in a previous section suggests a massive geological event or set of events that caused such large-scale geo-terrain alteration. Hence, it is possible that the extensive tectonic forces that fractured the dark terrain partially affected the bright terrain as well. Such tectonic weaknesses could display pathways to small cryovolcanic events that lead to resurface processes. However, tilt-block faulting and shears [START_REF] Head | Evidence for Europa-like tectonic resurfacing styles on Ganymede[END_REF] are some of the structures that appear bright within the dark terrain suggesting that tectonic cracks functioned as path for warm internai oceanic material to pass through like what occurred in the bright terrain. In terms of tectonics, and similarly to the other icy moons, there is no evidence of compressional deformation [START_REF] Showman | Tidal Evolution into the Laplace Résonance and the Resurfacing of Ganymede[END_REF]. Since the deformational pattern is extensional, the question is whether it is a global or a local phenomenon. [START_REF] Collins | Formation of Ganymede grooved terrain by sequential extensional épisodes: Implications of Galileo observations for régional stratigraphy[END_REF] studied observations of grooved terrains of spécifie stratigraphie âges that hâve consistent directions over hundreds of kilometers, something that indicates global stressing phenomena. On

Earth stressing phenomena could occur where severe forces cause convective currents in the océan. In a similar way, the plume convection within Ganymede's oceanic layer could create enough turbulence and temperaturepressure instabilities to cause global stressing phenomena with an impact on tectonism. Nevertheless, Ganymede présents styles of tectonism different from Europa's.

Following a pattern similar to the one mentioned for Ganymede, Enceladus' internai tidal stresses and radioenergetic decay produce warm pockets of material, the plumes, that subsequently form the geyser formations at the south polar région. The major and most valuable evidence of Enceladus' cryovolcanic activity, supporting as well the océan existence, is this geyser formation or geyser accumulation that form a fountain of more than 400 km, as observed by Cassini. The jets initiate from four sub-parallel linear dépréssions, which are tectonic in origin. Other surface expressions are scarps, ridges, and shields (Collins et al. 2009). On the other hand, Titan's surface structures related to its dynamic interior that also support the existence of a subsurface océan are the three cryovolcanic candidate régions Tui Regio, Hotei Regio and Sotra Facula and many morphotectonic structures like mountains, ridges and canyons (e.g. Solomonidou et al. 2010 and references therein). Generally, in the case of both moons it is thought that their ice shells transit from a conductive to a convective State; since they probably overlay a pure liquid océan (Tobie et al. 2005) this can hâve major effects on surface morphotectonics (Mitri and Showman, 2008). Indeed, thermodynamic oscillations within the ice shells may trigger repeated extensional and compressional events. In the presence of a subsurface océan undemeath Titan and as a conséquence of local stress mechanisms, parts of the icy crust could behave like rigid ice floes due to latéral pressure gradients. If such floating occurs, many morphotectonic features like faults and canyons can relate their formation to this event. Furthermore, radial contraction of the internai high-pressure ice polymorphs could possibly amend the radial expansion caused during the cooling stage of the moon (Mitri and Showman, 2008), in which the existence of a liquid layer plays a significant rôle. As a resuit, the overall global contraction could form mountainous chains (Radebaugh et al. 2007).

These four satellites are dynamic planetary bodies and the internai océan most likely plays an important rôle on the formation and modification of their surfaces. The properties on which the presence of an internai océan dépends were mentioned earlier (section 3). Another issue that should be taken under considération is the set of properties that détermine the possible exchange of material between the subsurface and the surface. These are the mechanical properties of the lithosphère as proposed by [START_REF] Tobie | Surface, Subsurface and Atmosphère Exchanges on the Satellites of the Outer Solar System[END_REF].

The formation of the surface signature of any upwelling activity dépends on:

(a) the Chemical interaction between the material in motion and the local environment, which travels through the conduit path from the source to the surface.

(b) the magnitude of the forces that triggers this motion.

(c) the complexity of the structure of the conduit path which could be an extensive tectonic zone of weakness.

From such a structure, multiple cryovolcanic éruptions could emanate.

(d) the influence of the atmosphère on the surface as described in [START_REF] Tobie | Surface, Subsurface and Atmosphère Exchanges on the Satellites of the Outer Solar System[END_REF].

Titan and Enceladus display cryovolcanic expressions in spécifie régions. On Titan these zones appear on a latitudinal ring around 20°S -30°S, while on Enceladus they lay at the southem pôle. On the contrary, Europa and Ganymede do not host evidence of such large and localized structures. Considering these surface expressions, as well as the morphotectonic structures described earlier, we infer that different endogenic conditions occurred. Firstly, the upwelling material within Europa and Ganymede could spread throughout the lithosphère, while in Titan and Enceladus it would pass through more localize exsolution paths. On Titan and especially on Enceladus the cryovolcanic expressions illustrate instant energy relief of the hydrodynamic activity, while Europa and Ganymede could expérience continuous relaxation. This can possibly explain the fractal development of lineae structures on Europa. Moreover, the gravitational field, as well as other heat and transfer mechanisms, play in each satellite a major rôle in the distribution of the upwelling material and the formation of the surface structures. In order to evaluate the above implications it is essential to record the spatial and temporal variations of the structures observed.

Habitability issues for outer planet satellites

The existence of an internai liquid océan undemeath the icy crusts of the giant planet satellites could serve as a potential abode for life. The location of the océan close to the surface provides food for thought on habitability zones, and conditions for life in general, in the Solar System.

Indeed, the discovery of hydrocarbon surficial lakes on Titan and the possible existence of subsurface liquid océans in Europa, Ganymede, Titan, and Enceladus reveal alternative concepts to the classical définition of the habitable zone, and suggest the need for reconsidering its limits (Lammer et al. 2009). Currently, more and more studies regarding the planetary habitability propose the icy moons with subsurface océans as potential worlds for initiating and/or sustaining some sort of life forms (Fortes, 2000;Raulin, 2008;[START_REF] Raulin | Exobiology and Planetary Protection of icy moons[END_REF]Coustenis et al. 2011 and référencés therein). Figure 19 présents the possible locations of the liquid layers within the icy satellites of Jupiter as well as their habitability potential. For Europa, the assumed internai processes that trigger the geodynamic phenomena are the tidal stresses, which produce considerably less energy than radioenergetic decay. Europa is being compressed and stretched when affected by the gravitational forces of Jupiter and the other Galilean satellites and thus, the resulting tidal friction provides enough energy, as well as température, for an extensive océan to exist undemeath the crust. The provision of large amounts of energy at the bottom layer of an océan can conceivably form hydrothermal vents like the ones seen on Earth. Notably, life could emerge around hydrothermal vents, which are important geological forms in terms of life propagation. [START_REF] Greenberg | Transport Rates of Radiolytic Substances into Europa's Océan: Implications for the Potential Origin and Maintenance of Life[END_REF] suggests that even ecology including complex organisms could exist on Europa. The author provided evidence of oxygen concentration within the océan greater than that of the Earth's, which is suggested as an indicator for aérobic organisms. Nevertheless in the case where the concentration of salts is large, only extremophile organisms (like halophiles) could survive (Cooper et al. 2001;Marion et al. 2003).

Ganymede lays between structures 1 and 2 in the scheme shown on Fig. 19 indicating that it is much colder than Europa, a factor that lowers its habitability potential. On the other hand, reinforcing Ganymede's possibility for life to exist, Barr et al. (2001) imply that magmatic events could form pockets of liquid, which would then ascent buoyantly carrying nutrients to the internai océan. Based on their calculations, a water plume could reach Ganymede's océan and carry nutrient-rich material with an éruption time of 3 hours to 16 days. Additionally, Ganymede's System could provide the necessary tools to concentrate biological building block ingrédients (Trinks et al. 2005) especially since it possesses a magnetic field that is able to protect life from harmful radiation and lies in a relatively quiet radio zone.

On Titan, terrestrial bacteria can absorb their energy and carbon needs from the tholins that exist in its thick atmosphère (Stoker et al. 1990). Furthermore, photochemically derived sources of free energy on the surface could maintain an exotic type of life, using liquid hydrocarbons as solvents (McKay & Smith, 2005). Other than the atmospheric properties that are favorable to life, the possible existence of an underground océan might support terrestrial-type life that had been introduced previously or formed when liquid water was in contact with silicates early in Titan's history (Fortes, 2000). Furthermore, the possible amount of ammonia dissolved within the océan is suggested to be ~10% (Lorenz et al. 2008), something that corresponds to a pH of 11.5, while at a depth of 200 km the pressure reaches ~5 kbars and hot plumes within the potential océan could be generated (Coustenis et al. 2011 and references therein). The aforementioned conditions are not unfavorable to the emergence and maintenance of life (e.g. Fortes, 2000;Raulin, 2008).

Enceladus' extremely high températures at the south polar région are probably generated by the hydrodynamic processes that form the fountain, as previously described, and thus enhance the potential for habitability. The most convincing theory, after Cassini data analysis, suggests that a liquid océan exists beneath the Tiger Stripes.

There are standards for life that Enceladus' possible océan is not consistent with: the sunlight, the oxygen compounds, and the organics produced on a surficial-crust environment. However, terrestrial régions like the deep océans that do no satisfy the aforementioned prerequisites for life, still function as active ecosystems (McKay et al., 2008;Muyzer and Stams, 2008). There, sulfur-reducing bacteria consume hydrogen and sulfate, produced by radioactive decay. Notably, the comparison with the terrestrial ecosystems suggests that plume's methane may be biological in origin (McKay et al., 2008). Hence, Enceladus displays an internai warm and chemically rich océan that may facilitate complex organic chemistry and biological processes (Coustenis et al. 2011).

In conclusion, the confirmation of the existence of a subsurface liquid océan undemeath the crust of the icy moons of Jupiter and Satum will revolutionise our perspective regarding the habitability potentials of such planetary bodies. Indeed, liquid water may exist well outside the traditional habitability zone, which is merely based on the presence of liquid water on the surface. For planetary habitability, the principal criteria are the presence of liquid water anywhere on the body, as well as the existence of environments able to assembly complex organic molécules and provide energy sources: this can well be undemeath the surface in some cases, if stability conditions are met. The four satellites described in this study seem to ftilfill some or ail of the above requirements. However, it is of high priority to revisit these bodies with new missions and advanced instrumentation (such as gravitational and magnetic field sounding Systems and in situ element detectors) in order to obtain altimetry and in situ monitoring of tidally-induced surface distortion data (Sohl et al. 2010) that could unveil in detail the internai stratigraphy of the moons and the specificity of the subsurface océans. Future large missions, or smaller dedicated ones, to the Galilean and Khronian Systems would allow us to better understand the mechanics behind the astrobiological potential of worlds with subsurface océans, and shed some light on the emergence of life on our own planet.

carbon, nitrogen, oxygen, sulfur and their compounds, as well as of the D/H, 4He/3He, 22Ne/21Ne/20Ne, 36Ar/38Ar, 13C/12C, 15N/14N, l80/(170)/160, 13c,Xe/134Xe/l32Xe/l30Xe/129Xe isotopic ratios will be measured by mass Giant planets are mostly made of (lie gas that was présent in the protosolar disk before the terrestrial planets aecreted (Fig. 1). Their comparative study is thus essential to understand planet formation in general and the origin of the Solar System. Satum in particular appears to be a natural target for near-future exploration, after the fine characterization of Jupiter by Galileo and Juno and before future ambitious missions to Uranus and Neptune. Satum, the ring planet, is mysterious in many aspects, and plays a key rôle to undcrstanding planet formation, the évolution of solar and extrasolar giant planets, planetary meteorology, magnetospheric interactions, dynamo génération and the physics of planetary rings.

Satum, like Jupiter, lias an atmosphère that appears to be enriched in éléments other than hydrogen and hélium with respect to the solar composition. This enrichment may be the resuit of planetary precursors formed at low températures, or of a progressive enrichment of the protosolar disk, with profound conséquences for understanding the formation of the Solar System. The different formation scénarios that resuit can be disentangled by a study of the atmospheric composition of Satum in noble gases and major volatile (H, C, N, S, O) elemental and isotopic compositions, which requires in situ measurements Table 1.

Noble gases are particularly relevant in this study, because: (1) they are chemically ineit and their abundances ai-e determined by physical processes such as phase partitioning, and ( 2) their isotopic compositions présent large-scale inliomogeneities between the original protosolar nebula composition and subreservoirs of the solar System that make theni isotopic tracers as well. hideed, analyses of meteorites point to the existence of a primordial réservoir of volatiles Giant planets are mostly raade of the gas thaï was présent in the protosolar disk before the terrestrial planets aecreted (Fig. 1). Their comparative study is thus essential to understand planet formation in general and the origin of the Solar System. Satuni in particular appears to be a natural target for near-future exploration, after the fine characterization of Jupiter by Galileo and Juno and before future ambitious missions to Uranus and Neptune. Satum, the ring planet, is mysterious in many aspects, and plays a key rôle to understanding planet formation, the évolution of solar and extrasolar giant planets, planetary meteorology, magnetospheric interactions, dynamo génération and tlie physics of planetary rings.

Satum, like Jupiter, lias an atmosphère that appears to be enriched in éléments other than hydrogen and hélium witli respect to the solar composition. This enrichment may be the resuit of planetary precursors formed at low températures, or of a progressive enricliment of the protosolar disk, with profound conséquences for understanding the formation of the Solar System. The different formation scénarios that resuit can be disentangled by a study of the atmospheric composition of Satum in noble gases and major volatile (H, C, N, S, O) elemental and isotopic compositions, which requires in situ measurements Table 1.

Noble gases are particularly relevant in this study, because: (1) they are chemically inert and their abundances are detennined by physical processes such as phase partitioning, and ( 2) their isotopic compositions présent large-scale inhomogeneities between the original protosolar nebula composition and subreservoirs of the solar System that make them isotopic tracers as well. hideed, analyses of meteorites point to the existence of a primordial réservoir of volatiles with an isotopic composition different ffom that of the Sun. The discovery of this réservoir through its isotopic signature in Satum's atmosphère would hâve a direct impact on the primordial history of the Solar System and on the study of meteorites and cornets. Generally, the elemental and isotopic déterminations of Satum's atmospheric composition would permit to explore sources of matter and processes of formation for the giant planets. Another surprising fmding of the Cassini-Huygens mission has been the fact that Satum's magnetic field is considerably modificd by the rings, with two conséquences:

(1) our inability to measure the deep rotation rate in the magnetic dynamo région; (2) the impossibility to measure the planet's true magnetic field, but instead, a filtered, perfectly asymmetric field. A measurement of tlie planetary magnetic field insidc tlie D ring would provide the data necessary to understand the planet's magnetic field génération and to analyze gravity data using proper constraints on the rotation rate.

Tlie next section discusses the primaiy scientific goals that are driving the mission, namely the study of tlie origin of the Solar System and tlie study of planetary atmosphères, as well as the secondary goals which require a feasibility assessment, and fmally the physics of magnetic fields and of planetary rings.

2 Science objectives 2.1 Planet formation and the origin of the solar System Satum and Jupiter formed 4.55 Ga ago, from the same disk of gas and solids that formed the Sun and eventually tlie entire Solar System (e.g., [1]). A significant fraction of their mass is composed of hydrogen and hélium, the two lightest and most abundant éléments in the Universe. Disks with hydrogen and hélium are almost ubiquitous when stars appear, but these disks fade away rapidly, on timescales of only a few million years (e.g., [2]). This implies tliat Jupiter and Satum had to foim rapidly in order to capture their hydrogen and hélium envelopes, more rapidly than e.

g. terrestrial planets which took tens of millions of years to attain their présent masses, and retained only negligible amounts of tlie primordial gases as part of their final composition. Thus by studying these giant planets, we hâve access to information on the composition and early évolution of the protosolar disk that led to the birth of tlie entire Solar System.

In spite of tlie récognition of tlie importance of sucli knowledge, data on the composition and structure of tlie giant planets, which hold more than 95% of tlie mass of the Solar System, excluding tlie Sun, remains scarce. The masses of £) Springer the central cores of Jupiter (0 to 15 M© where M© is the mass of the Earth) and of Satum (8 to 26 M©), as well as the total masses of heavy éléments (10 to 40 M©, and 20 to 30 M©, respectively) are poorly constrained [3]. Along the same lines, the abundance of oxygen, the third most abundant element in the Universe after H and He and a key element for planetary formation, is unknown in the well-mixed atmosphères of both Jupiter and Satum [4]. Jupiter's composition was measured in situ by the Galileo probe in 1995 down to a pressure level of 22 Bar [5]. On Jupiter, the abundances of 8 major éléments hâve been measured, while in contrast, on Satum we only hâve reliable data for C, He, and model-dependant results on N and S (Fig. 2). The Galileo measurements at Jupiter include a highly précisé détermination of the planet's hélium abundance, crucial for calculations of the structure and évolution of the planet.

Several key species (carbon, nitrogen, sulfur, argon, krypton and xénon), were found to be enriched in Jupiter's atmosphère compared to the solar composition mixture, which directly impacts théories on the formation of the Solar System.

Figure 2 shows that planet formation models that attempt to reproduce the abundances measured in Jupiter yield very different results for Satum, with key éléments being oxygen (in tire fomi of water) and the noble gases. Specifically a model based on direct adsorption on low-température planetesimals yields a high, uniform enrichment of ail species other than H, He and Ne [5]. A model in which gases are trapped into crystalline ice as clathrates yields a very non-uniform enrichment of these species (e.g. [7,8]). If the noble gases were not delivered directly with the solids but as gas in a protosolar disk enriched in heavy éléments by photoevaporation, one then expects enricliments in noble gases Üiat are decoupled from other speeies and comparable to thosc found on Jupiter [3,9]. The processes tliat formed the giant planets hâve implications for the other planets as well.

Kronos can tlrerefore directly test scénarios of the formation of the solar System.

In paiticular, Jupiter 's nitrogen is highly depleted in 15N by up to 40% compared to "planetary" N found in terrestrial planets and mefeorites [10]. Togetlrer with measurements of lunar soils irradiated by the solar wind [11], this has led to tlie concept of an originally l:,N-depleted protosolar nebula gas, in which variable extents of 15N-rich solid material were injected so as to yield the large range of N isotopic variations found among the solar System objects. Thus tire l5N/14N ratio of the giant planetary atmosphères is expected to vary depending on the mixing ratio of solid and gaseous components tirât contributed to tlreir formation, and a comparison betvveen Jupiter and Satum will yield information on the nature and strengtlr of the contributing cosmochemical sources.

Noble gases are also excellent tracers for the origiri and the évolution of major solar System réservoirs. Measurements of noble gas (Kr and Xe) isotopic ratios will also be of considérable importance for (1) assessing potential genetic relationship between solar System réservoirs and giant planet annospheres, and (2) investigating possible atmospheric processes tirât might hâve ifactionated atmospheric éléments.

For instance, cometary graiirs retumed by the S tard ust mission host a Ne component with air isotopic composition reminiscent of organic matter trapped in primitive meteorites rather tlran tirât of the solar nebula [12]. The measurement of the noble gas composition of tire Satum atmosphère is thereforc requested to investigate tire origin of Satum atmospheric gases.

Other key measurements include the spéciation and stable isotopic compositions of liglrt éléments (H, C, N, S) as well as trace gas speeies abuirdances (notably AsH3, PH3, GeH4, CO). Last but not least, the very précisé measurement of the hélium and néon abundances will be the key to air understanding of the formation of hélium droplets iir tire planct's interior and thus the évolution of Satunr aird Jupiter (e.g., [13]).

Kronos will also allow us to obtain a measurement of Satum's gravit y fteld to be performed with unprecedented accuracy. Even from a single flyby of the carrier spacecraft, the close polar orbit and the use of boür X and Ka bands guarantees improvements of several orders of magnitude compared to the measurements by Cassini-Huygens (Fig. 3). These accurate measurements will yield stringent constraints on the planetary interior structure. Air important source of uncertainty for models is the unknown interior rotation of tire planet. The détermination of high order harmonies of the gravity fteld will allow a détermination of whether Satum's interior rotâtes as a solid body or not [14]. Additional constraints from the magnetic fteld measurements, radiometric measurements and similar observations at Jupiter with .Tuno will help précisé the State of rotation of Satum's interior. With further constraints on the hélium and oxygen abundance measurements aird on the basis of présent models of Satum's internai structure, we expert to detennine the mass of the central core to within 10-20% accuracy, and to evaluate whether or not a layer of hélium surrounds its core.

A combitred analysis of in situ and remote Kronos measurements of Satum's composition, global water abundance, high accuracy gravity and magnetic ftelds will £) Springer Degree of Zonol Hormonics provide us with a detailed view of the planet's structure (Fig. 3). The comparison with similar measurements at Jupiter will yield a unique set of data that will permit us a global understanding of the évolution of the giant planets and their fondation within the solar System. Within the next 20 years, we expect thousands of extrasolar planets to be characterized, but with only limited information (e.g. mean density, major atmospheric gases). The parallel study of Satum and Jupiter enabled by Kronos is fundamental to understand giant planets both from a global statistical perspective, and in fine detail.

Meteorology and atmospheric dynamics

The dynamics and circulation of gaseous giant. planetary atmosphères are important physical attributes for many disciplines, not only for planetary meteorology. Heat deduced from a simple thermochemical model [4,16,17] Satum with 1x. 5x, 10x solar O, N, S, Ar suggest an interior rotation period as short as 10 hours 32 min [26]. As with Jupiter, Satum's cloud-level meteorology may extend much deeper into its convective interior than the shallow 'weather layer' within which the observed patterns of wind and clouds may be observed from space. Observational constraints on such (deep) atmospheric dynamics are very difficult to obtain, not least because of some fundamental 'degeneracies' in the near-surfàce signatures of deep and shallow circulations i.e. very different processes (deep or shallow) may produce surface signatures that are more or less indistinguishable. A major objective of Kxonos is therefore to obtain measurements (especially of Satum's detailed gravity field; see Fig. 3) that can résolve some of these ambiguities and ascertain which aspects of Satum's meteorology and circulation are deep-seated and which are shallow.

Images in the visual range by spacecraft hâve captured a plethora of meteorological phenomena immersed within the altemating pattern of zonal jets (Fig. 5), including: (1) cyclonic and anticyclonic eddies with closed circulations and sizes ranging from ~1,000 5,000 km [27,[START_REF] Garcia-Melendo | Numerical models of Satum's long-lived anticyclones[END_REF]. Particularly interesting is tLie strong cyclonic vortex found around the Southern Pôle [18], (2) Convective storms are relatively common at mid-latitudes [29,30] and probably fuelled by "moist" ammonia and water vapour latent heat release [31]. A major event are the "Great White Spots" (GWS) that occur sporadically in Satum (mainly at équatorial latitudes) attaining a size of 20,000 km before they spread zonally [32]. (3) Waves of different types hâve been detected at cloud level and in température maps. Most significant are those seen at cloud level on Satum, as for example Üie mid-latitude northem "ribbon" that moves with a speed of 145 ms-1 [33] [START_REF] Sromovsky | Voyager-2 observations of Satums northem mid-latitude cloud tèatures-morphology, motions, and évolution[END_REF], and the "hexagon" that surrounds the northem pôle at 78°Nortli [34] that appears to remain roughly stationary with respect to the planet rotation, at least as measured by the Voyager spacecraft.

This variety of meteorological phenomena aie observed at clouds and hazes vertically distributed above the ammonia condensation level. The properties of the The location of most convective storms appear marked with green dashed boxes clouds and their temporal changes hâve been studied mostrecently by [35] and [36].

Related features are also seen at other levels from observations and retrievals of température and composition at infrared wavelengths [27]. The zonal jets are seen to extend into the stratosphère though decay with height [37]. Some eddy features are also évident at higher levels though some features such as Satum's north polar hexagon seem to be confined to beneath the tropopause.

As with Jupiter, Satum's cloud-level meteorology may extend much deeper into its convective interior than tlie shallow 'weather layer' within which the observed patterns of wind and clouds may be observed from space. Seasonal effects are also much stronger on Satum than on Jupiter, and hâve a major impact on thermal structure in the stratosphère and upper troposphère. While Kronos will not monitor such effects directly, they will influence the State of tlie atmosphère at the time of its encounter with Satum, and will provide information by comparison with earlier missions (such as Cassini). The natui'c and distribution of convection in Satum's atmosphère is a major objective of tlie Kronos mission. It is now widely thought that convection, modified by the latent heat effects of water condensation and conséquent changes in molecular weight, is a critical component of tlie tropospheric circulation of both the gas giant planets [38].

Moist convection is highly intermittent and leads to massive storm cells --1000 km wide which carry water and otlier material vertically over -3 scale heights (~75 km) into tlie upper troposphère. In ways that are still not fully understood, moist convection works rather differently from terrestrial storms, due to tlie very low mean molecular weight of the atmosphère. As on Jupiter [START_REF] Dyudina | Lightnmg on Jupiter observed in the H-alpha line by the Cassini imaging science subsystem[END_REF], lightning may also accompany vigorous convection at times, Üiougli détection at Satum in association with particular convective events has been more elusive [30].

The horizontal distribution of convection is also highly variable. pressure of up to 10 bar, at two dynamically distinct locations on Satum. The profile of the water abundance will also be measured in order to constrain its rôle in moist convection processes. For the camer, üiese goals will require measurements of température, deep cloud distributions, ammonia and water abundance distributions and deep winds, using a combination of near infrared and microwave radiometry, and gravity field measurements ffom a close periapsis pass.

Magnetic dynamo, magnetosphere and radiation environment

Like Jupiter, Satum has a largely rotation-dominated planetary magnetosphere. But like Earth, solar wind convection and radial transport play major rôles [START_REF] Blanc | Magnetospheric and plasma science with Cassini-Huygens[END_REF],

Understanding Satum 's magnetosphere (Fig. 7) will provide the "missing link" between the Jovian magnetosphere (then explored in-depth by Juno) and the Earth's magnetosphere (explored by many satellites).

The Kronos mission has the potential to résolve major open issues about Satum's magnetic field, magnetic dynamo and internai rotation, its polar magnetosphere and radiation environment, and their couplings with the planet itself. Beyond While ail other magnetized planets in the Solar System hâve magnetic dipole axes tilted at an angle of 10°or more with respect to their rotation axes, Satum's global magnetic field is almost perfectly axisymetric (angle of dipole tilt <0.7°) [START_REF] Davis | A model of Satum magnetic-field based on ail available data[END_REF],

although a high-latitude 'anomaly' has been tentatively deduced ffom Satum Kilometric Radiation (SKR) measurements [START_REF] Galopeau | Evidence of Satums magnetic-field anomaly from Satumian kilometric radiation high-frequency limit[END_REF]. Classical fluid dynamo theory [START_REF] Cowlmg | The magnetic field of sunspots[END_REF] does not allow to maintain a field with perfect rotational symmetry, but new models do allow for the dipole tilt to be as small as 0.5°. Satum's dynamo is thought to be produced by liquid 'metallic' hydrogen in its outer core, combined with stnong convective motions and the rapid rotation of the planet itself (e.g. [START_REF] Raedler | Can tlie highly axisymetric magnetic field of Satum be maintained by a dynamo?[END_REF]). Accurate global field mapping is thus a key to both the composition and conductivity of Satum's interior.

The Cassini spacecraft has already provided magnetometer measurements over several tens of close passes within 5RS (average = 3.4RS) ffom Satum's centre (Rss -Satum radius = 60,330 km). In addition the spacecraft made a closest approach at 1.33 Rs at the Satum orbit insertion (SOI). Fluxgate magnetometer measurements near SOI suggested a possibly less symmetrical internai field within the radius of Satum's D ring (1.11 Rs), as compared to the field model built ffom the more-distant periapses. Analysis is on-going, however further multiple passes and £) Springer high-resolution scalar magnetometry at these very small distances-which are beyond the scope of Cassini's trajectory-are required to improve our knowledge of higher-order components or 'mullipolcs' of Satum's field.

Another unique aspect of Satum's magnetospheric field is related to the drifting rotational signatures of the azimuthal field and SKR. A mission such as Kronos, especially with an orbiter element, is expected to provide additional close-range (<1.3 Rs) magnetic data for Satum, of great value for addressing the origin and nature of its 'rotational anomaly', and characterize the internai (Enceladus ?) [START_REF] Gumett | The variable rotation period of tlie inner région of Satum's plasma disk[END_REF] or extemal (solar vvind ?) [25] control of the period variability, in order to déterminé Satum's tme internai rotation rate.

At Jupiter, the most dramatic signature of the planet-magnetosphere coupling and dynamics is the aurora (best observed in tlie UV), including the near polar signatures of the satellites-magnetosphere interaction. Ihis coupling involves electric currents flowing along field lines, energetic particle accélération, and associated plasma waves and radio waves. The latter permit remote sensing of the coupling variability as well as insights to tlie microphysics of the émission (électron distribution function tapped). Juno will significantly advance our understanding of Jupiter's polar magnetosphere via its polar orbit with low altitude perijove, allowing in situ measurements of fields, plasmas, currents, particles, at the key région where tlie magnetospheric activity is focused along converging magnetic field lines. Similar measurements should be performed at Satum (top right inset of Fig. 7) to better understand the magnetosphere-planet coupling. The origin of Satum's rings is still not understood. They were possibly fonned by fragmentation of cornets destroyed by tides after a too close passage [START_REF] Dones | A recent cometary origin for Satums rings[END_REF] or by tlie destruction of an ancient satellite [50]. Our limited knowledge of the physical properties of the ring particles prevents us irom discriminating these different scénarios. A key information would be the précisé size distribution of ring particles down below the 10 cm scale. Direct observation of ring particles would al low us to achieve a major ring science objective: tlie physics of accretion (similar to planet formation). Despite tlie strong tidal field of Satum, limited accretion is tlieoretically possible (e.g., [START_REF] Canup | Accretion in the Roche zone-coexistence of rings and ringmoons[END_REF]). Recent works show tliat moonlets and ringlets may be the two si des of a same object [START_REF] Canup | Accretion in the Roche zone-coexistence of rings and ringmoons[END_REF], Cassini lias recently provided new indirect proofs of tins [START_REF] Tîscareno | 100-Metre-diameter moonlets in Satum's A ring from observations of 'propeller' structures[END_REF]). Direct observation of these aggregates would be valuable and would unveil a new elass of Solar System objects: temporary moons. Thus, an exotic accretion £) Springer 536 physics could take place, resulting in the formation of temporary structures that need to be characterized in detail.

3 Mission profile

General mission architecture

The nominal mission includes two atmospheric probes dclivercd by a single carrier spacecraft. The primaiy' science objectives of both éléments is the formation of Satum and the origin of its atmosphère. The probes enable in situ measurements of Satum's atmosphère to be performed down to a pressure of 10 Bar. The carrier spacecraft is strongly inherited from Juno and enables a more global investigation of Satum's atmospheric H20 and NH3, gravity and internai magnetic field, rings, and the imier magnetosphere.

Launcher requirements

Since Kronos was proposed as part of an international collaboration betvveen ESA and NASA, the launch would hâve been possible with either a NASA or a European launch vehicle. Since the total mass of a flyby-based architecture with the two probes is estimated at less than 3000 kg, the projectcd launch vehicle for the mission could be an Atlas V-551. 2 and3.

Springer Preliminary analysis indicates that a data rate of ~60 bps from each probe is achievable without major new technology. This translates to a total data volume of Data link types on the relay architecture include: X-band telecom link for tlie carrier to the DSN; and 400 MHz UHF between the probes and relay spacecraft.

Assuming a 400 K Satum hot-body température noise at 400 MHz UHF frequency, and a corresponding 1.5 dB atmospheric atténuation at 10 Bar, the probes could support a data rate between of 512 bps, giving a total data volume of ~2 Mb from each probe.

On the carrier, the X-Band System uses ~35 W power and a 3 meter high gain antenna (HGA). Medium gain antenna and low gain antennas are likely necessary for maintaining links during cruise and safing scénarios. On the probes, the UHF

LGA will work with an Electra transmitter at 20 W RF output. This requires an upgrade to tlie Elecfra transmitter, which currently opérâtes at 12 W. Both ends of the radio link between the probes and the relay spacecraft must contain Ultra Stable Oscillators (USO) needed for radio interferometric (VLB!) and Doppler measurements (Doppler Wind Experiment -DWE-and Planetaiy Radio Interferometer and Doppler Experiment -PRIDE).

Ring science architecture

For the ring science purposes, we suggest an additional dedicated probe: Lora (Landers on Rings Array). Lora would not be a large probe, but rather a collection ( 2to 3) of very small and simple probes, like "MicroProbes", containing only one efficient caméra. The fleet of 2 to 3 identical probes would be dropped by Kronos during the flyby, and the probes would go directly into the rings at different locations, allowing to obtain high resolution images. One of the probes should target tlie ring plane, through a gap near a dense ring (B ring), in order to get information on tlie vertical structure of tlie ring.

Carrier

The spacecraft design is directly inherited from Juno replacing Juno's radiation vault mass with the probes. Juno is designed as an outer planet orbiter and already includes many of the components necessary to implement the mission. Low Intensity Low Température (LILT) solar panels from Juno and Bepi Colombo missions are baselined for Kronos, but require further tests for Satum environment.

The carrier is designed to deliver the probes to Satum and to perform remote sensing measurement during its flyby. An orbiter configuration would include additional propellant tanks, and a modified power System to account for the higher power demands. Studies performed at both CNES and NASA [START_REF] Balint | In 5th International Planetary Probe Workshop[END_REF][START_REF] Langevin | Inst, for Space Astrophysics in Orsay[END_REF] demonstrated tliat shortlived missions could operate at Satum using photovoltaic power génération. A flyby mission is expected to use batteries on both the probes and on the carrier spacecraft, and employ solar-panels for backup only, in case not ail data is transferred back to Eartli in a single 8 hours telecom pass. For an orbiter mission the power requirement is expected to be continuous: near the periapsis the spacecraft vvould perform its science measurements, while during craise to and from apoapsis it would perfonn data transfer to Earfh, liousekeeping, and recliarge its batteries. This would hâve a significant impact of power System sizing. While solar power genenttion is expected to be the baseline configuration for the flyby architecture, future mission studies should also address the use of RPSs on an orbiter, in order to assess its impact on mission cost and architecture. is not feasible for the probes to retum data from the 50 to 100 Bar pressure levelswhich would be required to measure the water abundance in situ-because of significant microwave absorption. Consequently, the probes will only be required to measure atmospheric composition and dynamics to about 10 Bar, while water and ammonia abundance will be measured using passive microwave radiometry from the carrier (like Juno) to 50-100 Bai-.

The Satum probes will use significant héritage from the Galileo probe (Fig. 8), including the thermal protection system (TPS), aeroshell design, and subsystems.

Except for the TPS, héritage from tire Huygens Titan probe will also prove AJAA'83, 21st Aerospace Science Meeting, Jan. 10 13, 1983, Reno, NV) The probes are considered autonomous objects, consisting of a set of science instruments, and snbsystems. Each probe is housed inside a protective aeroshell. Its total mass of~330 kg consists of a 222 kg décélération module, and 117 kg descent module. The descent module houses the science instruments and subsystems, while the décélération module includes the aeroshell, TPS, parachutes and séparation hardware. héritage ffom previous probe missions such as Galileo [START_REF] Niemann | The Galileo probe mass spectrometer: Composition of Jupiter's atmosphère[END_REF] (Fig. 9) and Cassini/ Huygens [START_REF] Waite | Ion neutral mass spectrometer results from the first flyby of Titan[END_REF].

Composition and isotopes

Light éléments to be measured with an atmospheric probe mass spectrometer include H, C, N, S, P, Ge, As, and the noble gases He, Ne, Ar, Kr, and Xe. These measurements together with a détermination of D/H, 13C/12C, 15N/14N, 3He/4He, 2ûNe/22Ne, 3gAr/36Ar, 36Ar/'4ûAr, and the isotopic composition of Kr and Xe will provide a complété comparison with the Galileo Probe measurements at Jupiter and reveal key différences or similarities between these giant planets. However, we know that such a probe will not altow us to measure the deep, well-mixed abundance of oxygen, because of water condensation and meteorological processes. As the £) Springer below enables these précisions. The Jovian elemental ratios of C/H, N/H, S/H, and D/H in this table were measured in the molecular species CH4, NH3, H2S, and H2

respectively. The GPMS made continuons direct measurements of gases ingested into the vacuum of the instrument in the 2-150 Dalton mass range during the descent except for short periods of time when processed gases were introduced into the mass spectrometer.

GAS spécifications

The proposed Satum Probe Gas Analysis System lias a high héritage in both successfiil atmospheric probe composition/isotope investigations conducted at

Jupiter [4,[START_REF] Niemann | The Galileo probe mass spectrometer: Composition of Jupiter's atmosphère[END_REF][START_REF] Mahafty | Noble gas abundance and isotope ratios in the atmosphère of Jupiter from the Galileo probe mass spectrometer[END_REF] and Satum's moon Titan [START_REF] Niemann | The abundances of constituents of Titan's atmosphère from the GCMS instrument on the Huygens probe[END_REF], Eléments of the GAS designed to improve the précision of the measurements are adapted front in situ calibration, static mass spectrometry, gas séparation, Chemical processing tecltniques and advanced pumping Systems developed for the Rosetta mass spectrometer Ptolemy [START_REF] Wright | PTOLEMY-an instrument to measure stable isotopic ratios of key volatiles on a cometary nucléus[END_REF] and presently under development for the 2009 Mars Science Laboratory [START_REF] Mahafiÿ | Exploration of the habitability of Mars: Development of analytical protocols for measurement of organic carbon on the 2009 Mars Science Laboratory[END_REF].

The mass spectrometer is a quadrupole mass spectrometer similar to more than a dozen tliat hâve been developed at the NASA Goddard Space Flight Center and successfully operated in the upper atmosphère of the Earth and Venus.

Noble gas enrichment and static mass spectrometry'

The combination of a noble gas enrichment System, a gas scrubber that can remove ail chemically active gases, a Chemical getter pump in the MS analyzer région, and a high conductance valve between the MS and the WRP will enable static MS operations for higher précision noble gas measurements than were possible at Jupiter with the GPMS. The static measurements intenupt direct sampling of the atmosphère for a portion of the descent, bat are necessary to obtain tire desired précision. The noble gas enrichment System was proven on the Galileo Probe MS and consists of Chemical getters and traps tirât can enricli the trace noble gases Kr and Xe. The Galileo Probe noble gas enricliment System was able to provide measurements to the sub ppb levels for isotopes of Jovian Xe.

Nitrogen combustion for !>N/I4N isotope measurement

The l5N/14N ration on Jupiter was obtained from the NH3++ signal at 8.5 and 9

Dalton for the two nitrogen isotopes. Due to the difficulty of removal of spectral interférences in this approach we provide an NH3 oxidation reactor on GAS to produce N2 so that this measurement can be carried out with higher précision. This element of GAS could be developed using techniques similar to those developed for Rosetta and Beagle 2, and expertise and techniques developed for the analysis of solar wind nitrogen implanted in the GENES1S targets.

Meteorology and atmospheric dynamics

In situ measurements that address the meteorology and atmospheric dynamics goals of the Kronos mission include density, pressure and température sounding to Springer characterize the atmospheric structure tutti instruments to measure the radiative balance, cloud structure, and zonal winds also down to a pressure of 10 Bar.

Measurements achieved vvith a Saturn probe nephelometer

The composition and précisé location of cloud layers in Satum are largely unknown.

They may be composed of ammonia, ammonium hydrosulfide, or simply water.

Because of this relative paucity of information on Satum's clouds, the demands vve place on a cloud particle sensor (nephelometer) are significant.

For Kronos, we hâve baselined a new instrument which would measure not only the amplitude phase fonction of the light scattered by the clouds iront a laser source on the probe, but also Ütc polarization ratio phase fonction as well. It does this at two wavelengths, separated by about an octave in wavelength near 1 pm. The phase fonctions are sampled at six discrète angles, chosen to maximize their leverage in distinguisliing betvveen different size, shape and aérosols compositions. These measurements at an optical depth unity as long as ~10 km permit to détermine the aérosol number density, their particle size, an indication of the particle shape, and the paiticle's index of réfraction at the two wavelength, which may be used to infer the molecular constituency of the aérosols.

Atmospheric structure instrument

The Kronos Atmospheric Structure Instrument (AS1) consists of three primary sensor packages: (1) a three axial accelcrometer (ASI-ACC), ( 2) a pressure profile instrument (ASI-PPI), ( 3) température sensors (ASI-TEM). Tire proposed instrument will benefit from a strong héritage of the Huygens AS! experiment of the Cassini/ Huygens mission [START_REF] Fulchignoni | The characterisation of Titan's atmospheric physical properties by the Huygens atmospheric structure instrument (HASI)[END_REF]. The key in situ measurements will be atmospheric density, The design of tlie microwave radiometers will follow closely tlie Juno design, with tlie exceptions that tlie measurement at Satum is considerably less driving.

Without the strong synchrotron émission from Jupiter, the measurement noise will be less andthus tlie instrument design less complicated (beam pattern, etc.). Fuithermore, the sélection of wavelengths (ffequencies) will be tuned to the Saturn atmosphère. As with Juno, it is expected that the radiometers will be a set (approximately 6) of frequencies ranging from 1 cm to 100 cm. Having tlie full latitude coverage is essential for understanding the rôles of ammonia and water in Jovian meteorology and for placing the probe measurements in context. For a Cassini-like configuration, eacli sensor would be capable of measuring fields up to -44.000 nT, well below the atmospheric depths at which the Kronos probes would be expected to operate. Chrétien reflector, with lm focal length and diffraction limited at 5 * 10"6 rad (> a 15 cm diameter) would be suited. As the probes get doser to the ring, the resolution increases and allows to achieve the ring science objectives.

A low inclination impact angle is preferred (-1°), its feasibility depending on the arrivai trajectory of the Kronos spacecraft. A high resolution imager for ring science that gives the required <1 10 m résolution images requires narrow angle caméras on the Lora probes that will be evaluated in the Kronos Phase A study. An example of synthetic images of A ring wakes at different résolution is shown in Fig. 10.

Assuming a resolution of 6*10~b radians/pixel (like the Cassini ISS-NAC), the minimal distance of approach to achieve a given resolution is given in Table 5.

The maximum scientifîc retum is expected when the distance to the rings is <100 km. However, once the Lora probes are below 8,000 km, they could achieve a better resolution than the best Cassini images.

In order to stabilize and orient the spacecraft to allow a good pointing, a set of reactions wheels and gyros onboard are necessary. The small inertia of the probes may allow easy re-orientation during the descent.

Timeline for the Lora probe

In the current State of the design of the Kronos mission, the exact trajectory of the Lora probe has not been precisely determined. We assume below that the Lora Probe would be launched on an hyperbolic trajectory, with a ~1°inclination above the ring resulting timcline for the Lora probe is presentcd in Table 6.

During ils descent to the ring plane, the Lora probe would transmit the eollected data as soon as they are eollected (in the spirit of what was done for the Huygens descent on Titan), because of the possible destruction of the probe during its Crossing of the ring plane. If the Lora probe survives the ring-plane Crossing, then new data could be eollected and transmitted again to tire Kronos spaeeeraft.

Since the Lora probe is meant to be a "microprobe" (i.e. with a single instrument and a simple communication System) its communication package would be designer! primarily for communicating with the Kronos spaeeeraft only, which would record and transmit üie data to Earth. Kronos also provides a possibility to study Satum's rings in fine details. Altogether, tlirough innovative approaches this international mission allows a combination of essential measurements that concem tlie ringed planet but impact many disciplines and the understanding of our origins. Introduction

The Cassini-Huygens mission has been investigaling the Saturnian System since J'uly 2004 when it performed the Satum Orbit Insertion (SOI). Titan, the largest satellite of Satum, is one of the major targets of the mission mainly due to its dense atmosphère. Indeed, the Cassini-Huygens mission has revealed the complex organic nitrogen-dominated atmosphère of Titan, rich in methane (1.48% [5]) and the diversity of its multivariable surface [e.g. 6],

The température on Titan's surface has been measured in situ at almost 94 K [7], which is close to the triple point of methane, meaning that methane on the satellite can exist as liquid, gas and ice [8], Literally, Titan hosts a complex and active organic global chemistry where methane plays a critical rôle in a cycle resembling the terrestrial hydrological one [9].

The presence of large hydrocarbon liquid deposits on Titan's surface was assumed before the Cassini-Huygens mission, considering the thermodynamics which dominate the satellite globally [e.g. 10,11]. Although the existence of a planetary-scale océan has been rejected, the Cassini orbiter has unveiled concentrations of large organic pools close to both pôles [1,12,13]. However, the Cassini RADAR beams are not able to penetrate through the liquid surfaces hindering sounding their interior. Therefore such features can be described only by modeling [e.g. 14]. However, it is crucial to understand how these hydrocarbon liquids contribute in the Mathieu Hirtzig, Athéna Coustenis LESIA Observatoire de Paris-Meudon Meudon, France active methane cycle. According to the contemporary photo dissociation rate of methane, it should be vanished within a period of 10-100 Myrs [9], since currently no methane source has been identified to replenish it. Towards this direction, as well as the necessity of the thorough investigation of these features, a lake touchdown by a probe has been proposed, carried by a future mission [3],

The usage of Micro-Electro-Mcchanical-Systems (MEMS) devices as infrared emitters has been recently proposed [15], as part of the science surface properties package of a future probe to Titan like the TSSM Lake Lânder or the Titan Lake Explorer [3,4], In this paper we suggest MEMS devices to operate as micro-laboratories by including also radio frequency RF wavelength emitters and température and pressure sensors.

Thus, these micro-machines could obtain the 3D sounding of the liquid deposit, its Chemical composition and detect the presence of any biological building blocks within the liquid. Likewise, the température and pressure micro-sensors could provide the vertical pressure, température and density profile of the liquid deposit

The MEMS pattern, owing to their very small shape and size without reducing their operational performance, seems idéal payload for a lake lander probe on Titan as well as outer planetary space missions in general. MEMS implémentation in Titan's exotic environment is a great challenge for science, engineering and space physics.

In section II we présent the advantages of MEMS devices, while in section III we report MEMS already used in space applications. The experiment concept and the MEMS technology embedded is described in Section IV and Section V contains the concluding remarks.

II.

Micro-Electro-Mechanical-Systems (MEMS)

Devices

Any small-size product within the range of a micron to a ccntimeter, which also combines mechanical and electrical structures, can be identified as a MEMS device. Although the first device was constructed by H. C. Nathanson in late 1960s [16], the industrial techniques were unable to produce such minor mtegrated circuits and only after the 1980s did start the massive production in microscopie scale [17],

Due to their small size (20 jim to 1 mm) and shape, MEMS présent properties that increase the performance of every scientific experiment for a wide range of uses and applications.

In general, they provide an excellent rate of the optimum shape to the rendering performance relation. Micro-devices can now be included in evety part of any experiment and also visit impossibly approachable places by traditional instrumentation.

They can successfblly replace quite larger devices giving the user the oppoitunity to explore the micro-structure of the nature. Therefore, due to the minimization of the cost both in their manufacture and operation, MEMS instruments aie préférable for scientific use by improving the quality and the amount of the obtained data range.

III. MEMS Devices In Space Sceme

The benefit of implementing MEMS techniques in Aerospace and Space Systems is obvious especially due to the reduced requirements in size. mass, povver issues of many Aerospace/Space applications [18]. In order to overcome portability and consumptiori limitations in any future outer planetary missions, new alteiTiative technologies hâve to be implemented. The minimization of the embedded electronic instrumentation is a promising approach when it employs MEMS Systems. Therefore. the presence of MEMS in space Systems otïers important monetary advantages to space agencies by severely shi inking tire cost of any mission, while they are reducing the weight of future spacecrafts.

Among many other applications, MEMS devices hâve already been used on the adaptive optics iinaging tcclinology in major télescopes [19][20][21][22], Moreover. JPL has already'

developed a new MEMS déformable iniiror (DM) sy'stem for NASA's adaptive space-based télescopes and in particular for Terrestrial Planet Finding (TPF) mission [START_REF] Gumett | Radio and plasma wave observations at Satum from Cassini's approach and first orbit[END_REF]. Likewise, several proposais hâve recently been submitted for futui e space missions and télescopes, exploiting the new leading edge of MEMS technology. Especially, MEMS techniques hâve been incorporated as detectors, a spectrophotometer and an IR caméra of the new proposed telescope called MTEL (MEMS Telescope for Extrême Lightning) in order to observe the extreme lightning occurring in tiie upper atmosphère [24].

MEMS devices seem also appropriate for obseiving fast moving objects and transient events by the proposed spacebased telescope called Obscura [25], Moreover, MEMS hâve also been implemented in Space Infrared telescope for Cosmology and Astrophysics (SPICA), the Japanese coronagraph which is planned for a launch in 2017 [26].

Without doubt, although the use of MEMS in space applications is in fairly early stage, they will certainly optimize the scientific potential of any future mission.

IV. Experiment Description MEMS as components of the science surface properties package of a future Lake Lânder probe to Titan aie being discussed hereafter. We propose MEMS machines to operate as (a) infrared emitters inside tiie liquid and (b) micro-laboratories by including also radio frequency' RF wavelength emitters as well as température and pressure sensors.

During the last stages of the descent of a probe in Titan's atmosphère and a few meters before landing on a liquid surface, the probe will releasethe MEMS devices (Fig. 1). This procedure will be continued aller the landing, when the probe will release more MEMS directly into the liquid. From this point, they commence their operation by transmitting instantly as IR and RF emitters. Some of these devices will float on the surface of the liquid and other will dive without stopping transmitting depending on their weight and exterior design. Schematically, the proposed experiment will be consistent of two phases:

A.

Phase One: The Deployment Few meters before the touchdown of the vehicle, MEMS capsules will be released into the atmosphère and begin simultaneously to operate as micro-laboratories. Another group of MEMS can be also released aller the landing of the lake vehicle inside the liquid environment (Fig. 1).

Figure 1. The deployment of the MEMS netwerk during the entry descent and landing (EDL) procedure of the Lake Lânder.

B.

Phase Two: The Operation Subsequently to Phase One, these devices will flow or dive into the liquid hydrocarbons and continuously transmit data to the Lânder depending on the duration of their power supplies (Fig. 2).

Phase B: Operation This deployment will provide a broad network of surface and submarine IR and RF sources, while the signal receiver sensor will be aboard the Lake Lânder. The MEMS micro-probes will follow the swarm pattern [27] in which each apparatus single operation is a component of a large group of similar devices. Apparently, their combination will enhance the results.

C. MEMS components

As mentioncd before, each MEMS device will include a combination of scientific instruments in micro scale. In particular, one pair of infrared (IR) and radio frequency (RF) emitters and two sensors for measuring the température and the pressure will be enclosed in each micro-shell. Additionally, MEMS technology can be part of the IR spectrometer's equipment aboard the Lakc Lânder.

1) Infrared Sources

Apart ffom traditional IR sources, micro-machined infrared emitters hâve been constructed, comprised of a photonic crystal modified micro-hotplate that emits thermally stimulated infrared radiation in a narrow band. As far as the perfonnance of the device is concerner!, it exhibits efficiencies in excess of 10 % greaterthan compétitive technologies. These applications are like the deposited filament emitter, with the avantage of minimization of the filament which reduces the thermal mass of the System and enhances the modulated performance [START_REF] Garcia-Melendo | Numerical models of Satum's long-lived anticyclones[END_REF],

2) Micro-mirror MEMS technology can be implemented at the IR spectrometer's architecture of a classical Michelson interferometer, being a transducer of the IR sensor. MEMS translational mirrors, covering an area of l.lxl.5 mm2 hâve been already tested giving reliable measurements at high scanning speed -in milliseconds-which enhances signal-tonoise ratio, suspended on a two long bending springs pattern [29]. A recent approach is the pantograph, which even though it exhibits less stability, it covers a surface mirror of 7 mm2 providing a spectral resolution of 20 cm'1 [30]. The IR source can be based on Silicon structure, a pattern which is already in long-term use for detecting gas in oil platfonns without a failure recording [31], In Titan's case, the target gases are mainly hydrocarbons, nitriles and carbon dioxide.

The surface température in Titan's environment is approximately 94 K, as it has been determined by the Huygens probe [7]. MEMS devices, released by the Lake Lânder at low height and inside the liquid, hâve to be subjective to such low températures. In fact, micro-devices are already adaptable in cryogénie conditions. In the SPICA coronagraph [26] the included micro-scale machines are designed to operate in cryogénie conditions, while the whole telescope will be cooled down to 4.5 K [32], This cryogénie environment can cause serious damage to the mirrors such as deformation and flaking.

To achieve robust function in this harsh température conditions these micro-mirTors can be based on a silicon-on-insulator wafer, while their back surface can be made of polycrystalline Silicon [33],

3) RF transmitter

The exact composition of the lake as well as its physical parameters are still unknown. If the liquid malerial behaves as a partial conductor (when consisting of i.e. ionic contaminants), it will preventRF propagation in long distances.

Such liquid conductors set the independent operation of these RF emitters inside the liquid under discussion, since their high permittivity and permeability cause strong atténuation of the signal power. For this reason the capsules which will be designed to dive in the lake can be connected by a wire with the mother ship as it shown in Figs 1,2. This pattern can also support the micro probes with energy originated by the carrier.

An overview of RF MEMS transmit/receive switches can be found in [34,35] 4) Température and pressure sensors

The température and pressure sensors hâve the functional différence that they should interact with the environment surrounding the MEMS capsule.

Any température measurement is correlated on the variance of an attribute of the exposed material with any ambient température change. MEMS resonator-based oscillators hâve been proposed to be used as Complementary metal-oxidesemiconductor (CMOS) température sensor with a resolution of 0.008°C [36],

The MEMS pressure sensing technique is based on the Silicon piezoresistive effect, where any pressure change will cause deflection and internai stain change, which will resuit output voltage variation [e.g. 37,38], Usually, a Wheatstone bridge circuit delivers such voltage measurements, in order to correlate them with the applied pressure.

V. CONCLUSION

The MEMS experiment will vertically measure the température and the pressure of the liquid deposit of an hydrocarbon lake on Titan's surface during a future mission, while at the same time it will operate as an internai source of infrared wavelengths giving the opportunity to the Lake Lânder to analyze the composition of the lake by receiving its IR signais. Moreover, the recorded reflections of the RF signal emitted by MEMS will construct the topographical map of the bottom of the lake.

Although it is intentional to hâve many devices in order to construct a network, the major issue of portability cornes up.

How'ever, micro-electro-mechanical devices suit perfectly in this kind of requirements because of their small shape, size and weight.

Obviously, the utilization of any innovative instrumentation requires complété knowledge of the environmentai issues and parameters and of course the deep understanding of their complexity.

In the contemporary market, MEMS can be easily found combining low cost and high reliability. These devices feature low power consumption, high-modulation depth, high emissivity and a long lifetime.

In a sense, this ambitious experiment can easily accomplis!)

both the 3D sounding of the lake as well as its Chemical composition.

MEMS devices combine:

1) Very small size (20 pm to 1 mm)

2) Fast pulsing (thanks to low mass of the em itter)

3) Limitation of low output power (450m\V) We also suggest possible target areas with internai dynamic potentiai and multivarlable surface expressions. Ilence, it will bc possible to idciitlty active régions on tlic satellites, which will provide important information regarding the fluid transfer Introduction Since 2004, the Cassini-Huygens mission has made exciting discoveries in the Satumian System and especially with regard to its satellites Titan, the largest one, as well as the enigmatic Enceladus. Not only hâve their observations furthered our understanding of the complexity of these harsh environments, but many questions about their surface and interior hâve also been raised e.g. [1,2].

Titan, the second largest moon of the Solar System may be représentative of many planetary bodies and consequently, lire understanding of its internai geology may allow scientists to enlighten the inner structure of an entire class of planets and moons [3], Similarly, Enceladus, although it has minor dimensions (radius of 252 km), it is extremely interesting due its huge geysers which mainly feed the E-ring of Satum. These plumes show that the moon is geologically active hosting a possible internai water océan, the source of its vents [4,5].

Likewise, the Galilean satellites of Jupiter présent many similarities with rocky planetary bodies in tenus of their surface features and atmospheric environments. Europa, the smallest of the Galilean satellites, hosts a relatively stable Mathieu Hirtzig, Athéna Coustenis LE SIA Observatoire de Paris-Meudon Meudon, France environment, a putative internai océan in combination with a young evolving surface, and therefore the moonhouses a great astrobiological potentiai [6]. On the other hand Io seems to be the most geologically active planetary moon through the solar System considering its huge volcanic activity occurring at the time of observation [7],

Without doubt, the icy satellites orbiting outer planets can significantly contribute to the study of geological processes across the solar System. Comparative Planetology can enlighten fundamental mysteries conceming the origin and the évolution of Earth and the solar system in general and predict future geological events. However, the internai geological structure of each planet or moon can only be detemuned by in situ measurements.

The local or global tectonic field, meteoroid impacts and moon's tidal deformations induced by .Tupiter/Satumian's gravity field as well as température and pressure fluctuations may cause ground vibrations within the icy moons. Such ground vibrations provide information about the nature of the subsurface material, its fracture and its Chemical composition.

Therefore, the Cassini-Huygens mission héritage encourages scientists to discuss about launching future missions towards the outer planets' satellites. For that reason, two large L-Class missions set these moons as primary targets: Europa Jupiter System Mission (EJSM) [8] will investigate the subsurface, surfaces and atmosphères of Jupiter's satellites, while the Titan Satum System Mission (TSSM) [9] will advance our current knowledge regarding the Satumian icy moons. For further information the reader is referred to visit the NASA's website of the Outer Planet Flagship Mission (http://opfm.jpl.nasa.gov). This paper examines the possibility of installing seismic experiments on icy satellites of the outer planets, setting the minimum technical requirements and describing scientific achievements and possible problems. The sections II and III contain a brief description of such space seismic equipment and the micro Gas Chromatograph (pGC) for gas sampling respectively, while in section IV the risks of such an instrument are discussed. Finally, the major goals of the proposed experiment are depicted in section V.

Data obtained by a seismic experiment, in combination with the analysis results of the pGC will provide us with details about the interiors of the moons and tlieir connection with the observcd surface featurcs.

II.

Description Of The Seismic Instrument

The seismograpli is the basic instrument for measuring any Botli low and high frequencies can be recorded by broadband seismometers on Earth. Newer seismographs measure ground movements smaller than 1 nm. There are several kinds of these instruments depending from the surface's location.

By using seismic instruments we extend our knowledge of planets' interior at the Solar System. Additionally, both these space and telluric observations vvill help us understand better the origin and the evolutionaiy path of our own planet.

.-4. Instrument 's spécifications -technical requirements

The Table I below indieates the main seismic instruments requirements considering the spécifications of the shortest modem seismic instruments which operate in extreme conditions on Earth like the deep océan floor [10]. Similar sensors can be easily found in the global market. Future mission's instruments on the icy satellites' surface of the outer planets will operate in abnormal physical conditions during the entry descent and after landing of their carrier. To achicve the seismograph's accuracy in such a location the envirotimental factors (mainly température and pressure) should be considered.

B. System Hardware Architecture

The ground vibration is the input signal to the System. If the ground vibrâtes within a range of 0.001 Hz to 100 Hz, it will be recorded by each one of the three sensors. Tlie sensor System contains a transducer, the device that couverts the mechanical motion into electrical signal. A piezoelectric accelerometer can be used as a transducer to sense any vveak or strong ground motion in a low frequcncy range (up to 100 Hz) without using extra power to operate [11],

Once the sensor records a motion, a signal vvill be sent al the Main Processor Unit (MPU) of the System. The MPU is responsible for the operation, the administration and the maintenance of the instrument and contains the core of the application software. After receiving and recording the signal from the MPU, it will be transferred at the next component of the instrument, the Multiplexer (MUX) by using a line interface circuit, which provides the connectivity between MPU and MUX.

At the MUX, data files will be comprcssed, converted to the right format and prepared for transmission. Because of the continuous function of the seismograpli, a service for bulky data needed. Tlie transmission frame protocols will be defined similar to the protocols ofthe lander's instrumentation. The purpose of the micro Gas Chromatograph (uGC) is to provide accuratc information about the composition of the local gaseous environment. A successfiil relative experiment with the Gas Chromatograph Mass Spectrometer (GC-MS), was performed by the Huygens probe during its descent phase [12]. This instrument on Huygens only made measurements during the descent from 170 km to the surface and it provided us with valuable data in harsh environmental conditions. We propose the micro-GC as a payload of the lander to operate during the descent phase of the probe like in the Huygens case as well as after the touchdown. During terrestrial earthquakes close to volcanoes, amounts of gas are released, triggered by the ground vibration. In some cases this relief is used for earthquake prédiction. Such instrumentation can be extremely useful 011 tlie surface of icy moons' locations where gas relief lias been identified, like the tiger stripes of the Enceladus' South Pôle.

Tlie iiGC tube/coluirm mounted on the probe hosting tlie seismograph will be triggered when any vibration will be recorded by the seismic sensors. A présentation of a low power uGC colunni can be found in [13]. The instrument will also take samples of the local atmospheric envelope in several spécifie times which will enable us to détermine any temporal variations in Chemical composition due to any recorded vibration as well as during in the descent phase of the lander. Such measurements can be extremely useful in the case of Enceladus if the probe landcd close to the tiger stripes, the great linear surface fractures from where its plumes spread out In fig. 2 the procedure of the operation of the pGC is depicted. Briefly, the instrument will collect the sample from the local gas and collected material will be droved into the pre concentrator device where the concentration will be incrcased.

Then, the sample will be heated in order to release the gas and it will be passed through the tube. This is the séparation phase, critical to the analysis procedure where each component of the sampled gas needs its spécifie tinie to traverse the tube. Eventually, the identification stage follows this séparation. The design of the signal condition circuit which converts the content of the separated gas to electric signal can follow the standards of [14]. The pGC will perform batch sampling like the Huygens/GC-MS did during Huygens Descent phase [15]. Hydrogen will be selected as the carrier gas like in the Huygens/GC-MS case [12].

By correlating the seismic soundings with the gas composition, we can infer about the significance of the internai gases and the mechanism beyond their relief. These measurements can help us solve the puzzle of icy moons inleriors. Since we lack any knowledge of seismic events on icy moons, data from a seismic instrument and global radar mapping for long periods of time will be of extreme importance. Obviously, due to power limitations, any instrument on these surfaces will operate continuously from the touchdown until its battery discharging. Therefore, a more durable energy source is necessary.

IV. Problems And Risks

A radioisotope thermoelectric generator (RTG) is to be considered An appropriate location for the placement of a seismic equipment will ensure that its recordings will represent accurately and separately every ground vibration. For this purpose, on Earth the seismic instruments should be placed in a hole of approximately 0.5 m depth. Thus, less ground noise will be recorded. Noise on an icy surface can be originated by the atmosphère due to its seasonal and diumal effects as described in [16] for the Titan case. Noise can be produced also from the local wind and meteor impacts.

In the case of Titan, its dense atmosphère protects the surface from meteoroid impacts -few craters hâve been observed [17] -therefore, the seismometers will record merely interior events. On the other hand the same instrumentation on Enceladus, Io and Europa will also measure vibrations caused by extemal sources.

The proper contact between the feet of the seismograph and the local surface at the landing area stands for another issue to be conffonted. Any particles and dust can easily perturb the seismic sensors during their function and cause random errors at the sensor's record. This type of surface features has already found on Titan. The Huygens probe landed on a relatively soft solid surface whose properties consisting of analogous to wet clay, lightly packed snow and wet or dry sand [18]. Moreover, the Descent Imager and Spectral Radiometer's (DISR) surface images showed rounded stones approximately 15 cm in diameter to lie on top of a finer-grained surface in variable spatial distribution [19]. If the icy pebbles lying over the instruments feet move or/and melt, the equilibrium position of the seismic equipment will disturbed and the sensors will lose their orientation. The measurements should be corrected if such a micro-movement of the probe is noticed.

Expérience gained from the Apollo Passive Seismic Experiment (PSE) on the Moon will be extremely useful for such an ambitious effort. PSE was the first extraterrestrial network of seismic instruments and it consisted of four seismometers deployed by the astronauts on the lunar surface between 1969 and 1972 while Earth-based stations received data for eight years. The Apollo seismometers recorded 12,558 events and these seismic data hâve been recently reanalyzed [20]. A seismic experiment has also been perfonned on Mars as part of the Viking mission payload [21]. The Viking seismograph performed a long-term operation and measured ground vibrations without any significant seismic signais. The primary source of the recorded noise was the local wind. The dimensions of the equipment were 12x15x12 cm, weighted 2.2 kg and needed 3.5 W to operatc.

V.

Major Scientific Goals-Discussion

Icy moons of the outer planets seem to be or hâve been active like our own planet, as far as their atmospheric circulation, surface geological processes and tectonics are concemed. Each moon has its own atmospheric and geological record and thercfore it should be treated separately. The internai structure of these satellites may be involved in some surface processes.

Hence, since only the seismic instrumentation can map the subsurface layers, detennine their composition and structure and measure their thickness, we will consider mounting it in a future mission as part of a landing probe payload.

Seismic waves from distant events travel deeper into the interior of a planetary body than waves from nearby events. Hence, measuring events at various distances by seismometers can provide the variance with depth of seismic velocities within icy moons.

Future missions to the outer planets icy moons will be a great opportunity for comparative planetology providing proofs of active planetary Systems. A seismic experiment can identify the existence of liquid internai deposits, with a great astrobiological potential. Such isolated environments that consist of water and organics, components that hâve already been identified on rnost icy satellites, provide idéal conditions for the survival of biological building blocks.

Figure 1 . 1 -

 11 Figure 1.1-The Pioneer 11 structure (crédit: NASA).

1. 2 )Figure 1 . 2 -

 212 Figure 1.2-The Voyager sketch with its main instumentation (source: http://pds-rings.seti.org/voyager/spacecraft/sc_cartoon.gif).

  identical twin spacecraft, the sophisticated instrumentation of which is briefly listed below (Space Science Reviews, Spécial Voyager instrumentation issue, vol. 21(2): 75-232, 1977): a) the Imaging Science System (ISS), the spacecraft's two-camera System b) the Radio Science System (RSS), the télécommunications System of Voyagers used to investigate physical data of planetary objects c) the Infrared Interferometer Spectrometer (IRIS), dedicated to atmospheric composition exploration d) the Ultraviolet Spectrometer (UVS), used to investigate atmospheric properties e) the Triaxial Fluxgate Magnetometer (MAG), which measured the magnetic fields, the solar wind coupling of the planetary magnetospheres f) the Plasma Spectrometer (PLS), which measured plasma ions and électrons from 5eV to 6 keV h) the Low Energy Charged Particle instrument (LECP), which investigated energetic particles, their differential energy fluxes and angular distributions i) the Cosmic Ray System (CRS), which studied the interstellar cosmic rays in the interplanetary and trapped in the planetary energetic particle environment g) the Planetary Radio Astronomy Investigation (PRA), which studied the radio émissions from giant planets k) the Photopolarimeter System (PPS), observed the surfaces of giants and Satum rings l) the Plasma Wave System (PWS), which measured electron-density profiles and electromagnetic waves.

Fig. 1 . 3 :

 13 Fig. 1.3: The latitudinal distribution of trace gaseous mixing ratios. The panel (a) shows the hydrocarbons and C02, while the panel (b) shows the nitriles. The error bars depict the 3a uncertainities for HC3N and C2N2 at the probed latitudes below 50°N[START_REF] Coustenis | Titan's atmosphère from Voyager infrared observations. IV: Latitudinal variations of température and composition[END_REF].

Figure 1 . 4 :

 14 Figure 1.4: The Cassini spacecraft. The main parts of the instrumentation is shown as well as the Huygens probe

Figure 1 . 5 :

 15 Figure 1.5: Cassini's Fields and Particles pallet where the instruments for in situ investigations are located

  moons. It perforais both in situ and remote measurements. It senses remotely the magnetospheric ion plasma with energies greater than 7 keV through the détection of energetic neutrals produced through charge-exchange interactions of energetic ions with cold neutrals. It also measures in 3-dimensions the ion composition and charge States for ion energies between 3 keV/e and 220 keV/e. MIMI consists of 3 sensors the Ion and Neutral Caméra (INCA), the Charge-Energy-Mass-Spectrometer (CHEMS) and the Low Energy Measuremens System (LEMMS) which detects energetic électrons (>15 keV) and energetic ions with cold neutrals. During each Titan flyby,

1. 6 )Figure 1 . 6 :

 616 Figure 1.6: The Remote Sensing Palette on board Cassini (crédits: NASA).

  Figure 1.7: Huygens probe instrumentation, the upper and the bottom side of the vehicle

Figure 1 . 8 :

 18 Figure 1.8: The space joumey of the Cassini-Huygens mission. The spacecraft took advantage of the gravitational fields of the Earth, Venus, Sun and Jupiter in order to reach the Satumian System. (NASA)

(

  ENAs). The Ion and Neutral Caméra (INCA) of the Cassini Magnetosphere Imaging Instrument (MIMI) measured the ENAs flux and modeled Titan's exosphere following the Chamberlain formalism in combination with Ion and Neutral Mass Spectrometer (INMS)

Figure 1 . 9 :

 19 Figure 1.9: Tholin production in Titan's upper atmosphère according INMS findings (adapted from Waite et al. 2006).

Figure 1 . 10 :

 110 Figure 1.10: Titan lakes as discoverer by the Cassini RADAR located in the North polar région taken during the Tl 6 flyby of Titan on 22 July 2006. The smallest is at about 1 km wide, while the biggest are more than 30 km wide (crédit: NASA).

Figure 1 . 11 :

 111 Figure 1.11: Model of the interior of Titan (crédit NASA, http://scitechdaily.com/satums-largest-moon-titanseen-in-unprecedented-detail/).

Figure 2 . 3 :

 23 Figure 2.3: CIRS fields of view with correct relative positions and sizes. FP1 is the far-IR focal plane, while FP3

FarFigure 2 . 6 :

 26 Figure 2.6: CIRS Titan timeline during a flyby. The fields of view of the instrument are shown. With the circle the FOV of FP1 is indicated and with the two parallel lines the linear arrays FOV of FP3 and FP4 are also plotted.

Fig. 2 . 7 :Figure 2 . 8 :

 2728 Fig. 2.7: FP3 averaged spectra as recorded from CIRS. The upper panel (a) shows the FP3 spectral average of September 2010 at équatorial latitudes, while the lower panel shows the FP4 one. The sine wave ripples are located below 600 cm'1 and above 1000 cm'1 at the FP3 sélection and below 1050 cm'1 at the FP4 sélection.

  CIRSDATA database in Meudon (Observatoire de Paris). The available database inferences are: (a) v2.5 current pipeline with local calibration, (b) v3.2 current pipeline with local calibration, (c) DS4000 pipeline with deep-space calibration, (d) Grand Average with unique average calibration over the mission.

Figure 2 . 9 :

 29 Figure 2.9: FP3 average spectrum from the current database v2.5 (in red) with calibration problems. It is for the time period 01 March -30 November 2009, while the latitude range is from 20°S to 30°S at high resolution (0.5

Figure 2 . 10 :

 210 Figure 2.10: The différence between DS4000 (blue) and Grand Average (red) queries is showed by the green graph (Ron Carlson, Ever Guandique pers. comm).

Figure 2 .

 2 Figure 2.11 : (panel a) Average spectrum best fit of the Grand Average sélection file. The data are depicted in red color, while the model is in blue. (panel b): Average Spectrum best fit of the DS4000 sélection file. The data are also illustrated in red color, while the model is in blue.

2. 14 )

 14 hâve lower line radiances between 1020 and 1250 cm'1 than the v2.5 spectra.

Figure 2 . 16 :

 216 Figure 2.16: Fit of the December 2006-FP3 average spectrum from the DS4000 database at 50°N and high resolution (data in red, model in blue).

Figure 2 . 17 :Figure 2 . 18 :

 217218 Figure 2.17: Fit of the December 2006-FP3 average spectrum irom the v2.5 database at 50°N and high resolution (data in red, model in blue).

Figure 2 . 19 :

 219 Figure 2.19: Fit of the January 2010-FP3 average spectrum from the v2.5 database at 50°N and high resolution (data in red, model in blue).

Figure 2 . 20 :Figure 2 . 21 :

 220221 Figure 2.20: Plot of the December 2006-FP3 average spectrum at 50°N and 401 rti. The continuum from 570 to 650 cm"1 of both versions of CIRSDATA is higher in DS4000 (v2.5 in red, DS4000 in blue).

Figure 2 . 22 :

 222 Figure 2.22: Plot of the December 2006-FP4 average spectrum at 50°N and high resolution. The continuum from1050 to 1200 cm'1 of both versions of CIRSDATA is higher in v2.5 (v2.5 in red, DS4000 in blue).

Figure 2 . 23 :

 223 Figure 2.23: Plot of the January 2010-FP4 average spectram 55°N and 401 rti. The continuum from 1050 to 1200 cm'1 of both versions of CIRSDATA are equal (v2.5 in red, DS4000 in blue).

  .2 versus v2.5, DS4000 and DS4000 with phase correction I hâve queried three sélections from CIRSFRANCE database in Meudon on spectra produced during Titan flybys. I hâve asked for the same sélection at FP3 and FP4 focal planes(and FP1 for the first two queries) within three versions of CIRSDATA:(a) the calib_v3.2, which refers to the current pipeline with local calibration which has replaced the previous v2.5, hereafter v.3.2, (b) the calib_v4.2, which is the DS4000 dataset, hereafter DS4000, and (c) the calib_v4.3, which is also the DS4000 dataset but with the application of the new phase algorithm by N. Gorius, hereafter DS4000-PhC.The first sélection picked up spectra recorded by Cassini/CIRS on September 2010 (T72 flyby) focused on équatorial latitudes taken at high spectral resolution (0.5 cm'1). The second sélection contains averaged spectra from the first Titan flyby on July 2004 (T0) in medium resolution (2.5 cm'1) focused on mid Southern latitudes (50°S). It should be noted that the choice of T0 data (July 2004) was due to the fact that only these spectra were uploaded at the moment of the July 2004 query from the DS4000-PhC database. Although we do not take into account the data from this flyby in our radiative transfer calculations, this comparison may be helpful for the algorithm developers. The last sélection refers to January 2007 at 50°N at high resolution from Cassini Titan T23 and T24 flybys.

Fig. 12 .Fig. 2 . 26 :

 12226 Fig. 12.24: FP3 averaged spectra from CIRSFRANCE. The upper panel shows the results of the database inferences (local calibration in blue, DS4000 in green and the new DS4000-PhC algorithm in red). The lower panel shows the différence between the two DS4000 versions (black). Ripples in the recorded radiation are obvious in both spectral edges.

(Fig. 2 . 27 :

 227 Fig. 2.27: As in the previous case, fits of the C2H6 signature in each database extract are plotted in this figure.The model is depicted in blue, while the data are in red. The panels a, b, c show the fit plot versus data of current (v3.2), DS4000 and DS4000 phase corrected database, while the panels d, e and f show the corresponding différences between model and data.

Fig. 2 .

 2 Fig. 2.29: FP4 averaged spectra from CIRSFRANCE. The upper panel shows the results of the database éditions (local calibration in blue, DS4000 in green and the new DS4000 with the phase correction algorithm in red). The lower panel shows the différence between the two DS4000 versions (black).

Fig. 2 .

 2 Fig. 2.30: FP1 averaged spectra from CIRSFRANCE. The upper panel shows the results of the database éditions (local calibration in bleu, DS4000 in green and the new DS4000 with the phase correction algorithm in red).

Fig. 2 . 32 :

 232 Fig. 2.32: Plot of the FP3 averaged spectra of January 2007 at 50°N from CIRSDATA. (a) Comparison between the previous édition v2.5 database (blue) and the current v3.2 database (red). Négative radiances and spike from 570 to 574 cm'1 as well as the spike at 1085 cm'1 hâve been vanished in the current édition, (b) Comparison between the current v3.2 database (red), the nominal DS4000 database (green) and the DS4000-PhC database (black).

Fig. 2 Fig. 2 . 33 :

 2233 Fig. 2.33.

  This chapter introduces the fiindamentals of the radiative transfer theory and the tool I use to model the stratosphère of Titan. It is a radiative transfer code that has been used in the past for retrievals of the thermal and Chemical structure of Titan by the analysis of Voyager, ISO and Cassini data. I describe here the code upgrade we made in order to analyze the recent Cassini results using the latest laboratory spectroscopic line lists and advances in haze modeling. The simulations produced by this code are compared with the actual CIRS data to dérivé the abundances of the atmospheric constituents, which are given in the next chapter. I discuss here the inferences of the température profiles we need to extract first and how they are used

Figure 3 .

 3 Figure 3.1 also illustrâtes the engineering model based on Voyager data of Yelle et al.

Figure 3 . 1 -

 31 Figure 3. 1 -HASI density and température profile. The Yelle et al. (1997) (dashed line) and the Lellouch et al.

3. 2

 2 HASI and CIRS température profiles déterminations and discrepanciesTwo separate instruments of the Cassini-Huygens mission hâve retrieved the vertical température profiles of Titan's stratosphère. CIRS on board the Cassini orbiter senses remotely the emitted infrared radiation during Cassini Titan flyby. HASI on board the Huygens probe measured in situ the température during its descent following the procedure described previously. The HASI data correspond to the Huygens trajectory through the atmosphère near the equator. Instead, CIRS profiles are retrieved remotely from recordings that correspond to different latitudes.

Figure 3 . 2 -

 32 Figure 3. 2 -Methane mole fraction as measured in situ by GC-MS during Huygens descent (adapted from Niemann et al. 2010). With blue diamonds leak L1 data are depicted, while with red squares leak L2 data are plotted.

Figure 3 . 3 -

 33 Figure 3.3-Stratopause location: Comparison among HASI CIRS and IRAM (crédit F. Ferri).

Figure 3 . 4 -

 34 Figure 3. 4 -Comparison of HASI température profile and CIRS and ground based observations (crédit F. Ferri).

Figure 3 . 5 -

 35 Figure 3.5-Température profiles retrieved by assuming both HASI and CIRS-T3 profiles as the initial guess. The lower panel (c, d) focuses on the 1-30 mbar région.

Figure 3 . 6 -

 36 Figure 3. 6 -A complété composite averaged CIRS spectrum of 2007 flybys (T23-T39) at high resolution (0.5 cm'1) at 50°N. The positions of the major molécules and trace gases are illustrated in the figure.

  = kAx , whereris the optical depth of the influenced layer, the above équation takes the general form of the RTE: dr Then, one can write this differential équation such as:

Fig. 3 .

 3 7 depicts en example of the contribution fonctions of methane v4 band Q-branch at 1304.4 cm'1 and its R-branch at 1291.8 cm'1. The contribution of the central methane band peaks at about 0.5 mbar (approximately 225 km).

Figure 3 . 7 -

 37 Figure 3. 7 -The contribution function related to the methane v4 band. The solid [red] line represents the contribution function of the center of methane band, while the dashed line shows the contribution of the left wing.

  , one in the troposphère from 20 to 40 km and the other in the stratosphère, beginning from the height of 40 km and upward. The continuum is fitted by using the collision-induced absorption (CIA) opacity of the combination of the main molécules of the atmosphère (N2-N2, CH4-N2, H2-N2, CH4-CH4, CH4-H2, and H2-H2). The recently upgraded radiative transfer code adapts Vinatier et al. (2012) vertical distribution of the aérosol extinction coefficient for

transfer calculation takes 2 -

 2 3 minutes. When I started my PhD thesis in 2008, the version of the CIRS database also installed in Meudon servers was at its v2.3 version with local calibration applies to the CIRS spectra up to T44 flyby (28/5/2008). Today, the available inventory of CIRS spectra has been expanded to include the most recent Cassini Titan T86 flyby (26/9/2012). While waiting for the validation of the most recent calibration, sevaral CIRS databases are currently available in Meudon: the v3.2 with local calibration, the DS4000 with deep space calibration, the DS4000 with the application of phase correction algorithm and the Grand Average with deep space calibration over several years. The version of the radiative transfer software I received from my predecessors was the ARTT 0.3.8 (Atmospheric Radiative Transfer for Titan, version 0.3.8), based on spectroscopic parameters adopted from GEISA11 (97 and 2003 éditions), HITRAN12 (2004 and 2004 éditions) and line lists from individual researchers. The spectroscopic line parameters contained in these databases change frequently to include necessary updates. One of my tasks was to search for weak features in the CIRS spectra and to try to search for new molécules.The lack of a perfect fit of the major molécules did not allow for such a fine spectroscopic research in the régions around the major émission bands of C2H2, HCN and C2H6. Within the past four years, I hâve searched regularly in the available databases in order to retrieve any updates of the molécules included in the model and to improve the fits. However, the ARTT 0.3.8 version did not allow us to easily include such updates. Several years after the last upgrade of the ARTT radiative transfer code we had to move to a newer version in order to: a) Make the software more user-friendly and improve its flexibility b) Include new aérosol description[START_REF] Vinatier | Optical constants of Titan's stratospheric aérosols in the 70-1500 cm'1 spectral range constrained by Cassini/CIRS observations[END_REF] to fit the spectrum continuum c) Improve and update the spectroscopic parameters 1) New spectroscopic atlases hâve been released with more accurate parameters and enhanced line lists following contemporary laboratory work such as the HITRAN 2008(Rothman et al., 2009) and the GEISA 2009 (Jacquinet-Husson et al., 2011).

  Etude des Informations Spectroscopiques Atmosphériques: Management and Study of Spectroscopic Information (Jacquinet-Husson et al., 2005) 3.4.4 The improved ARTT software The new software upgrade took place in LESIA (Laboratoire d'études spatiales et d'instrumentation en astrophysique) at the Paris Observatory in Meudon, France and in the Astrophysics Laboratory of the University of Athens, Greece. Stefanos Stamogiorgos, software engineer, assisted in this procedure, focussing on the programmatic part of the upgrade.

Figure 3 . 8 -

 38 Figure 3. 8 -Aérosol number density versus the altitude by DISR measurements, adapted from Tomasko et al.

Figure 3 . 10 -

 310 Figure 3. 10 -The real (upper panel) and the imaginary (lower panel) parts of the refractive index of Titan's aérosols (adapted from Vinatier et al. (2012b)). The indices as obtained of[START_REF] Khare | Optical constants of organic tholins produced in a simulated Titanian atmosphère -From soft X-ray to microwave frequencies[END_REF] for tholins are illustrated in blue.

Figure 3 . 11 -

 311 Figure 3.11-Output of the model simulation of ail the major molécules of the FP3 part of CIRS spectrum.

Figure 3 . 13 -

 313 Figure 3. 13 -The output of the model by taking into account only the two 13-C isotopologues of

Figure 3 . 14 -

 314 Figure 3. 14 -The results of the model simulation within 664 and 720 cm'1, where the weak v4 band of C6H6

Figure 3 . 15 -

 315 Figure 3.15-The outcome of the model simulation at FP4, where methane v4 band dominâtes the spectra.

Figure 3 . 16 -

 316 Figure 3. 16 -Model simulation at FP4 with methane v4 band discarded. The signature and the position of 13-C isotopologues of molécules in Fig. 3.15 is obvious.

Figure 3 . 17 -Figure 3 . 18 -

 317318 Figure 3. 17 -Model simulation at FP1. Panel a shows the CO émission, while panel b shows the DCN model signature. The latter has been not yet detected in CIRS data.

Figure 3 .

 3 Figure 3. 19 -Panel (a) illustrâtes the outcome of model simulation of 13-C isotopologues of ethane at FP3. Panel (b) shows acetonitrile model calculations.

Figure 3 . 20 -

 320 Figure 3. 20 -Water vapor isotopologues model calculations at FP1. Only the signature of the major isotopologues has been identified in Titan's atmosphère (panel a).

3. 4 . 7

 47 Improved simulation of the observationsTests on the main molécules We hâve compared the output of our new upgraded radiative transfer code (ARTT 1.0.0) with the previous version (ARTT0.3.8) in order to quantify what has changed. For this purpose, we hâve applied the spectral CIRS FP3 sélection of May 2010 at medium resolution at 50°N.

Figure 3 . 21 -

 321 Figure 3.21-Comparison of the outcome of the two models for the sarne CIRS sélection and comparison to the data (panel a). In panel b, the différence between the two models is illustrated. In panel c, we depict the différence between the model of ARTT 0.3.8 and the data. Similarly, the différence between the model of ARTT 1.0.0 and the data is showed in panel d.

(Figure 3 . 22 -

 322 Figure 3. 22 -HCN lines sériés comparison between the two files adapted in old and new radiative transfer code within the range 640-800 cm"1. Although the line lists files corne from the same version of the HITRAN database (2008) some lines hâve been discarded in the existing file. The HITRAN 2008 line lists used in 0.3.8 version is shown in blue solid line, while the HITRAN 2008 file used in 1.0.0 version is shown in red solid line.

Figure 3 . 23 -

 323 Figure 3. 23 -The same comparison with Fig. 3.22, but focused on the HCN Q-branch.

Figure 3 . 24 -

 324 Figure 3. 24 -GEISA 2009 upgrade for C2H2 (adopted from Fig.8of Jacquinet-Husson et al., 2011).

3. 21 -

 21 b). In this région, except for the P-wing of acetylene, the contribution of V26 C3H8 émission band exists. The différences are indeed significant. Since the acetylene is the same in both models, we hâve to focus on the propane spectroscopic files. This différence can be explained by the fact that the GEISA 2009 lines contains the GEISA 2003 ones multiplied by a factor of 0.420(Jacquinet-Husson et al., 2011) because the 2003 line lists for propane V26

Figure 3 . 25 -

 325 Figure 3. 25 -Focus on the Q-branch of C02 line lists comparison between GEISA 2003 line lists (blue solid line) and GEISA 2009 (red solid line). Here, the significantly increased number of transition lines is depicted.

Figure 3 . 26 -

 326 Figure 3. 26 -Line lists comparison between Arie & Johns (1992) (Elue solid line) and GEISA 2009 (red solid line).

Figure 4 . 2 -

 42 Figure 4. 2 -Sources of ionization in Titan's atmosphère. The EUV solar flux has been calculated for a solar zénith angle at 60°. Electrons T5 represent the nighttime électron distribution in Titan's ionosphère which was observed by Cassini flyby T5[START_REF] Krasnopolsky | A photochemical model of Titan's atmosphère and ionosphère[END_REF].

  5 yrs) and reaches almost the abundances recorded by Voyager 1 in 1980. The end products of the Chemical évolution of methane in the atmosphère are complex refractory organic compounds and ethane. Figures 4.3

- 4 .Figure 4 . 3 -

 443 Figure 4.3-Plot of the relation among observed molécules and parent molécules. Numbers refer to the Chemical cycles described inYung et al. (1984).

Figure 4 . 4 -

 44 Figure 4. 4 -Plot of the hydrocarbon chemistry in Titan according toWilson & Atreya (2004). The stable species

Figure 4 . 6 -

 46 Figure 4. 6 -Flow chart of the model. The model opérâtes in a self-consistent manner (adapted by Lavvas et al. 2008a).

  latitude

Figure 4 . 7 -

 47 Figure 4. 7 -Illustration of the Cassini changing viewing aspect for the same surface intersect latitude at 50°N.

  yearly spectral averages from the Cassini-Huygens nominal mission

Figure 4 .Figure 4 . 8 -

 448 Figure 4. 8 -The contribution functions of CH3D at 1156 cm'1 (black), CH4 -Q branch at 1304 cm'1 (red), CH4 -R branch at 1212 cm'1 (blue) and CH4 -P branch at 1327 cm'1 (magenta). The corresponding sélection is the one of May 2009 at 50°S in high resolution.

Figure 4 .

 4 Figure 4.9 below. The 3-o standard déviation due to noise alone in this région is about 2x10'10 2 1/1

Figure 4 . 9 -

 49 Figure 4. 9 -(Upper) Fit of CII3D at 1156 cm"1 and CH4 at 1304 cm"1. (Lower) Différence between data and the model. The sélection is the one of May 2009 at 50°S in high resolution (0.5 cm"1).

Figure 4 . 12 -

 412 Figure 4. 12 -Latitudinal variations reported in Coustenis et al. (2010b) with 3-a error bars. For latitudes lower than 40°S, only upper limits are illustrated for bcnzene.

  retrieved from 15°S observations by use of CIRS nadir and limb data from the T3 Cassini flyby. The gap in altitude between nadir and limb measurements in the stratosphère was fulfilled by interpolation using as initial guess the updated Voyager radio occultation profile. The characteristics of this profile are given in the online supplementary material of[START_REF] Flasar | Titan's atmospheric températures, winds, and composition[END_REF] paper and in Chapter 3. The results are shown in Figure4.13a-e below. The précision of these température profiles dépends on the number of spectra and the calibration accuracy. The random error imported in the measurements due to the instrumenté noise is at about 0.1 K. Systematic errors corne from uncertainties in methane abundance and the absorption coefficients(Coustenis et al., 2007). The température error bars are about 0.7 K at 1 mbar and about 5 K at 5 mbar(Achterberg et al., 2011).

Figure 4 .

 4 Figure 4.13f shows the température profiles retrieved from limb data obtained by Sandrine Vinatier. Figures 4.13a-e show the results on the température as a function of latitude from nadir data inferences described in Table4.7-4.10, using the altitude information

Figure 4 .

 4 13a,b (upper panel) shows the température structure inferred at two latitudes (50°N and 30°N) at certain dates during the Cassini mission.

  cooling from 2007 up to 2009, while in 2010 a slight warming is observed again at both 50°N and 30°N. Température decrease (by 8 K maximum) with time in the lower mesosphere is also observed at the équatorial and southem latitudes. No significant changes are found in the équatorial mid-stratospheric températures (Figure 4.13c).Similarly, in the southem hemisphere (at 30°S and 50°S -Figure4.13d,e respectively), small température variations are observed that are always less than 10 K and more typically 2-4 K between 0.5 and 10 mbar within the error bars. Mid-stratospheric measurements for southem latitudes at 30°S(Figure 4.13d) show cooling between 2010 and early 2012 by 3 K, while no significant variations occurred from 2006 to 2008. More importantly, the general shape of the thermal structure is changing and we find a more isothermal profile at higher latitudes losing the marked stratopause présent in earlier years.

Figure 4 . 13 -

 413 Figure 4. 13 -Retrieved thermal profiles from CIRS nadir data at (a) 50°N, (b) 30°N, (c) equator, (d) 30°S, (e) 30°S and (f) from limb data for the 2006-2010 limb data and from several dates. Panel f illustrâtes the température profiles corresponding to June 2010 at 50°S (pink solid line), 1°N (red solid line), 51°N (violet solid line) from CIRS limb observations acquired with a 15.5 cm'1 spectral resolution. They are compared with thermal profiles of the 2006-2007 period (northem winter) at similar latitudes. For the nadir température profiles the typical error bars uncertainty are about 0.7K at 1 mbar and about 5K at 5 mbar, while for limb retrievals is 1 -2K.Dashed lines indicate the altitude levels above or below which, the température has higher uncertainties or is not accurately constrained(Bampasidis et al. 2012a).

  n) we describe the weighting function and B(z) the Planck fonction. The contribution fonctions of the molécules présent in the FP3 région are plotted in Figures 4.14

  -

Figure 4 . 14 -

 414 Figure 4. 14 -Contribution functions of most of the hydrocarbons, nitriles and carbon dioxide that hâve émission bands in Titan's CIRS FP3 spectra at 50°N high-resolution data.

Figure 4 . 15 -Figure 4 . 16 -Figure 4 . 17 -

 415416417 Figure 4. 15 -Contribution function of most of the hydrocarbons, nitriles and carbon dioxide that hâve émission bands in Titan's CIRS FP3 spectra at 30°N high-resolution data.

Figure 4 . 18 -

 418 Figure 4. 18 -Contribution function of most of the hydrocarbons, nitriles and carbon dioxide that hâve émission bands in Titan's CIRS FP3 spectra at 50°S high-resolution data.

Figure 4 .

 4 Figure 4. 19 -(Upper panel) The best fit achieved (in blue the model) between 600 and 1000 cm'1 for the December 2007 Titan spectrum at 50°N (in red). (Lower panel) Plot of the différence between model and data.The peaks at 765 and 955 cm'1 are caused by electrical artificial inferences -spikes (see Chapter 2).

Figure 4 .

 4 Figure 4. 20 -(Upper panel) The best fit obtained (in blue the model) between 610 and 690 cm'1 for the March 2009 Titan's spectrum at 50°N (in red). (Lower panel) Plot of the différence between model and data. The C4H2, C3H4, HC3N and C02 émission bands are at 628, 633, 663 and 667 cm'1 respectively.

Figure 4 .

 4 Figure 4. 21 -(Upper panel) The best fit found (in blue the model) between 690 and 750 cm"1 for the March 2009 Titan's spectrum at 50°N (in red). (Lower panel) Plot of the différence between model and data. The HCN and C2H2 émission bands are at 713 and 730 cm'1 respectively. The isotopologues of H°CN, HC15N and 13C12CH2are also fitted at 706, 712 and 729 cm'1 respectively.

Figure 4 .

 4 Figure 4. 22 -(Upper panel) The best fit found between 750 and 1000 cm'1 for the December 2009 Titan's spectrum at 30°15I (in red the data, in blue the model). (Lower panel) Plot of the différence between model and data. The négative peaks at 765 and 955 cm'1 are caused by spikes in the spectrum. The 13-C isotopologue of C2H6 (I3C12CH2) is included.

Figure 4 .Figure 4 . 23 -

 4423 Figure 4. 23 -CIRS NESR as a fiinction of wavenumber (provided by Dr. R. Carlson, pers. communication).

Figure 4 .

 4 Figure 4. 24 -Time-latitude composition variations for the major trace gases of Titan's stratosphère: (a) C2H2, (b) HCN, (c) C4H2, (d) C3H4, (e) C2H6, (f) C2H4, (g) C3H8 and (h) C02 with 3-c estimated error bars. The latitudes mapped from 2006 to 2012 are: 50°S (violet), 30°S (light blue), equator (red), 30°N (green) and 50°N (blue).Connected filled circles are high resolution observations(0.5 cm'1), while open circles are medium resolution data (2.5 cm1) for 2008, 2009 and 2010 (which sometimes coincide with the higher resolution values). The 3-g estimated error bars are indicated(Bampasidis et al., 2012a).

Figure 4 . 25 -

 425 Figure 4. 25 -Comparison of the vertical température profiles betwcen limb and nadir sélections at northem spring equinox. The nadir température profile has been retrieved by R. Achterberg by using 0.5 cm'1 spectra, while the limb one has been retrieved by S. Vinatier by using 15 cm'1 spectra. The profiles présent a 10 K différence at 1 mbar.

Figure 4 . 26 -

 426 Figure 4. 26 -Fits of the limb data (in black) of the July 2009 sélection by using the température profile of June 2009 nadir query to model the observations (in red). Data and model fit only at 213 km or 0.3 mbar.The plot was provided by S. Vinatier. Radiance is in xlO"7 Wcm^sr'Vcm'1, while the Wavenumber is in cm1.

Figure 4 . 28 - 29 -Figure 4 . 30 -

 42829430 Figure 4. 28 -Contribution fonctions of V4 methane band at 1304 cm'1 and its wings derived from the limb température profile. CH4(G) is the left wing, while CH4(D) is the right wing.

Figure 4 .

 4 31 illustrâtes the outcome of this simulation.

Figure 4 . 31 -

 431 Figure 4. 31 -Fit of the nadir sélection of June 2009 at 50°N when the thermal profile retrieved from limb measurement was adapted. Values predicted from General Circulation Models (Rannou et al. 2005) hâve been used. The misfit due to the température profile is obvious.

Figure 4 . 32 -

 432 Figure 4. 32 -Fit of the nadir sélection of June 2009 at 50°N when the thermal profile retrieved from nadir measurement was adapted. Vertical distributions provided from S. Vinatier hâve been used. Now the model spectra fit well the data.

Figure 4 . 33 -

 433 Figure 4. 33 -Daily averaged insolation as a function of time at the top of the atmosphère. Cassini entered into SOI at Ls 293°(adaptedfrom Lora et al., 2011).

Figure 4 . 34 -

 434 Figure 4. 34 -The Solar Cycle from 1980 (Voyager encounter at the Saturnian System) up to date. Cassini began touring the Saturnian System in mid-2004. This plot shows the Composite Mg II index which is frequently used as a proxy for spectral irradiance variability. It refers at 1 AU and it provides the solar activity at EUV and UV As I hâve discussed with Prof Moussas and Dr. L. Didkovsky (Univ. of Southern California/SSC) Mg II index represents the UV solar radiation (e.g.[START_REF] Viereck | The NOAA Mg II core-to-wing solar index: Construction of a 20-year time sériés of chromospheric variability from multiple satellites[END_REF].

Figure 4 .

 4 35 shows the contribution function of water vapor. The contribution functions peak at 0.01 mbar (about 400 km), c.f.Coustenis et al. (1998).

Figure 4 . 35 -

 435 Figure 4. 35 -Contribution functions of water vapor (dashed line) and the vertical profile used for simulating the best-fit (adapted from Coustenis et al. 1998).

  observations in search of water vapor Water vapor has rotation lines in the FP1 région of CIRS. CIRS/FP1 is a detector with a circular field of view and probes the far infrared région in the range of 10 to 600 cm'1. The stronger water lines are located from 90 to 260 cm'1. The selected datasets were queried from December 2004 up to December 2008. Three sélections cover Titan up to 30°N (80°S-45°S, 45°S-10°S, 0°-30°N) in nadir viewing mode (Table4.13).With Dr. A. Coustenis, we hâve gathered large nadir FP1 spectral averages within the SCET14 limits, as required by V. Cottini and analyzed them following a study already started about 2 years ago aiming to explore the FP1 région, but which was impeached by the difficulty to model the aérosols which are very important in that région. By focusing on the water lines and simulating summarily the continuum level, I hâve injected the température profile retrieved by R. Achterberg from the corresponding FP4 spectrum taken at the same time and at the same latitudes and conditions of geometry in order to fit the FP1 observations. I hâve simulated the observations and when the best fit was achieved, I hâve inferred the abundances. Also into account were taken the other molécules présent in the région, namely C4H2, C2N2 and C3H4 and furthermore, we hâve extended the range of our calculations 300 (check) cm'1 région, while the Cottini et al. paper is constrained from 150 water vapor détection by ISO with CIRS Water has émission bands at the FP1 spectral range of CIRS at 150.

cm' 1 .

 1 The contribution functions of the water vapor lines are plotted in Figure4.36 below. Ail of them peak at about 50 km.

Figure 4 . 36 -

 436 Figure 4. 36 -Contribution functions of water vapor. Ail the émission bands peak at about 50 km.

Figure 4 .

 4 37 below, shows the best-fits of the sélection 0°-30°N when using a constant-to-height mixing ratio for water vapor at lxlO'10. The vertical black lines in the upper plot indicate the positions of water lines. The model fits well the water lines. The peak of the data at 220 cm'1 is discussed in Section 4.10.

Figure 4 .

 4 Figure 4. 37 -(Up) Best-fit plot at the 0°-30°N latitudinal bin achieved for water vapor abundance of lxlO'10.

Figure 4 .

 4 Figure 4. 38 -(Up) Best-fit plot at the 45°S-10°S latitudinal bin achieved for water vapor abundance of lxlO'10.

Figure 4 .

 4 Figure 4. 39 -(Up) Best-fit plot at the 45°S-10°S latitudinal bin achieved for water vapor abundance of 2xlO~10.

4. 40 )

 40 , while the value of 2xlO'10, fits better the 250 cm"1 line but not the others (Fig.4.41).

Figure 4 .

 4 Figure 4. 40 -(Up) Best-fit plot at the 80°S-45°S latitudinal bin achieved for water vapor abundance of lxlO'10.(Down) Plot of the différence between model and observations. The vertical black lines depict the positions of water lines in the spectrum. Except for 250 cm"1, ail the water lines are fitted.

Figure 4 .

 4 Figure 4. 41 -(Up) Best-fit plot at the 80°S-45°S latitudinal bin acbieved for water vapor abundance of 2xlO'10.

4. 9 . 2

 92 Modeling benzene's vertical distribution In collaboration with Dr. P. Lawas, I hâve modeled the vertical distribution of CôHô in the atmosphère of Titan, expending the work presented in Coustenis et al. (2010b) from a latitudinal point of view.

  photochemical calculations. P. Lawas has updated the photochemical model with the necessary ion reaction based on Vuitton et al. (2008) (see Figure 4.42). Following our recent work in Coustenis et al. (2010b) on CIRS data analysis and photochemical modeling, we hâve thus applied this advanced version of Lawas et al. (2008a;b) 1-D model, which is described in section 4.3.2. This advanced model then includes ion-chemistry at different latitudes of Titan's globe.

Figure 4 . 42 -Figure 4 .

 4424 Figure 4. 42 -Principal chemistry reaction sequence for benzene (adapted from Vuitton et al. 2008).

Figure 4 .

 4 Figure 4. 44 -(a-d) Vertical distribution of benzene's abundances retrieved by the advanced 1-D photochemical model by Lavvas et al. (2008a;b) at the équatorial and sourthem Titan's hemisphere. CIRS results are plotted

Figure 4 .

 4  Vertical distribution of benzene's density for latitudes, (b) Vertical distribution of benzene's abundance for latitudes(Bampasidis et al., 2010).

4. 10

 10 Winter condensâtes évolution between the North and South pôle on Titan By comparing large yearly averages of CIRS/FP3 at different time periods, I hâve discovered a dramatic déplétion in the abundance of HC3N at 50°N in 2010 of more than 50% (see Fig.

4. 46 )

 46 compared to 2006-2009 values. This was indicative of changes in the North pôle concentrations at recent times.

Figure 4 .

 4 Figure 4. 46 -CIRS/FP3 annual data at 50°N. The déplétion of HC3N abundance is obvious in 2010 spectra.

Figure 4 . 47 -

 447 Figure 4. 47 -Haze observations in seasonal scale (adapted from Jennings et al. 2012, in press.).

  the relative température profile derived from the FP4 spectral sélection in the FP3 and then fitted the major molécules during a best-fit itérative process. The results of the radiative transfer calculations are plotted in Figs.4.48 and 4.49.

Figure 4 . 48 -

 448 Figure 4. 48 -Comparison between CIRS nadir spectra and radiative transfer models. Hydrogen cyanide and the

Figure 4 . 49 -

 449 Comparison between CIRS nadir spectra and radiative transfer models. We hâve derived only a 3a upper limit of 8x10'10 for the isotopologue 13C180160 of carbon dioxide. Top (a): Averaged spectra of T50 to T64 flybys at high resolution (0.5 cm'1) (in blue) at 50°N, best-fit model with main gases and isotopologues (in red) and best-fit model without |JC1X0160 isotopologue at 662.5 cm'1 (in green).Bottom (b): Différence between model and data for both cases: with ail gases and isotopologues (in green) and without 13C180160 at 662.5 cm"1 (in blue).

Voyager 1

 1 Infrared Radiometer Spectrometer (IRIS) measurements from 1980, Cassini Composite Infrared Spectrometer (CIRS) continuous recordings from 2004 to 2010 and the intervening ground-based and space-bome observations with ISO provide almost a complété picture of the stratospheric évolution within a Titanian year. The VI measurements were taken at the same season compared to the one of Cassini/CIRS in mid-2009. The next step is to compare the previous results in température and composition variations from Voyager 1/IRIS and ISO with Cassini/CIRS ones. This can be complemented with the ground-based observations I hâve assembled in Table 4.4. Thus, we can describe the thermal and Chemical évolution of Titan's stratosphère within a Titan year (29.5 years) and evaluate General Circulation Models prédictions.

Figure 5 .

 5 2 below shows the global map of Titan obtained by the RADAR instrument.

Figure 5 . 2 -

 52 Figure 5. 2 -Global map of Titan's surface obtained from the beginning of the mission up to flyby T50 (February 2009). It shows lhe coverage derived from RADAR synthetic aperture (SAR) images overlaid onto the ISS map, in a simple cylindrical projection centered at 0°N and 180°W. Full resolution SAR images are shown in gray, lower resolution liigh altitude SAR (HiSAR) observations in blue, and the ISS base map in brown. Blackand dark blue correspond to -20 db, white to 5 db[START_REF] Stephan | Mapping Products of Titan's Surface[END_REF].

5. 4 . 2

 42 Study of Mayda Insula in Kraken Mare vvith an eye to morphotectonicsDuring the near-polar Titan T25 flyby onFebruary 22, 2007, the Cassini/SAR instrument observed a big island or a peninsula in the northem shoreline of Kraken Mare, the Mayda Insula(Fig 5.4). The same région was covered by the T28 flyby(Fig 5.5).

Figure 5 . 4 -

 54 Figure 5.4-Part of the northem shoreline of Kraken Mare from Cassini/SAR taken during the T25 flyby.Mayda Insula is in the center of the image (79°N and 310°W). North is to the left (Crédit: NASA/JPL-Caltech/ASI: PIA09180).

Figure 5 . 5 -

 55 Figure 5. 5 -Mayda Insula from the Cassini Titan flyby T28. This figure is part of the swath which began at 20°S and 37°W and continued in the north-northeast direction. T28 flyby overlaps T25 (Crédit: NASA/JPL-Caltech/ASI: PIA09217).

Figure 5 . 6 -

 56 Figure 5. 6 -Mayda Insula drainage basins and stream classification (Cassini/SAR image PIA09180-PIA09217).

Figure 5 .Figure 5 . 7 -

 557 Figure 5.7-2nd order Northern and Southern basins classification according to their orientation.

Figure 5 . 8 -Figure 5 . 9 -

 5859 Figure 5. 8 -Histogram of the Island-region A.

Figure 5 .

 5 Figure 5. 10 -Histogram of the SW Mountains-region C.

Figure 5 . 11 -

 511 Figure 5. 11 -Histogram of 2nd order Streams orientation of Northern and Southern poorly drained basins.

5. 6 . 2

 62 Speckle noise in Cassini/SAR images RADAR observations are almost unaffected by the atmospheric conditions. The recorded echoes can distinguish different types of surface régions through the measurement of their emissivity or radar reflectivity. The instrument measures the normalized backscatter crosssection of Titan's surface.

Figure 5 . 14 -Figure 5 . 15 -

 514515 Figure 5. 14 -The initial Cassini/SAR image PIA08630 (Crédit: NASA/JPL-CalTech/ASI).

Figure 5 . 16 -

 516 Figure 5. 16 -The segmented image after filtering. Dark lake with low backscatter are illustrated with blue, the granular lakes with green and the local background is colored red.

Figure 5 . 17 -

 517 Figure 5. 17 -Application of the filtering and segmentation procedure to Sinlap crater: Initial image (left) PLA07368 (Crédit: NASA/JPL-CalTech/ASI). Filtered image (center) and Segmented image (right).

  the paper ofSolomonidou et al. (2010), which I hâve co-authored, we hâve worked on the two recognized (at that time) cryovolcanic candidate areas, Tui Regio and Hotei Regio, on which we applied the Principal Component Analysis tool. The paper was published in the anniversary volume of the Hellenic Journal of Geosciences19 for the 40 years since the foundation of Geology and Geoenvironment Department of the University of Athens. The main goal was to identify régions of interest that présent brightness or color alterations using VIMS data within the possible cryovolcanic areas that suggest déposition of material coming from the interior and having a different Chemical composition than the surrounding surface.Hence, a variety of materials within a certain geological terrain suggests either past or still ongoing activity of different and distinct geological processes or the déposition of different and distinct material from various sources.The surface contributes significantly to Titan's methanological System[START_REF] Hayes | Transient Surface Liquid in Titan's Polar Régions from Cassini[END_REF] and its lakes are an intégral part of the methane cycle. By studying the filtered images of the same régions of Titan, taken at different time periods, we can study the existence of possible ephemeral or seasonal changes on the surface. These studies are also of interest to the astrobiological research, which I discuss in Chapter 6, the design of future mission instrumentation and the détermination of régions of interest for future landing sites, which I describe in Chapter 7.19 formely Annales Géologiques des Pays HélléniquesChapter 6Habitability of Titan and other icy moons Despite great strides in biosciences, the roots, the sources and the initial conditions of life emergence on Earth remain largely unknown. Due to the erosional and tectonic activity occurring on Earth, the traces of the geological record hâve been inevitably erased.Extraterrestrial environments can provide the necessary information for finding the missing link, looking at the conditions for the emergence of life or at least for characterizing Earth as a habitat.Current missions and data analysis revealed that among the main candidates for finding signs of past and/or current life within our Solar System, besides Mars, are the icy moons of the giant planets Europa, Ganymede, Callisto, Titan and Enceladus. These planetary bodies may host the proper conditions for the emergence or the maintenance of life, providing the concentration of the necessary ingrédients and the proper Chemical inventory for biochemical reactions. Due to their high astrobiological significance, they hâve been selected as targets of future exploration programs, ESA's Jupiter Icy Moons Explorer (JUICE) and and the NASA-ESA joint study of the Titan-Satum System Mission. Indeed, one of the major objectives of both missions is to explore the habitability potential of these icy moons. It was that astrobiological significance that triggered my interest and formed my final decision to engage with Titan studies. For this reason, I hâve studied and compared Titan and Enceladus, in collaboration with Dr. A. Coustenis, Prof. F. Raulin and A. Solomonidou considering the recent discoveries of the Cassini-Huygens mission. This study entitled as "Life in the Satumian Neighborhood" has been published as a chapter in Cellular Origin, Life in Extrême Habitats and Astrobiology

Figure 6 . 1 -

 61 Figure 6. 1 -The habitable zone in our galaxy. Habitability should not be restricted to the places where liquidwater may exist at the surface of moons and planets. A much larger domain exists beyond the snow-line, where very large liquid réservoirs can exist below the icy crusts of the moons and planets, see alsoLammer et al., 2009.

Figure 6 . 2 -

 62 Figure 6. 2 -Schematic illustration of the connections between Titan's interior, surface, atmosphère, and cosmic environment. Images show lakes close to the north and south pôles, mid-latitude terrains with dunes, and fluvial features carved in the ice crust. It is based on an original figure from Lunine (1993) with Cassini VIMS, Radar and Huygens DISR images added (adapted from the TSSM Final Report).

Figure 6 . 3 -

 63 Figure 6. 3 -Composition of Titan's aérosols from the Huygens/ACP(Raulin, 2008a).

Figure 6 . 4 -

 64 Figure 6. 4 -Illustration of methane sink in Titan(Atreya et al., 2006).

Figure 6 . 6 -

 66 Figure 6. 6 -Linear dunes on Titan from T8 flyby, Oct. 2005 at 8°S, 264°W[START_REF] Radebaugh | Planetary science: Titan's sticky dunes?[END_REF].

-Figure 6 . 7 -

 67 Figure 6. 7 -Map of almost 655 lakes and sea features observed by the Cassini/Radar. These maps are in azimuthal projection at the North Pôle of Titan. The projection A (left) shows the radar swath mosaic up to May 2007 flybys. The projection B represents the spatial distribution of mapping units. Lake Michigan is illustrated

  Tobie et al. (2005) suggested a layered interior structure, consisting of a rocky core overlaid by high pressure ice, a liquid layer overlaid by low pressure ice and finally a solid icy crust.Fortes (2000) claimed that undemeath Titan's icy crust, at 200 km, lies an ammoniawater solution océan.Mitri et al. (2008) assumed pockets of methane clathrates trapped within an ammonia-water océan.

Figure 6 . 8 -

 68 Figure 6. 8 -Model of the interior structure of Titan. The ammonia-water océan lies at about 250 km in depth, below the crust and the methane clathrate hydrate ice layer (crédits: L. Spilker).

  Clark et al. (2010) focused on the complex Chemical activity on the surface of Titan, based on Cassini data.

Figure 6 . 9 -

 69 Figure 6. 9 -Enceladus from Cassini/ISS. It is a mosaic of the moon's leading (or western) hemisphere. The Southern polar régions are geologically younger than the northem ones. This image was taken during the November 21, 2009 flyby over Enceladus (crédits: PIA11684, NASA/JPL/Space Science Institute).

Figure 6 . 10 -

 610 Figure 6. 10 -The south pôle of Enceladus with the linear dépréssions "Tiger Stripes" indicated by the red circle (captured by Cassini/ISS). The plumes are ejected from these fractures (crédits: PIA12566, NASA/JPL/Space Science Institute, Porco et al. 2006).

  Cassini/INMS data identified H2O vapor as the prédominant component of Enceladus' plumes, CO2 as the second most abundant, in addition to methane and trace quantities of acetylene and propane and other organics. During the flyby of 9 October 2008, Cassini dived into the south polar plume at a distance of 339 km and INMS recorded the presence of

Figure 6 . 12 -Figure 6 . 13 -

 612613 Figure 6. 12 -Schematic model of the interior mechanism occurring in Enceladus' south polar région: #1 indicates the internai water océan, #2 shows the fracture from where the water océan exsoluted gases rises towards the surface. #3 indicates thermal anomalies where heat from the water is being conducted through ices to the surface and radiated to space (#4). The label 5 marks the plume chamber which feeds the geysers. The label 6 indicates the retum of the water back to the océan after the éjection of the volatiles (adapted from Matson et al. 2012).

Figure 6 .

 6 Figure 6. 14 -Natural color view of Ganymede from Galileo/SSI. The picture was taken during the first Galileo encounter of the satellite. North is on the top of the image and the Sun illuminâtes the moon from the right. The bright spots are young impact craters and their ejecta. The dark areas are geologically older, while the light areas are younger, tectonically deformed régions. The brown/gray color is due to mixtures of rocky materials and ice (crédit: NASA/JPL-PIA00716).

Figure 6 . 15 -

 615 Figure 6. 15 -The interior structure of Ganymede (adapted by Nakamura & Ohtani, 2011).

Figure 6 . 16 -

 616 Figure 6. 16 -The training hemisphere of Europa from the Galileo mission in natural color. Dark brown areas represent rocky material from the interior. The fractures in the crust are also shown (crédits: Galileo/SSI-PIA00502, NASA/JPL/DLR).

Figure 6 .

 6  The layered interior of Europa and (right) a proposed model for its interior suggesting a thin icy crust of 200 km (crédits: NASA/JPL/Galileo/SSI-PIA01669).

Figure 6 . 18 -

 618 Figure 6. 18 -The only complété global color image of Callisto obtained by the Galileo orbiter. Callisto is uniformly cratered (the bright scars on the darker surface) (crédits: Galileo/SSI/PIA03456 NASA/JPL/DLR).

Figure 6 . 19 -

 619 Figure 6. 19 -The interior of Callisto from Galileo data. An internai océan lies probably beneath the 200 km icy crust (crédit: NASA/JPL/PIA01478).

  evolved since its formation 4.5 billion years ago and most of the traces of the initial environmental conditions hâve been erased. It is therefore difficult to retrace the processes of the emergence of life on Earth around 4 billion years ago, or even of the formation and évolution of its atmospheric organic content. Astrobiology could provide crucial information from extraterrestrial locales with similarities to our planet's early stages.According to theLammer et al. (2009) classification, Saturn 's Titan and Enceladus and Jupiter's Ganymede, Europa and Callisto can be considered as habitats, since they hâve encapsulated internai water deposit. The question is if these moons hâve compositional, structural and other similarities among them or compared to the Earth with astrobiological implications. In[START_REF] Coustenis | Life in the Satumian Neighborhood[END_REF] and in[START_REF] Solomonidou | Water Océans of Europa and Other Moons: Implications For Life in Other Solar Systems[END_REF] we try to provide answers. Titan's surface differs significantly from Enceladus and the other

  : (a) The heat production, which mainly originates front radiogenic heating or other triggering mechanisms (e.g.McKinnon, 1999, Tobie et al, 2005 ); other heat sources are the dissipation of tidal energy due to the orbital interaction between the satellites and their planets, the exothermal geochemical production of heat like hydration and crystallization ofsolids (e.g.Sohl et al., 2010, Hussmann et al., 2010); (b) The efficiency of heat transfer, which is based on thermal diffusion and thermal convection (e.g.Hussmann et al., 2010);(c) The components that decrease the melting point of ice and support the océan 's liquid State (e.g.Sohl et al., 2010, Tobie et al., 2010); an antifreeze component like a solution of water with ammonia should exist in the océan in order to remain in the liquidState; and (d)

  , followed by Satum are the largest planets of the Solar System. They both hâve their own System of numerous natural satellites in multivariable sizes (more than 60 moons each). Unique among these moons are Jupiter's Ganymede, Callisto and Europa and Satum's Titan and Enceladus. Ail of them are active worlds likely to hâve a differentiated internai structure with encapsulated océans. The scientific interest in these objects is considérable, conceming their corrélation with the origin of the Solar System, the terrestrial-like processes and features they exhibit, as well as their habitability potential.The planetary environments of Jupiter and Satum hâve been studied by the missions Pioneer 10 and 11, Voyager 1 and 2, Ulysses, Galileo, Cassini-Huygens and New Horizons for Jupiter and Pioneer 11, Voyager 1 and 2 and Cassini-Huygens for Satum. The Galileo mission to Jupiter and the Cassini-Huygens mission to Satum hâve significantly advanced our knowledge about the biggest outer planets of the Solar System and their Systems. The Cassini-Huygens mission, in particular, studies in depth Satum, its satellites and its rings and will continue touring up to 2017. However, numerous questions hâve been raised so far, which cannot be addressed by these missions, mainly due either to the instrumentations' limitations. Without doubt, both Systems of Jupiter and Satum should be revisited by new missions.22 Pluto orbits also outside the asteroid belt, but the International Astronomical Union (IAU) categorized it as a Titan and Europa were identified, among others, as main targets for future exploration by NASA in their 2003 and 2011 Decadal Surveys. ES A has recently performed extended studies for the outer planets satellites exploration (TandEM, TSSM, EJSM) and finally selected JUICE, as the first large mission to be launched within the Cosmic Vision 2015-2025Program. In this Chapter, I présent some of the studies for future missions in which I hâve had the chance to somehow contribute or to participate and in particular some individual experiments concepts for icy moons that I hâve explored and proposed.7.1Recent mission studies to Jupiter: From EJSM to JUICE Jupiter is the biggest planet of our Solar System. It has a mass of 1,898.6x1024 kg and a radius of 71,492 km at the equator. Its mean distance to the Sun is 778.57xl06 km and it complétés a full révolution of the Sun every 11.86 years. Jupiter has the strongest magnetic field of the Solar System (4.2 gauss on the surface at the equator, 14 gauss on the surface at the North pôle and 11 gauss on the surface at its South pôle) and the biggest magnetosphere, which extends up to 75 Jupiter radii.Jupiter is considered as a gas giant along with Satum, Uranus and Neptune since it mainly consists of a gaseous envelope, without a surface. Its atmosphère is hydrogendominated (89.8%) with hélium (at about 10%), while the rest of the gases are methane, ammonia, hydrogen deuteride, ethane and water vapor in traces.It has 63 natural satellites and among them the four so-called "Galilean" (because they were discovered by Galileo Galilei in 1610), Ganymede, Europa, Callisto and Io. Ganymede is the biggest satellite of our Solar System. The Jovian System has been investigated by the spacecraft Pioneer 10 (1973) and 11 (1974), Voyager 1 (1979) and 2 (1979), Ulysses (1992 and 2004), Cassini-Huygens (2000) and New Horizons (2007), during their flyby maneuvers. The Galileo spacecraft entered into Jovian orbit in December 1995 and operated until September 2003, when it fell onto the giant planet. Galileo also released a probe, which entered the planet's atmosphère in December 1995 and collected measurements for 57.6 min until the atmospheric pressure crushed it.

Figure 7 . 3 Figure 7 . 4 -

 7374 Figure 7. 3 The JUICE spacecraft design (crédit: ESA)

  is hélium. It expériences the most intense storms of the Solar System with its équatorial zonal winds to présent a slowing from 450 m/s down to 250 m/s in respect to the Voyager era. Satum has 62 natural satellites known today and a unique planetary ring System. Its rings are mainly composed of ice particles, rocky débris and dust and they are separated into nine continuous rings and three discontinuous arcs. I got involved in future mission planning to the Satumian System in February 2007, during my MSc studies. I hâve participated in the starting workshops for planning future missions in response to ESA's Cosmic Vision 2015-2025 Call, conducted in Meudon-Paris Observatory from 12 to 15 February 2007. There were two workshops in which Prof. Xenophon Moussas (my PhD co-supervisor) and myself were invited to attend, representing the Astrophysics laboratory of the University of Athens. The outcome of the work of the Consortium formed at the time were two separate L-scale (large) mission proposais to Satum,

  planet. Kronos would be an equally ambitious mission compared to Cassini-Huygens with instrumentation improved by several orders of magnitude in resolution and enhanced capabilities. A. Coustenis was co-Lead of the proposai with Bernard Marty and Tristan Guillot. The Kronos study is well-described in the papers ofMarty et al. (2009a;b) published in Experimental Astronomy Journal, which I co-authored.The ESA contribution to the mission were the two probes, while NASA (via its New Frontiers program) would provide the carrier which would also perform some remote measurements and telecom. Kronos would analyze in depth the Chemical and isotopic composition of the Satumian atmosphère. Such studies provide essential dues about the formation and the évolution of Satum and the Solar System in general, since the giant planets are made of protosolar nebula gases. The study of the meteorology and atmospheric dynamics would be also a major goal of the mission in order to understand the évolution of its atmosphère since the planet's formation. The Kronos orbiter and probes would measure the température vertically and zonally, the winds, the cloud properties, the radiative balance and the convection mechanisms.

Figure 7 . 5 -

 75 Figure 7. 5 -The cross section of the Galileo probe. The Galileo probe's héritage would be used as the basis of the Kronos probes design (Galileo Probe Décélération Module Final Report, Doc No. 84SDS2020, General Electric Re-entry Systems Operations, 1984; Project Galileo Mission and Spacecraft, 1983) (Marty et al., 2009a).The carrier would arrive to Satum after an interplanetary joumey by taking advantage of Earth, Venus and Jupiter gravitational assists as Cassini did. The initial proposai mission

2 The

 2 Indeed, the Galileo probe survived for about 58 minutes before melting from the high atmospheric pressure of 20 bar and the température of about 150°C inside Jupiter''s high pressure environment, using a Wheatstone bridge circuit configuration.My main concem was also to minimize the probe's payload weight and that research led to the study of implementing Micro-Electro-Mechanical System (MEMS) technology in future missions to outer planets. A further study of MEMS devices and their implémentation as distinctive future experiments in icy moons are described later in this Chapter. 32 http://www.kistler.com/mediaaccess/000-614e-09.09.pdf 7.2.Titan and Enceladus Mission (TandEM) My involvement in future mission planning to Titan and Enceladus began in early 2007 in Meudon, when Prof. X. Moussas and I participated to the Workshop of the TandEM Consortium conducted in Salle des Conférences du Château of the Paris-Meudon Observatory (Atelier CIAS). This Workshop set up the baselines for the study of a future flagship mission dedicated to these planetary objects. TandEM was a multi-national cooperative effort, in which 155 expert scientists and engineers participated (Coustenis et al., 2009a). The Proposai Lead was Athéna Coustenis. The TandEM mission was planned to be launched around 2018 or later and wanted to address several of the Cosmic Vision 2015-2025 call thèmes with the help of an orbiter, a balloon, and several landers/penetrators to be delivered to the satellite around 2030. Its design was based on the Cassini-Huygens architecture and it was expected to surpass the latter's legacy by exploring Titan and Enceladus in full close-up and in situ coverage over long periods of time, which is not currently possible. The study's task assignments of the Astrophysics Laboratory of University of Athens were the study of Titan surface, the mission architecture and the outreach campaign. During my MSc dissertation, I studied the surface of the northem part of Titan and tried to identify tectonic signs on the crust. However, we found that the resolution of the available Cassini/Synthetic Aperture RADAR (SAR) images was not sufficient to produce topographie and géologie maps. For this reason, I took advantage of the involvement of the University of Athens in TandEM's Titan surface and mission architecture working groups to suggest improved RADAR instrumentation characteristics for the mission. A high resolution RADAR on an orbiter revolving around Titan for a long period of time would help in the study of surface évolution and its interactions with the lower atmosphère. TandEM's RADAR was expected to provide solutions to a sériés of under-study matters such as: 1. Détermination of the active (or not) local tectonic field 2. Détermination of the âge of Titan's surface 3. Chemical composition and source of liquid surface deposits 4. Depth measurement of lakes 5. Corrélation between the surface lithology and the surface features.In order to understand the internai structure of Titan Dr. A.Coustenis and I worked together on including seismic instrumentation in the TandEM strawman payload. I hâve studied the technical difflculties of robotic emplacement of this very sensitive sensor System on Titan's surface and set the surface constrains for its installation. The idéal position of a seismometer is the one which will guarantee the accurate recording of the ground vibration above the instrument's threshold. Moreover, the instrument's legs should contact firmly with the local surface at the landing site. The presence of dust and sand grains can influence significantly the instrument's performance. The location should be on a fiat surface away from dunes and drainage networks. This study led to the proposai of implementing seismometers in future missions to icy moons, which is described later in the Chapter.

Figure 7 . 6 -

 76 Figure 7. 6 -TandEM mission concept (Coustenis et al., 2009a). Design by Dr. M. Hirtzig.

  were the two finalists selected by ESA in October 2007 for further study in the framework of the 2015-2025 Cosmic Vision Plan as L-class missions in collaboration with NASA. Both studies were accepted as fully responsive to the Cosmic Vision Goals.

  Mission (TSSM) is a Flagship (large) mission which was to be developed jointly byNASA and ESA (TSSM Final Report JPL D-48148 NASA Task Order NM0710851). TSSM could study for several years Titan's atmospheric and surface environment and explore extensively one of Titan's lakes. TSSM is planned to be launched after 2025 and after a 9-year interplanetary voyage it would reach the Satumian System. The duration of the nominal mission will be four years. I hâve been given the unique opportunity to participate in the design of a future mission dedicated to Titan, the subject of my PhD Thesis. I got involved in the TSSM mission study ffom its very beginning. During the 15-month préparation for the sélection procedure, I supported Dr. A. Coustenis, one of my Thesis supervisors and European lead of the TSSM Consortium, in administration issues. One of my administration duties was, and still is, the web support of the mission. I am the web master of the TSSM official web site and I hâve also created the TSSM main article in Wikipedia, the web sites of which are the following: • http://www. lesia. obspm.fr/cosmicvision/tssm/tssm-public/ • http://en. wikipedia. org/wiki/Titan_Saturn_System_Mission I hâve also, more essentially, been involved in the définition of Cassini-like instruments on board TSSM, which would study the atmosphère and the surface like the infrared spectrometer which would follow in the steps of the Voyager/IRIS and Cassini/CIRS instruments (see Chapter 4), as well as the spectro-imager similar to VIMS with improved capabilities from my expérience described in Chapter 5. The corresponding instruments on

Figure 7 . 8 -

 78 Figure 7. 8 -The TSSM basic concept: the orbiter, the balloon and the lake lander (TSSM Final Report. JPL D-48148 NASA Task Order NM0710851).

Figure 7 . 9 -

 79 Figure 7. 9 -The TSSM orbiter (TSSM Final Report. JPL D-48148 NASA Task Order NM0710851).

Table 1 . 2 -

 12 The orbiter's payload (TSSM Final Report. JPL D-48148 NASA Task Order NM0710851) fluxgate sensors. Noise level -1 lpTm,s. Interaction of field with ionosphère: internai and induced field. TOF Analyzer w/ss detectors to measure magnetospheric particle fluxes, -10 keV to >MeV with 150°x 15°FOV. of spacecraft telecom System. Lower stratosphère and troposphère T profile. Gravity field.

Figure 7 . 10 -Figure 7 . 11 -

 710711 Figure 7. 10 -The TSSM Montgolfière gondola and a view of the entry System aeroshell {TSSMFinal Report. JPL D-48148 NASA Task Order NM0710851).

Figure 7 . 12 -

 712 Figure 7. 12 -The TSSM lake lander (TSSMFinal Report. JPL D-48148 NASA Task Order NM0710851).

Figure 7 . 13 -

 713 Figure 7. 13 -The positions of the lake lander payload {TSSM Final Report. JPL D-48148 NASA Task Order NMO710851).

Figure 7 . 14 -

 714 Figure 7. 14 -The TSSM Lake Lander {TSSMFinal Report. JPL D-48148 NASA Task Order NM0710851).

Table 7 . 4 -

 74 The Lake Lânder Payload (TSSM Final Report. JPL D-48148 NASA Task Order NMO710851). spacecraft/lander telecom SystemThe TSSM competitor for the final sélection as the next L-scale mission was the Europa Jupiter System Mission (EJSM). In February 2009, a down-selection meeting was conducted in Washington DC by the NASA and ESA management and the Europa Jupiter System Mission was selected as the most technically feasible to pursue first. ESA's Solar System Working Group recommended, and NASA agreed, that TSSM would be the EJSM follow up mission within the next decade. However, as I hâve mentioned above, the EJSM project has been terminated and ESA carries forward alone with JUICE mission to the Jupiter System. The future exploration of the Satumian moons, and especially of Titan and Enceladus, has been supported by several white papers which I hâve co-authored with the Titan Working Group (TWG). These white papers were submitted in 2009 in response to National Research Council (NRC) Planetary Science Decadal Survey 2013-202234. In the Lunine et al. white paper, we highlighted the significance and the scientific merit of a future mission to Titan. In the white paper of Spilker et al. we recommended the continuation of Cassini mission operation by exceeding its fonction for 7 more years, the Cassini Solstice Mission, which NASA accepted. In Nixon et al. article we recommended Titan exploration to be as a high priority target, the continuation of fonding for large ground and space based télescopes in order to perform long-term observations of Titan and to support laboratory experiments that provide constraints in Titan science. Titan was identified in that paper as the only crédible analogue to the Earth's seasonal-dependent and climatic variations and as such an important study object. Seasonal variations of Titan's atmospheric chemistry are an important part of my PhD studies (see Chapter 4). The laboratory support as well as the launch of an orbiter to sample in situ Titan's upper atmosphère with high resolution instrumentation was the subject of the Yelle's et al. white paper. In the Allen et al. white paper, we recommended a mission concept similar to TSSM to be delivered at Titan, considering its 34 http://sites.nationalacademies.org/SSB/CurrentProjects/ssb_052412 high astrobiological interest. In the Coustenis et al. white paper, we also recommended a balloon mission to Titan as high priority for decadal science and to invest in reducing the risks of such a flight. In the Atkinson et al. white paper, we recommended that probes be sent to giant planets beyond Jupiter. The NRC Committee recommended in March 2011 the continuation of the study of TSSM and considering the mission's flexibility, NASA anticipâtes future studies to divide this large flagship mission into separate flight opportunités in varying sizes (see NASA responseto Planetary Science Decadal Survey Report(9)(10)(11)(12)(13)(14)(15)(16)(17)(18)(19)(20)(21)(22), 26/7/2011).7.2.4 Titan Aerial Explorer (TAE)Besides TSSM other concepts for future missions in medium scale sizes to Titan hâve been proposed such as Aerial Vehicle for In-Situ and Airbome Titan Reconnaissance (AVIATR), Titan Mare Explorer (TiME), Titan Lake Probe and Titan Aerial Explorer (TAE).) is a proposai suggestive of alternative concept to the Titan balloon mission. Since Titan expériences low gravity and a dense atmosphère, a nuclear powered airplane could fly more than 20 times easier than on Earth. It could sample directly the atmosphère and cover huge swaths of Titan's landscape[START_REF] Bames | AVIATR-Aerial Vehicle for In-situ and Airbome Titan Reconnaissance. A Titan airplane mission concept[END_REF]. Although very interesting, this proposai did not get selected among the Discovery mission studies.

Figure 7 . 15 -

 715 Figure 7. 15 -AVIATR air vehicle exterior blueprint[START_REF] Bames | AVIATR-Aerial Vehicle for In-situ and Airbome Titan Reconnaissance. A Titan airplane mission concept[END_REF].

Figure 7 . 16 -

 716 Figure 7. 16 -View of the AVIATR vehicle interior from below. The locations of the scientific instrumentation is showed[START_REF] Bames | AVIATR-Aerial Vehicle for In-situ and Airbome Titan Reconnaissance. A Titan airplane mission concept[END_REF].

Figure 7 . 17 -

 717 Figure 7. 17 -Artist's view of the TiME sailing Leigia Mare on Titan (Stofan et al., 2010a; b).

Figure 7 . 18 -

 718 Figure 7. 18 -Titan Lake Probe (crédit: Waite et al. 2010, Decadal mission concept study)

Figure 7 . 19 -

 719 Figure 7. 19 -The gondola of TAE (crédits: Hall et al., 2011)

Figure 7 . 20 -

 720 Figure7. 20 -The entry, descent and inflation sequence of TAE (crédits:[START_REF] Hall | Titan Aerial Explorer (TAE): Exploring Titan By Balloon[END_REF]

Figure 7 . 21 -

 721 Figure 7. 21 -Terrestrial Stromatolites (http://www-eaps.mit.edu/geobiology/biomarkers/whatis.html photo: R.V. Bume).

  low power consumption, high-modulation depth, high emissivity and a long lifetime. In general MEMS devices combine: 1) Very small size (20 pm to 1 mm) 2) Fast pulsing (thanks to low mass of the emitter) 3) Limitation of low output power (450 mW) 4) Réduction of the thermal mass of the System 5) Enhancement of the modulated performance 6) Extremely low cost Design considérations

Figure 7 . 22 -

 722 Figure 7. 22 -Deposited filament MEMS emitter (source: Brian Elias, Photonics Spectra (Elias, 2008))

Figure 7 .

 7 Figure 7. 23 -A pattern of the MEMS 5-inces package (source: NASA, Nano ChemSensor Unit (NCSU) Experiment http://128.102.216.35/factsheets/view.php7icHl 18)

Figure 7 .

 7 Figure 7. 24 -Phase One of the MEMS experiment during the EDL procedure.

Figure 7 . 25 -

 725 Figure 7. 25 -Phase Two of the MEMS experiment.

1 .

 1 Unveil the internai structure of the moon 2. Detect tectonic activity3. Confirm or not the existence of internai océanThe icy satellites of the outer planets can significantly contribute to the study of geological processes across the Solar System. Comparative Planetology can enlighten fondamental mysteries conceming the origin and the évolution of Earth and the Solar System in general and predict future geological events. However, the internai geological structure of each planet or moon can only be determined by in situ measurements.The local or global tectonic field, meteoroid impacts and the moon's tidal deformations induced by Jupiter or Satum's gravity field, as well as température and pressure fluctuations, may originate ground vibrations within the icy moons. These ground vibrations provide information about the nature of the subsurface material, its fracture and its Chemical composition. The Cassini-Huygens mission héritage encourages scientists to discuss about launching future missions towards the outer planets' satellites. In this section, I examine the possibility of installing seismic equipment on icy satellites of the outer planets, setting the minimum technical requirements and describing scientific achievements and possible problems. Data obtained by a seismic experiment, in combination with the analysis results from remote observations, will provide us with details about the interiors of the moons and their connection with the observed surface features(Bampasidis et al., 201 la).7.4.2The payload of the experiment A seismograph is the basic instrument for measuring the ground vibration. It mainly consists of the seismometer and the unit that records the signal. The seismometer consists of three sensors placed in the same sealed case. Each sensor is a pendulum that when it is triggered by the ground movement, it moves from its equilibrium position. This movement is the record of the instrument. The sensors can measure any ground motion within a frequency range usually of 0.001 Hz to 100 Hz, at the north/south, east/west and vertical component in orthogonal System.

  s instruments on the icy satellites' surface of the outer planets will meet abnormal physical conditions during the entry descent and after landing of their carrier. To achieve the seismograph's accuracy in a spécifie location the environmental factors (mainly température and pressure) should be considered. The ground vibration is the input signal to the System. If the ground vibrâtes within a range of 0.001 Hz to 100 Hz, the motion signal will be recorded by each one of the three sensors. The sensor System contains a transducer, the device that converts the mechanical motion into electrical signal. A piezoelectric

Figure 7 . 26 -

 726 Figure 7. 26 -Functional procedure of the seismic instrument.

Figure 7 .

 7 Figure 7.26 are parts of the same physical equipment, the seismograph.

0 180Figure 7 . 27 -

 0727 Figure 7. 27 -Seismic waves propagation through the Earth. The Earth is a differentiate medium. The two types of body waves, initiated by earthquakes or by explosions, travel through the interior and provide crucial information about the internai structure. The primary waves (P) compress the material and can travel either through solid or fluid media. On the other hand, the secondary waves (S) shear the medium and can only travel trough solids. Hence, the recordings of P and S waves on the surface can give us dues about the nature of the interior of a planetary body (crédit: JL Ahem, 2009).

Figure 7 . 28 -Figure 7 . 29 -

 728729 Figure 7. 28 -Titan network mission scénario. Detailed long-term surface studies require data collection though out a geophysical network[START_REF] Lange | Concept Study for a Titan Geophysical Network[END_REF]. Seismic equipment can be included as payload of such a mission.As in other planetary bodies except for the Earth (Moon, Mars, Venus) Satum's icy moons seem idéal targets for seismic experiments since they possibly possess active interior.

  discouraged to get involved in Astronomy by the complexity of modem technology. Despite the intense interest in astronomical phenomena and the contemporary facilities like Internet or a great amount of exceptional publications and TV broadcasts, most people hâve not got familiar with astronomical phenomena yet. The rôle of Space Scientists is vital in reducing the distance between Astronomy and the public. The Committee on the Planetary Science Decadal Survey Steering group of the National Research Council of USA strongly recommended and NASA administration agreed to spend at least 1% of the cost of each space TA mission for éducation and public outreach activities .

8 . 2

 82 The educational value of Titan and the Cassini-Huygens mission I hâve been studying the Cassini-Huygens discoveries since 2006. My participation and contribution to the Kronos, TandEM and later TSSM and TAE mission studies gave me the opportunity to interact with many scientists worldwide, experts in Planetary Science. The number of these exchanges increased during my PhD Thesis studies and had significantly benefited me and advanced my knowledge in Titan science as well in Physics and Mathematics. In the meantime, I hâve been thinking of the possible benefit for the broader public from similar interactions.

Figure 8 . 1 -

 81 Figure 8. 1 -Annual report of activities listed in Table8.1 and the number of participants

- 18 years

 18 old) to describe where they think the caméras on Cassini should be pointed during upcoming observations scheduled for every year's fall. A science team of expertise designs the scientific plan for each year by setting three possible targets within the Satumian System, including Satum itself. Each of the three targets has a scientific importance that justifies a possible future observation by Cassini. The task is to choose a target focused on the best science retum and explain the reason, in an essay of about 500 words. They also hâve the option to do team work through groups of maximum three students. For the years 2010 and 2011 (2012 contest is currently under planning), the Space Physics Group of the University of Athens in association with extemal colleagues has been selected as the coordinator of NASA for the international compétition in Greece. This kind of school compétition in Greece is particularly important since students rarely hâve the opportunity to expérience and participate actively in similar tasks. Under the guidance of the Cassini Outreach team, A. Solomonidou, Prof. X. Moussas and I hâve informed, explained and spread the rules of the compétition at primary, secondary and high schools ail over Greece. We kept open communication with students, teachers and parents for questions and guidance. For the 2010 contest, the target number one was a close-up view of the moon Rhea, number two was a video of Titan, Tethys and Enceladus dancing under Satum's rings and target three captured a day on Satum. The main preference of about 60% was target number two. For the 2011 contest, the target number one was Hyperion, number two Rhea & Titan and number three Satum. Again most arguments were in favor of target number two (67%), showing a preference trend on Titan.The participation in the contest for 2010 (first time in Greece) was unexpectedly high and thoroughly satisfied at about 150 original essays, while for 2011 the participation increased up to 75%. The winners awarded through a ceremony, which was held in the old amphithéâtre of the University of Athens, which was fully packed (Fig.8.1).

Figure 8 . 2 -

 82 Figure 8. 2 -The winner's ceremony of the NASA contest 'Cassini Scientist for a day' in Greece, May 2012

  objects to bring onboard in case of sélection as astronaut d)Mention a song which fits a future space mission

  ability to solve scientific problems in the future. With such activities, students hâve developed problem solving and their presentational and communicational skills. Moreover, this kind ofgroup work activities introduces students in the Science discipline. During my Thesis study, I hâve investigated different aspects of Titan's environment, focusing mainly on its middle atmosphère but with some non-negligible contribution to its surface science as well. The Thesis manuscript can be divided into two major parts, one conceming the techniques of analyzing Cassini/CIRS data and the other purely scientific. The outcome of the scientific research is then used as a basis for presenting Titan's habitability potential and experiments' design for future missions. During my research, while covering the Cassini mission, I hâve managed to infer variations as a fonction of latitude and time. I hâve confîrmed that a rapid change in the atmosphère took place at 50°N (Bampasidis et al., 2012a). I was led to this conclusion after finding an indication for a compositional enhancement from 2006 to mid-2009 during Titan's northem spring equinox (NSE) on 15 August 2009 for almost ail molécules compared to the 2008 and 2010 values. These results are compatible with the fmdings of Teanby et al. (2010).

  As far as the surface is concemed, I hâve searched for links of the atmosphère with the surface and the interior, since my Thesis subject was on the environment of Titan. Titan's surface liquids (lakes and their surroundings) can provide one such linking point. The surface liquids reveal information conceming the évolution and the dynamics of the "hydrological" processes on Titan surface. The corrélation between the geomorphic features that are expressed by the Hortonian morphometric parameters, on which I hâve worked, and the local topography can provide a qualitative description of the region's development. The analysis of the Mayda Insula drainage shows the existence of steep slopes on its Western part. Since the island basins hâve various slopes, I assumed that they hâve experienced various élévations.Moreover, I hâve contributed to the research of possible morphotectonic expressions on Titan through a comparative study between Cassini observations and terrestrial tectonic Systems. From the above, we managed to make suggestions about possible formation mechanisms of the observed features(Solomonidou et al., 2012). Furthermore, inBratsolis et al. (2012), I got involved in the application of a despeckle filtering technique for obtaining restored Cassini/SAR images. This tool is based on probabilistic methods and in combination with a segmentation technique, makes possible the extraction of régions of interest from the local background.

( 1923 )

 1923 any modeling of physical phenomena should satisfy the next three attributes: a) the existence of a solution, b) the uniqueness of the solution, and c) the stability of the solution (continuous dependence) on the observed data.

  vertical température and aérosol profile of Titan's atmosphère for both limb and nadir measurements. The column vector x is a [(2nq+l) x 1] vector.On the other hand, we define a forward spectral radiance model vector h{x) with melements such as the measurements Ik. Then, by calculating the partial dérivatives of the model h in respect to the n-elements vectors Tj and bj we dérivé: x n) Jacobian matrices. Additionally, the partial dérivatives of the model radiances h in respect to the p-elements shifts at each tangent point latitude are: matrices (Lyand M.) and the vector W/Construct the (m x (2nq+p)) array K:

  fonction39 in respect to a : Q = (AI-KFaf E~x (AI -KFa) + yara With E we represent the measurement error covariance matrix. The first term illustrâtes the usual penalty fonction of the least square fittings, while the second one is a measure of the departure of the solution from the reference profile. Since we want to minimize the previous fonction Q in respect to a, we should apply the following matrix identity: 3^+ bfc(px + + E)rCÏA + + b)rCD dx Then, the partial dérivative oïQ 'm respect tocrbecomes: -= (AI-KFaf (E~l J (-KF) + (AI-KFaf (e~1 f-KF) + 2yar and ffom the définition of the covariance matrix (^-1) = E~\ the dérivative takes the form: ^2 = 2(AI -KFaf (E-lf-KF) + 2yaT By setting -= 0, we quest for the minimization: da (AI -KFaf E~lKF = yaT Now, by applying the identity (ABCf = CTBTAT and by replacing the KF = G we dérivé:39 By the superscript T we define the transpose matrix.

(

  aIT -arGT)E-lG = Yar AITE~lG -aTGrE~lG = yaT MTE~lG = aT (yl + Gt E-lG) AIT = aT(yI + GTE-'G\E-lGy UT =aT(yI + GTE-ïG)G-lE Mt = ar{yG~'E + Gr£"1GG"1E) AIT =aT(yG~lE + GT) where I is the unit matrix, and by transposing the last expression we hâve: (a -[ar(yG l£ + Gr)J which leads to: AI = (yG-'E + GTja AI = ((yG-'E'f + G^! a/=(/e(g'')t + g\ (/E(G-'J + GJ Now its time to seek for an expression containing the température différence Arhence we multiply with the F matrix and replace again the G = KF : F^Ep'j + g)al = Fa ((££)"'J +Kf) AI = Fa f{yE((KFf J' + KFl) AI = Fa AT = f[{yE + (KFjKFj ^KF'f )'j AI Therefore, the minimization of Ogives the following expression:

  bands were replaced by Arie and Johns results[START_REF] Arie | The bending energy levels of C4H2[END_REF], while the rest replaced by the GEISA 2003 ones .The 2006 version of the code could perform calculations for 14 individual molécules: methane (CH4), (2Hi) methane (CH3D), carbon monoxide (CO), carbon dioxide (CO2), hydrogen cyanide (HCN), cyanoacetylene (HC3N), cyanogen (C2N2), methylacetylene (C3H4), acetylene (C2H2), ethane (C2H6), ethylene (C2H4), propane (C3H8), diacetylene (C4H2) and water (H2O).

  (CH3CN) vg, 2vg, V4, V7 bands were added to the code from Pacific Northwest National Lab -PNNL data (https://secure2.pnl.gov/nsd/nsd.nsf/Welcome)

  C2H6 bands were replaced by the ones provided from VanderAuwera et al. (2007). Following its détection(Coustenis et al., 2008), C2HD V4 and V5 bands were updated provided by the work of Jolly et al.[START_REF] Jolly | Measured integrated band intensities and simulated line-by-line spectra for 12C2HD between 25 and 2.5 pm, and new global vibration-rotation parameters for the bending vibrations[END_REF]. Moreover, three molécules are separated from their isotopologues: HC14N and HC15N adapted fromHITRAN 2004, 12CH4 and 13CH4, 12C2H2 and 13C12CH2 adapted from GEISA 2003. Eventually, this version takes into account for calculations 21 molécules and is named asARTT 0.3.8.

•

  Linux or Macintosh OS distribution The source distribution of ARTT (ARTT-src-*.*.*-yyyy-mm-dd.tar.gz) contains ail the source files for the installation of the program. In the name of the file we can find its version in the part of the name seen as **' and the date that its development ended in the 'yyyy-mm-dd' part of the name. The current version is 1.0.0. The source folder contains the following folders: • ARTT-work-*.*.*: This folder contains some sample files to be used for the modeling of the spectrum after the installation. • cirs: This folder contains Cassini CIRS datasets • db: Here we can find ail the spectroscopic data files which are used for the calculation of the model.

A:

  Wavenumber of the line B: Intensity of the line at 296K C: Air broadening pressure half-width D: Energy of the lower transition level F: Température dependence coefficient n of the air broadening half-width G: Identification code for isotope as in GEISA I: Identification code for molécule as in GEISA Adopt HITRAN linelistsThe procedure needed to get a molécule or an isotope from the HITRAN database is somewhat different from GEISA. HITRAN is also a web accessible database and it offers a program called javaHAWKS. After downloading and installing javaHAWKS locally we proceed to the "HITRAN2008/By-Molecule/Uncompressed-files" folder. Every file in this folder corresponds to a spécifie molécule. In order to track down this molécule we must advice the corresponding HITRAN database table. After downloading the data file we open the JavaHAWKS program and complété the following steps:Click on "Select" tab -In the field "HITRAN filename" enter the name of the data file -Give a name in the output file -Edit the sélection parameters and hit the "Run select" button Now the requested file is created locally and it contains the data specified in the "SélectionParameters" area.The last step remaining is to complété the transformation of HITRAN data to ARTT format. The problem that we meet after downloading the data from HITRAN is that the wavenumber values should be in the first column and the molécule and isotope id numbers should be in the last 2 columns. The following program written in Fortran makes

7 ) 8 )

 78 )-db/geisa_C3H8.dat' forma[START_REF] Waite | Ion neutral mass spectrometer results from the first flyby of Titan[END_REF]='(fl2.6,lx,dl 1.4,f6.4,fl0.4,36x,f4.2,i3,i3)' These lines describe the file we hâve added in the /db folder. The number inside the parenthesis is the unique identifier of the file under the fichier column in the spectroFP*.don files. In the first line we add the exact path to the new file and in the second line we edit the format in Fortran language. The argument fl2.6 means that we hâve a number with twelve digits, six of which are the décimal fraction. The argument 36x means that we hâve 36 blank spaces. The argument dl 1.4 means that we hâve a number in exponential form with four décimal digits. In the file compounds.inc we must augment the nbcomp parameter to the number of molécules that our radiative transfer database contains. It is always the same number we add in the file trans_rad.fi Then we must add the right number in the end of every set of variables. When we add an isotope of an already existing molécule ail the values we must add except the molecular weight are the same as the ones of the original molécule. It is crucial that the name of the molécule will be only 6 characters long.

Figure 1 .

 1 Figure 1. Example of fit of the FPÎ (a) and FP4 (b) spectrum of Titan taken by Cassini CIRS in June2009 al flybys T56-T57 al 50°N, containing 476 spectra for FP3 and 575 spectra for FP4. (A color version of this figure is available in the online journal.! 2. DATA ANALYSIS

  this new release (ARTT, for Atmospheric Radiative Transfer for Titan), we hâve included more constituents relevant to Titan's chemistry (26 species in total, ail the so-far detected hydrocarbons, nitriles, and oxygen compounds in addition to the main molécules Na. H2, CH4 and argon). Our list. of molécules includes in particular tlie i3CH4,13CH3D, 12C13CH2, and H15CN isotopologues which are essential in fitting the émission bands of methane, acetylene, and hydrogen cyanide and allow us to infer the 12C/I3C and 14N/15N ratios. Furthermore, new oxygen and hydrocarbon isotopologues as detected by Jennings et al. (2008) and Nixon et al. (2008b) are included. Spectroscopic parameters for the molécules found in the FP3 and FP4 spectral ranges are from GEIS A 2009 (Jacquinet-ITusson et al. 2011) and H1TRAN 2008 (Rothman et al. 2009) databases, with several updates, as for instance in the case of ethane and propane, both of which are simulated from the individual line lists of Valider Auwera et al. (2007) and Flaud et al. (2010), respectively.

Figures 2 (

 2 Figures 2(a)-(d) show the results of the température as a function of latitude, using the altitude information contained in the resolved Unes of the t>4 methane band. Figures 2(a) and (b) (upper panel) show the température structure inferred at two latitudes (50CN and 30°N) at certain dates during the Cassini mission. The température profiles discussed hereafter were calculated using as an initial guess the one from Flasar étal. (2005).But we hâve also run tests using thermal profiles retrieved from the same CIRS FP4 sélections but by adopting into the

Figure 2 .

 2 Figure 2. Retrieved thermal profiles from CIR S nadir data at (a) 50°N, (b) 30°N, (c) equator, and (d) 30°S. The typical error bar unceitainty is about 0.7 K at 1 mbar and al most 4 K at 5 mbar. Dashcd line» indicate the altitude levels above which the température lias Inglier unccitainties. (A color version of this figure is available in the online journal.)

1 Bamtasidis

 1 (a)) and about 7 K at 30CN (Figure 2(b)) at around 250-300 km in altitude, since the beginning of the Cassini mission. In the mid-stratosphere below 200 km, the northem polar région initially began cooling from 2007 up to 2009, while in 2010 a slight warming is observed again at both 50°N and 30°N.The cooling with time in the lower mesosphere (by 8 K maximum) is also witnessed at the équatorial and Southern latitudes. On the contrary, in the stratosphère (lower atmospheric levels), no significant changes are found in the équatorial températures (Figure2(c)) and this holds for latitudes down to'ltoAsTROFHYSiCAL Journal, 760:144 (8pp), 2012 December

Figure 3 .

 3 Figure 3. Time-latitude composition variations for the major trace gases of Tïtan's stratosphère: (a) C2H2, (b) HCN, (c) C4H2, (d) C3H4. (e) C2H6, (f) C2H4, (g) CjHg and (h) COa. The latitudes investigated from 2006 to 2012 are: 50°S (violet), 30°S (light blue), equator (red), 30°N (gteen), and 50°N (blue). Connected filled circles are high résolution observations (0.5 cm-'), while open circles are medium résolution data (2.5cm"1) for 2008, 2009, and 2010 (and sometimes coincide witlt tire tiigtrer resolution values). Tire 3a eslirnated error bars are indicated (sec tire text). (A color version of this figure is available in the online journal.) 30°S. More to the south (50°S, Figure 2(d)), some température variations are observed in the stratosphère, typically 2-4 K.

  profiles described above injected in our radiative transfer code, we derived the abundances of the molécules from the best fits obtained of their émission bands in tlie Titan CI RS FP3 spectra (Figure l(b)). We show the Cassini/CIRS rnixing ratio inferences for Titan's gases as found above the condensation level in Figures 3(a)-(h) with the associated errors.

Figure 3 )

 3 Figure3). We detect a maximum around the time of the NSE(Northern Spring Equinox, 2009 August 15) in mixing ratios with increases in abundances by about 30%-40% for C3II4 and GjHé, 50%-70% for C2H2 and C4H2 and about a factor of 2 or more for HCN and C2H4 (albeit with higher uncertainties for tire latter) relative to the adjacent time periods in 2008 and 2010.

  latitudes are radier uncertain. The results on the increase and the reported values near April 2009 presented for C2H2 (7 x 10~6) and HCN (2 x 10~6) in Üiese figures (just before die maximum in lune 2009) are in excellent agreement wiüi Üie findings of Teanby et al. (2010), as shown in their Figure 4, panels (c) and (d). Although there is a liint in their Figure 4, Teanby et al. do not yet see the increase in abundance that we report here for GtH2 but we are still compatible wiüiin error bars at values near 2 x 10~8 in 2009.The observed increase in abundance for sorne molécules is followed by a strong decrease which reduces ccnsiderably the observed enhancement by 2010. This decrease settles in quickly, within 1-2 terrestrial years and brings Üie abundances back to their levels prior to the ascenL Admittedly we hâve so far only 011e high-resolution nadir sélection after die NSE and one in lower resolution confirming this resuit, so that the finding probably requires further vérification. However, boüi liigh-and low-resolution nadir data taken in the 2010-2012 period seem to support the increase and follow-up decrease.

  consistent with the values given by Vinatier et al. (2007) and about 4.8 times lower than the terrestrial value (272; Fegley 1995).

  5. DISCUSSION OF THE RESULTS AND POSSIBLE INTERPRETAI'IONSIn this work, we estimate the abundances of the trace gases in Titan's stratosphère from 50°S to 50°N. We find no significant temporal variations at mid-and southem latitudes during the Cassini mission. We monitor and quantify tire compositional enhancement at 50°N, and find indication for a maximum at the time of the Titan northern spring equinox (NSE, mid-2009), followed by a sharp decrease of the gaseous Chemical content within the next Earth year or so. Our results are compatible with tire findings of Teanby et al. (2010) in tirât we find HCN and C2H2 to display a rapid increase in northern latitudes up to mid-2009 while the abundances at équatorial and southem latitudes remain stable. C02 présents no latitudinal variations anywhere because of its long photochemical lifetime. The peak in abundance is observed around the northern spring equinox, during which we know a rapid change in the atmosphère took place. Indeed, short-term var iations observed during the Cassini mission can arise from changes in the circulation around the time of the equinox. The collapse of the detached aérosol layer

  photochemical species. Changes in the solar output during the 11 year cycle can potentially affect the Chemical production rates in Titan's atmosphère. On the other hand, for the time period of the Cassini mission, the Sun has been remarkably stable going through an extended minimum with the first weak signs of increased output observed toward the end of 2009. The Chemical lifetimes in Titan's stratosphère (at 200 km) range between ~1 year for C2H4 and C3H4 up to ~20 years for HCN, which are longer than the timescales of some of the rapid changes observed. Thus, the temporal variability observed during the Cassini mission is more likely related to changes in the atmospheric circulation patterns due to progression of seasons.

  we analyzed and discussed a set of nadir CIRS data from the first year of Cassini obser-* Corresponding auchor. Fax: *33 1 45 07 74 26. E-mail address: athena.coustenis®obspm.fr (A Coustenis). 0019-1035/S -see front matter © 2009 Elsevier Inc. Ail rights reserved. doi : 10.1016/j.icarus.2009.11.027 vations (Tours TB-T10). The abundances of ail known gaseous constituées were retrieved at different locations of Titan's disk.

  Fig.1. Spectral averages performed over TB-T44 Titan fiybys as described in Table1for different latitudes binned in 10°: (a) at medium resolution (2.5 cm '); (b) at high résolution (0.5 cm 1 ).

Fig. 3 .

 3 Fig. 3. Contribution functions of methane at the center and wings and of CHjD, calculated at high resolution for équatorial latitudes and, in addition, at 70°N, for methane.

4. 3 .

 3 Fig. S. Contribution functions of some molécules at different latitudes for high résolution: (a and b) are at 5"Sand (cand d) areat 50"N. For latitudes near theequatoror the South, the contribution functions are similar to the ones shown here for 5"S.

Fig. 7 .

 7 Fig. 7. Fits in the 610-690 cm ' région of the 5"STitan spectrum. uslng constant-wirh-heighc profiles, ar medium (a) and high resolution (b), showing the QK;, Qilj. HC3N, COj, C«Ha and C2HD bands, as well as the residual radiances after fltting this région at 50"S, 15°S and 50"N. The 3 -a error bars (see text for details) on the radiance are indicated.

  withUMB 50*N distribution (x0.82

  withGCM 50 *S distribution (x0.50) [_ Doto-Model with UMB S0*S distribution (x 1.115) Doto-Model with constont C2*i (2.35e-0S) J y-'nfYfDfW'V ^|yV v y-y-'/V V' *"*V V'/' Wavenumber (cm'1) F**-9-The fit produced at high résolution for (a) 50°N, (b) 15'S and (c) 50"S as well as the residuals co the fits found when the CCM or a constant-with-height or the limbinferred vertical mixing ratios are used for CîHj. The constant vertical profile is forced to fit the cerner of the CjH2 Q-branch. The misfit near 730 cm ' is due to the fact that the very close right wing of the band éludés fitting on ail occasions, perhaps due to a noise feature in that région (see text). 3 -c error bars are indicated.

  in the North (contrary to inferences from the limited North-pole dataset we had used in Coustenis et al. (2007)), while C2Hb and C02 remain constant (Fig. lia). When relative errors are considered (toosmall to appear in Fig. lia and b), there is a possible indi cation for a notch (small unexpected decrease) near 50°N in the abundances of C2H2 and HCN and near 50°S for C2FL, and C3H8. For the 50°N possible minimum in the abundance of the aforeinentioned species, Vinatier et al. (2009) report such an occurrence at 0.01 mbar (around 350 km) for C2H2 and HCN in results they prés ent from CIRS limb data analysis. Teanby et al. (2008a,b) also in ferred such abundance minima which they interpret as a combined elfect of poleward advection of air from Southern lati tudes, a mixing barrier at around GO^N and eddy transport towards

5 and 2 .

 52 Fig. 1t. (a and b) Méridional variations of trace gascs on Titan. Error bars represent 3 -<r standard déviations including ail error sources (relative error bars, without systematic crrors) are too small to appear on this figure for most molécules. Only upper limite were inferred for QHs at latitude <40"S

Fig. 12b .

 12b Fig. 12b. Same as Fig. 12a but for lesser hydrocarbons. The model gives a good représentation of the CJRS findings from nadir data within uncertainties and tends to agréé with limb-inferred vertical profiles.

  Fig. 12c. Sa me as Fig. 12a but for nitrites and oxygen compounds within the CIRS range. With the exception of a marginal fit for HCjN, the model nicely reproduces the CIRS findings and the ISO-retrieved water vapor value. With respect to other vertical distributions retrieved from CIRS limb data, the model overestimates the HCjN abundance.

on

  Herschel measurements of the Enceladus torus combined with modeling of the fate of the species within the torus, Hartogh et al. (2011) showed that tire flux of O/O* into Titan is consistent with an Enceladus source for the oxygen seen in Titan CO, except for the fact that Enceladus does not seem to provide enough OH/H20.In this paper, we analyze the spectra acquired by CIRS in the far infrared spectral région in order to retrieve the water vapor vertical or spatial distribution in Titan's atmosphère. CIRS has been acquiring spectra ofTitan since the beginning of the Cassini prime mission (July 2004). After 2 years of the extended mission (XM), which included the 2009 equinox, in July 2010 Cassini entered in the Solstice Mission (SM), which is scheduled to last until 2017. Since an upper limit for H20 was reported by deKok et al. (2007a) tliere has been a

  observations hâve the FP1 Focal Plane centered around two different altitudes -hereafter limb 1 and 2 -and are therefore used to constrain the water vapor abundance in the stratosphère around 115 and 230 km respectively, well above the 45 km tropopause. Since the contribution functions of water for on-disk obser vations peak around 97 km (Fig.2), the retrieved water vapor abundance derived from these measurements can be compared with the lowest altitude targeted by our limb intégrations around 115 km.For the water détection and retrieval of quantitative informa tion together with possible latitudinal variations, multiple Titan flybys must be utilized to enhance the signal. To date, 35 limb inté grations of approximately 1 h in duration (~60 high-resolution spectra) hâve been obtained covering latitudes from 87°S to 8(hN. The nadir intégrations are more ntimerous (about 92 successfully executed, of typical duration 5 h, --300 spectra) as they occur in a less contested observing time further from the Titan closest-approach period. They also hâve a more or less complété spatial coverage of Titan's latitudes and longitudes with an average footprint size of ~15°great circle arc We focus on one season of on-disk observations acquired from December 2004 to December 2008 (northern winter on Titan) in order to reach a compromise between obtaining a large number of spectra and a sufficiently homogeneous dataset. Inside this time period for on-disk observations acquired from a maximum dis tance of 300,000 km and with a maximum émission angle of 60°, we selected latitudinal bins (80-45°S, 45-10°S, and 0-30°N) cen tered around three latitudes for which observation-derived tem pérature profiles were available (see model description in Section 3 and Fig.3). The number of spectra averaged in these latitudinal bins were approximately 1700, 3800 and 7000 respectively and their average émission angles were 35°, 38°and 34°.CIRS limb spectra are acquired in much smaller numbers, there fore to reach a sufficient signal-to-noise ratio we consider only one average of about 320 spectra acquired from December 2004 and September 2009, encompassing the entire Southern hémisphère and mid-Iatitudes within the range 90°S to 20°N. During this time period and at these latitudes data can be considered quite homoge neous as shown inTeanby et al. (

Fig. 1 .Fig. 2 .Fig. 3 .

 123 Fig. 1. (a) In black, the average of Cl RS far-IR on-disk observations is plotted (~7000 spectra acquired from December 2004 to December 2008 in the latitudinal range of 0-0"N), limb observations centercd around 115 and 230 km (respectively ~320 and ~280 spectra acquired from December 2004 to September 2009 in the latitudinal range of 90-20"S) and their fit (in green, blue, red respectively) assuming a constant water mole fraction above the condensation altitude. In (b) the retrieval spectral range is shown with the main water lines indicated by vertical dotted Unes. (For interprétation of the references to cobrin this figure legend. the reader is referred to the web version of this article.)

  Fig. la the on-disk and two limb observation averages are shown together with their fits. Fig. 1 b shows only the spectral range used for the water line analysis.

Fig. 3 .Fig. 4 .

 34 Fig. 3. The observed water mixing ratio indicates the absence of any significant latitudinal variations within the error bars in these

  Fig. S. Water vapor mole fraction retrieved from CIRS on disk and two limb observations assuming a water profile constant with altitude over the condensation level. Water profiles from three photochemistry models are also shown for comparison. The Hôrst et al. (2008) water vapor profile was derived assuming two different eddy diffusion coefficients -100 cm1 s 'and 400 cm2 s '(the second being the one recommended in their model; dotted curve plus solid curve). Also the profiles in Wilson and Atreya (2004; dashed curve) and in Lara et al. (1996; dotdash curve). latitude ranges. It should be stressed that in this work we did not analyze the water stratospheric content at high northem latitudes that were experiencing winter during this time period. To model spectra at these latitudes for a large average is particularly complex since the stratospheric température was changing quickly with latitude. In addition the température profiles for high north em latitudes are not yet available for the lower stratospheric ré gion sensed by the CIRS water fines.

  Titan and the water abundance profi le dérivés from the amount necessary to form the observed C02. This assumption was challenged by the Hôrst et al. (2008) model, in which oxygen species are assumed to arrive from outside the Moon and form carbon monoxide as well as carbon dioxide in the atmosphère. The values of the input fltixes of O and OH were adjusted to reproduce the observed abutidances of CO and C02. In the Hôrst et al. (2008) model water profiles were produced for six different values (from K0 = 100 cm2 s"1 to K0 = 1000 cm2 s-') oftheeddy coefficient in the lower atmosphère,

«i 2011

 2011 Published by Elsevier Ltd. î. Introduction Cassini-lluygens, cite extreinely successful NASA/ESA joint mis sion to the Saturnian System, has already accomplished six years of extended investigatioa Titan being one of the most intriguing objects in the Solar System, hosting an Earth-like environment, is one of the main targets of this mission. Titan's landscape has been observed by multiple flybys that examine the surface expressions through Cassini's infrared and radar instrumentation.

  when the sensors of this lattice select one scene from .F. Let us note the finite lattice of sites asS= (s}s ^with ScZ2. Each site se S has a set of neighbors r noted J\fs. If s=(ij) four-connectivity is assumed in the following : ,VS ={(f-lj);(i+l j);(ij-l);(ij+l». A random process {/s|seS} with fs random variables following the joint probability function P(f) is a "random field". In the spécial case where P(fs\[fr^s}) = P(fi\{fr}r,M,) the random field is 110 E. Bratsolis et aL } Planetary and Space Science 61 (2012) 108-113 ta lied a "Markov random field" (MRF). By che Hammersley-ClifFord theorem MRFs and Gibbs random fields on a finite lattice

Fig. t .

 t Fig. t. Initial image of lakes (P1A08630, NASA/JPL) after subtraction of négative <r0.

Fig. 2 .

 2 Fig. 2. Initial image of lakes (P1A08630, NASA/JPL) after subtraction of négative an and subtraction of values greater than <<r5 > +3 srdfoo).

Fig 4 .

 4 Fig 4. Image ratio between the initial (Fig.2) and the despeckled (Fig.3) image.

  images without missing important textural details from the régions of interest. Consequently, Fig.3seems smoother than Hg. 2. Hg. 4 depicts the ratio between the initial and the despeckle image. In this figure the lakes' interior is presented smoother than the surrounding area. This can be explained as the dark lakes are not really speckled régions. From the initial data (and <To in nondimensional values) the samples of the région 2 (granular lake)

112E.Fig. 5 .

 5 Fig. 5. Segmented image aller filtering.a critical parameter for the global methane cycle(Atreya et al., 2006}.

  geology is tlie fïeld of geological research that focuses on the study of features observed on the crust of the Eartli and that of other pianets investigating the processes, forces and movements that resulted in thetn. Tectonism encotnpasses geological events not caused by exogenous processes such as érosion and meteoritic impacts. Tectonism being compressional or extensional is related with important endogenous processes such as terrestrial volcanism and most probably with extraterrestrial cryo volcanism. Morphotectonics correlate landscape morphology to

  et aL, 2005), the Surface Science Package(SSP-Zarnecki et al., 2005), and the Cas Chromatograph Mass Spectrometer(CCMS-Niemann et al., 2005, 2010) provided additional information of Titan's geology. The actual landing site on the Saturnian satellite appears to be a relatively soft surface similar to tar or dry sand, tinted by methane ready to evaporate and providing ample evidence for fluvial and aeolian processes. 2.7. Surface expressions 2.1.1. Geological features formed by non-tecronic processes

  upper crust due to past high heat flux from the interior and high température gradients in the ice Shell Transitions of the ice shell over a liquid subsurface oceaa from a conductive State to a convective State, causes tectonic stresses and movements that influence the surface Localized compression due to thickening of the crust linked with the cooling of Titan at areas where fault structures stresses that penetrate the icy shell and create pathways for the internai material Recent crustal thickening due to localized extension Déposition of ejecta blocks around craters in a radial manner Erosion and incision of terrains that form régional uplifted due to Sevicr/Laramide Orogeny Tectonic and magmatic aspects on geological terrains: Mid Atlantic Ridge. Atlantic Océan Ahaggar Mountains. North-central Sahara Desert Mountains due to extension: Basin and Range Province (Harcuvar Mountains, Cila Mountains, MaricopaMountains)(Radebaugh et al., 2007) Ejecta patterns: Meteor crater, Arizona Erosional geomorphological structures: Colorado Plateau, USA[START_REF] Lorenz | Titan's surface before Cassini[END_REF] Radebaugh et al.. 2007). It is dissected by a number of long north-south trending normal faults while deep entrenchment of streams and differential érosion hâve formed high standing crustal blocks mountain ridges (in this paper we will refer to mountain ridges as ridges), cliains of elevated (uplifted) ground that extend for some distance The major mountains and ridges on Titan are listed along with their location and observational characteristics in Table

Fig. 2 .

 2 Fig. 2. Major mountainous régions on Titan, (a) Mountain extends for almost 240 km (NASA/JPL). (b) Long bright ridges with multiple mountain peaks were observed in T8 on October 28, 2005 (linear mountains extend from 13-5"S to 198-225"W) (e.g" Radebaugh et al., 2007; topes et al, 2010), and extend over 480 km (NASA/JPL) (c) Three radar bright parallel ridges (2'S, 127JW) within the mountainous area of Xanadu from T43 of May 12. 2008, the length of the image is almost 400 km (e.g,. Mitri et aL, 2010) (NASA/JPL/Space Science Institute).

Fig. 3 .

 3 Fig. 3. Rocky Mountains, USA (a) in LandsatTM scene with DEM data (Crédit: Fédération American Scientists FAS). (b)3-D perspective view by combining two spacebome radar images (PIA01840 NASA/JPL).

Fig. 4 .

 4 Fig. 4. (left) Fault-block mountains on Titan (portion of Fig. 2b)(PlAQ3566 NASA/JPL) formcd possibly through crustal extension (Radebaugh et al., 2007). (right) Relief map of the Basin and Range province in west-central Nevada (USGS) displaying the parallel ranges and valleys created by crustal thinning and fracturing by extensive stresses.

4. 1 .

 1 Tectonic contrat on Titan's linear features: Faults, fractures, canyons and drainage networks In geology, a failli is a rupture that séparâtes a rock unit into two parts, moving one relatively to another, in a microscale or a whole field in a macroscale. A variety of geological processes are associated with faults and therefore their analysis is very important For instance, geologists consider the direct relation of earthquakes with faults, or the pénétration of igneous rocks on Earth's crust along faults, or also the interaction of faults in the development of sedimentary basins. On the other hand, canyons are another ty|>e of cmstal scars, formed on Earth by the accelerated érosion by rivers entrenched aller tectonic activity and trying to reach base-line

Fig. 5 .

 5 Fig. 5. (left) Radar image showing canyon-like Systems on Titan (71*S, 240°W) (P1A12036 NASA/JPL). (right) Massive canyon formation on Earth. Grand Canyon, USA (USGS). Please cite this article as: Solomonidou, A., et al., Morphotectonic features on Titan and their possible origin. Planetary and Space Science (2012), http://dx.doi.Org/10.1016/j.pss.2012.05.003

  marks on VIMS data(Barnes et al., 2006) Circular tectonic features(Soderblom et al., 2009) Topographie élévation, mountain-like structures (unidentified)[START_REF] Bames | Near-infrared spectral mapping of Titan's mountains and channels[END_REF] Fl*. 6. Area on Tui Regio (left) with possible tectonic inlluence; two dark patches on Hotei Regio (middle) assumed to be volcanic caldera ridges (NASA/JPL/University of Arizona); Harrat Khaybar (right), massive volcanic terrain in western Saudi Arabia on Earth; the dashed Unes indicate the linear trend of the volcanic vents suggesting tectonic control (NASA). Please cite this article as: Solomonidou, A., et al., Morphotectonic features on Titan and their possible origin. Pianetary and Space Science (2012), http://dx.doLorg/10.1016/j.pss.2012.05.003 A Solomonidou et ai / Pianetary and Space Science fini) iii m zones of weakness on which volcanism and tectonism occurred.

Fig. 7 .

 7 Fig. 7. (a) Europa s major geo-structures that may hâve formed due to volcanic and tectonic processes acting together (NASA/JPL-Caltech). (b) Ganymede's grooved and tectonic terrain (NASA/JPL/Brown University). Please cite this article as: Solomonidou, A., et al., Morphotectonic features on Titan and their possible origin. Pianetary and Space Science (2012), http://dx.doi.Org/10.1016/j.pss.2012.05.003

  features resulting from tidal-induced convection hâve different âges, as can be shown from superposition and cross-cutting relationships. Mountains are old(20-100 million years-Rade baugh et al., 2007) and probable cryovolcanic areas are yotinger (less than 10,000 years for Hotei Regio-Soderblom et al., 2009). This is to be expected since tidal-induced convection will cause ice élévations with subséquent breaking of the surface ice and formation of fractures (extension). The ice floats may collide Fl*. 8. Location map or the major morphotectonic features on Titan (background map crédit: NASA/JPL)i In green: mountains, ridges and hills; in blue: linear features, faults, fractures, canyons; in red: probable cryovolcanic régions. (For interprétation of the references to color in this figure legend, the reader is referred to the web version A. Solomonidou et aL / Planetary and Space Science l (mi)

Fig. 4 a

 4 striking structural similarity exists in the morphology of the T8 mountain range on Titan and the Basin and Range Province on Earth. Alternative proposed mechanisms for mountain formation on Titan do not involve crustal movements: Radebaugh et al. (2007) argue that blocks of mountains viewed in T3 swath could hâve been formed by impact ejecta. Such deposits resemble those of Meteor Crater in Arizona, USA where blocks of ejecta deposited 1200 m away from

  originated through three distinct épisodes: the fîrst following the silicate cote formation, accretion and différentiation period; a second épisode upproximately 2000 million years ago wlien convection contmenced in tlie silicate core; and finally, a geologically recent period, circa 500 million years ago,

Fig. 2 .

 2 Fig. 2. Tiger Stripes on Enceladus (NASA).

Fig. 3 .

 3 Fig. 3. Enceladus' Plume front Cassini/ISS (Porcoétal., 2006).

Fig. 4 .

 4 Fig. 4. (a) Ridges and movmtains on Titan's surface. The radar bright featurcs arc part of the undiffcrcntiatcd plains(Lopes et al., 2010). The proposcd processes that formed this terrain hâve possibly tcetonic origin. (b) Dendritic networks as seen with SAR and morphological map(LORENZ et al., 2008).The System is locatcd at the western end of Xanadu close to our area of interest, Tui Regio. (c) Sand dunes around the Belot sand sea on Titan. The dunes are formed due to Aeolian prouesses. The bright figures are topographie obstacles that advance the formation of the dunes(RaüEBaUOH et al., 2009).

Fig. 5 .Fig. 6 .

 56 Fig. 5. Titan's surface from the Huygens probe(Tomasko et cil., 2005).

. 8 .

 8 Fig. 8. Reprojection of Tui Regio: (a) Gray scale 2.03 pm, (b) RGB false colors R: 5 pm, G: 2 g ni, B: 1.08 pm, (c) PCA R: l1' , G: 21"1, B: 3ri components, (d) PCA R: 3rt, G: 2"1, B: 1* components.

  ) showed areas of different colors and brightness suggesting diversity in surface composition. The PCA method is com patible to gray scale (a) and RGB (b) projections of Tui Regio. The visually brighter areas represent the highest I/F values and the darker areas the lowest, where I stands for the intensity of reflected light measured by the instrument and F the plane-parallel flux of sunlight incident

Fig. 9 .

 9 Fig. 9. Isolated areas wilh the use of visual colour alterations, which suggest areas of spectral différence.

Fig. 10 .Fig. 11 .

 1011 Fig. 10. Rcprojcction of Hotei Regio: (a) Gray scale 2.03 pm, (b) RGB false colors R: 5 pm, G: 2 pm, B: 1.08 pm, (c) PCA R: 1* , G: 2"1, B: 3"1 components, (d) PCA R: 2:i, G: 2*1, B: 1" components.

Plot 2 .

 2 Fig. 12. (left) Longitudinal Sand dunes in Saudi Arabia (NASA), (right) Longitudinal équatorial dunes on Titan[START_REF] Radebaugh | Linear dunes on Titan and earth: Ini tial remote sensing comparisons[END_REF].

13

 13

Fig 13 .

 13 Fig 13. (left) Supervolcano on Earth. Lake Toba (USGS). (right) Posàble supervolcano on Titan. Forboth figures, the central paît o fcaldtras b indicated by the white airows (Lofes et al, 2007b).

Fig. 15 .

 15 Fig. 15. (up left) Hairat Khaybar volcanic terrains, Nortli of Médina, Saudi Arabia (NASA), (up ri ght) Possible cryovolcanic terrain on Titan, Hotci Regio (Soderblom et al.. 2009). (down) Baglidad Sulcus, volcamc fracture on the south pôle of Enceladus, Tiger Stripes (NASA). Pockets of beat liave been appeared along the fracture. The whitc arrows indicate structures of tuffs. dômes and calderas forHarratKhaybar and possible doine and caldera formations for Hotei Regio. Additionally. the red arrows indicate lava flows and possible cryovolcame flows witliin the volcanic terrain. The ycllow dashed lines indicate areas that are possibly tectonic zones of weakness from which internai matenal may pass through.

  de Paris. CNRS, UPMC Univ Paris 06. Univ. Paris-Diderot-Meudon, France 2LISA-IPSL, CNRS/UPEC & Univ. Paris Diderot, 61 Avenue Général de Gaulle, F-94000 Créteil. France 3National & Kapodistrian University of Atliens, Athens, Greece 1. Introduction and Context Charles Darwin's vision, in 1859, of the origin and évolution of life on Earth, paved the way for future biological searches and studies on our planet and on other planetary bodies. The DNA decoding by James Watson and Francis Crick (1953) confirmed both the complexity and the symmetry of living particles.

Figure 1 .

 1 Figure 1. Map of the habitable zone's (HZ) limita Earthisin the middle ofthe HZ for thesolar System while Mars and Venus lie in its boundaries. The planetary bodies <?, b, c, and dare exoplanets. Gliese 581 is reddwarf star 20.3 light years from Earth (NASA).

Figure 1 .

 1 Figure 1. (lefi) Titan's thick, orange, and smoggy atmosphère (NASA); (right) Earth's atmosphère as seen from space (NASA); the cloudy nature of both atmospheric envelopes and also commonalities.

Figure 3 .Figure 4 .

 34 Figure 3. A scheme of possible théories on the origin of terrestrial and extraterrestrial life. The chemistry of the raw materials as well as energy sources plays key rôles to each evolutionary path (From[START_REF] Davis | Origins of life: a comparison of théories and application to Mars[END_REF] andMcKay et al. (2008)).

Figure 5 .Figure 6 .

 56 Figure 5. Titan's méthane (upper) and Earth's water (lower) cycles in the atmosphère (USGS). Titan h as a methane cycle resembling Earth's water cycle in its processes

Figure 9 .

 9 Figure 9. Lakes discovered in Titan's north-polar région by the Synthetic Aperture Radar (SAR) on board Cassini-Huygens mission (NASA/JPL). The dark patches are believed to be filled with hydro carbon liquid.

Figure 10 -

 10 Figure 10-Map of almost 655 lak.es and sea features by the Cassini radar System in azimuthal projection at the north pôle of Titan. Map A (left) shows the radar swath mosaic up to May 2007 flybys. Map B

Figure 12 .

 12 Figure 12. Enceladus's plumes ejected from the south-polar région captured by Cassini/ISS (Porco

Figure 13 .

 13 Figure 13. ([lefl) Modeling of a possible internai océan on Enceladus filled with water and Chemical compounds. This model tries to explain the significant abundance of sodium salts of Saturn's E-ring, probably originating in Enceladus plumes(Postberg et al., 2006). (right) Internai model of Enceladus based on tectonic meltwater. Heating along fractures is caused by tidal tlexing(Nimmo et al., 2007).

Figure 14 .

 14 Figure 14. Indications for biological and nonbiological processes in relation to molecular parameter (left) terrestrial abiogenic production of hydrocarbons from McCollom and Simoneit (1999) and ils biogenic pendant from Devai and Dclaunc (1996). (right) Organic distribution for abiotic. biological, and possible alien life on Enceladus (Both figures from McKay et al., 2008).

Fig. 1 .

 1 Fig. 1. Variety of interesting géologie features on Europa (NASA).

Fig. 2 .

 2 Fig.2. Schematic formation of intersect lineation[START_REF] Park | Foundations of Structural Geology[END_REF].

Fig. 4 .

 4 Fig. 4. (left) Lenticulae on Europa (reddish semi-circular spots in the middle of the image, (right) Conamara

Fig. 5 .

 5 Fig. 5. Bright terrain's major geological features in Ganymede. (a) Uruk Sulcus région is filled with ridges (red dashed line), grooves, craters (white arrows) and generally displays a smooth area, (b) Nippur Sulcus région display an extensive crater overlying ridges and troughs. (c) Sippar Sulcus région contains a large curvilinear scarp or cliff or possibly a caldera. If this structure is identified as caldera then it is the basic evidence for a surficial édifice of active past or présent cryovolcanism on Ganymede (NASA/JPL/Brown University).

Fig. 6 .

 6 Fig. 6. Dark terrain's major geological features on Ganymede. (a) Ancient impact craters visible at the middle left. (b) Grooves -northwest to southeast sets of fractures that possibly traverse a chain of craters. (c) Tectonics in the dark terrain: sets of ridges and grooves and fault blocks traverse the extensive crater located at the lower left part of the image and deforms it. (d) Sériés of scarps eut through the heavily cratered and old dark terrain (NASA/JPL/Brown University).

Fig. 7 .

 7 Fig. 7. Transitional région on Ganymede (dark terrain to bright terrain) separated by the red dashed line (NASA/JPL).

Fig. 8 .

 8 Fig. 8. Galileo Regio on Ganymede. Image taken by Voyager (NASA).

Fig. 9 .

 9 Fig. 9. Tectonic structures on Titan, (a) parallel mountains (NASA), (b) long, dark ridges spaced around one to two kilometers apart (NASA), (c) rectangular river network that possibly lie over faults that control the direction that methane can flow across the surface (NASA/JPL/Devon Burr), (d) canyon Systems along with bedrocks, channels and high cliffs (NASA/JPL).

Fig. 10 .

 10 Fig. 10. Sotra Facula, a possible cryovolcano on Titan. The peaks are more than 1 km high and the craterscaldera almost 1.5 km deep. Furthermore, flow features about 100 meters thick are obvious following a radial pattern around the craters (NASA/JPL-Caltech/USGS/University of Arizona).

Fig. sand 11 .

 11 Fig.sand

Fig. 12 .

 12 Fig. 12. Pôles of Enceladus. (left) Heavily cratered terrain of the north pôle, (right) Tectonized terrain (ridges, grooves, rifts) with fissures that emanate cryovolcanic material (NASA/JPL).

Fig. 14 .

 14 Fig. 14. Stratigraphie models of Europa's interior. Warm convective ice below the ice crust (a); Liquid océan under the ice coverage (b) (NASA/JPL).

Fig. 15 .

 15 Fig. 15. Two models for the icy crust thickness. The thin ice model ~200m (a) the thick ice model ~15 km (b) (NASA/JPL).

Fig. 17 .

 17 Fig. 17. Ice and minerai deposits on Ganymede. Surface features from Voyager in visible light (left); Minerais in red and ice grains in blue in infrared light from Galileo (NIMS, Galileo Mission, JPL, NASA).

Fig. 19 .

 19 Fig.19. Possible schematic location of océans in the icy moons of Jupiter as a function of depth. Europa probably agréés with internai structures 3 (thick upper icy layer <10 km and a thick océan) or 4 (very thin upper icy layer 3-4 km). Ganymede is doser to 1 (completely frozen) or 2 (three-layered structures impeding any contact between the liquid layer and the silicate floor)(Lammer et al. 2009). The color scales on the right side indicate the physical and Chemical constraints on which habitability dépends on.

  spectrometry on two probes entering tire atmosphère of Satum at two different locations near mid-latitudes, down to a pressure of 10 Bar. The global composition of Satum will be investigated through these measurements, together with microwave radiometry détermination of H20 and NH? and their 3D variations. The dynamics of Satum's atmosphère will be investigated ffom: (1) measurements of pressure, température, vertical distribution of clouds and wind speed along the probes' descent trajectories, and (2) détermination of deep winds, differential rotation and convection with combined probe, gravity and radiometric measurements. Besides these primary goals, Kronos will also measure the intensifies and cliaracteristics of Satum's

Fig- 1

 1 Fig-1 Interiors of Jupiter and Satum highlighting the impor tance of hydrogen and hélium for the structure of these two planets (yellow indicates that hydrogen is in molecular form, red that it is metallic, and the central dense cote is shown in blué)

Fig. 2

 2 Fig.2Elemental abondances measured in the tropospheres of Jupiter (top) and Satum (bottom) in units of their abundances in the protosolar nebula (from[6]). The elemental abundances for Jupiter are derived from the in situ measurements of tire Galileo probe. The abundances for Satum are spectroseopic déterminations from Cassini for He/H and C/H, and model-dependant gmund based measurements for N/ H and S/K. Note that the retrieval of the hélium abundance is indirect and un certain, so that very précisé He/H measurement is needed. Kronos \vill allow distinguishing between different formation scénarios whose prédictions are shown as green, blue and pink cimes, respectively (see text)

Fig- 4

 4 Fig-4 Mean vertical distribu tion of cloud layers on Satum,

5

 5 General circulation of Satum and relevant atmospheric features on its atmosphère. Wînds at cloud level, relative to the interior reference frame measured by Voyager, traced by the Voyagers (grey lùié) and Cassini data of the Southemmost latitudes (blue) and équatorial région in different filters (red and violet). Relevant meteorological structures appear on the insets4, a North polar hexagon in visible (Voyager 1) and infrared light (Cassini); b The Ribbon; c Satum Great White Spot in the Equatorial Région in 1990 and tlie State of the equator as seen by Cassini in the methane absorption band and continuum filters; d The South Polar jet and the inner polar vortex; e Convective storms seen by Cassini; f Anticyclones from Voyager 1.

  Observational constraints on such (deep) atmospheric dynamics are very difficult to obtain, not least because of some fondamental 'degeneracies' in tlie near-surfoce signatures of deep and shallow circulations i.e. very different processes (deep or shallow) may produce surface signatures that are more or less indistinguishable. A major objective of Kronos is therefore to obtain measurements that can résolve some of these ambiguities and ascertain which aspects of Satum's meteorology and circulation are deep-seated and which are shallow. step, Kronos will obtain the first detailed in situ measurements of the vertical structure of Satum's atmospheric température, winds, radiative balance and cloud structure, to depdis of up to 10 bar wliere duvet solar heating is essentially negligible. The factors determining the atmospheric thermal and density structure in the outer planets dépend on a complex mix of radiative heating and cooling, turbulent convection (both 'moist' and 'dry') and larger scale circulation Systems. The distribution of radiati ve heating in the upper troposphère is affected strongly by the presence of clouds and hazes, whose distribution is itself detennined by tlie ambient circulation, and the concentration of various volatiles. Figure 5 shows a much simplified 'typical' vertical distribution of clouds, modeled by assuming a particular composition at depth and adiabatic uplift until different components become saturated and begin to condense. Such a distribution may reflect a global mean picture, but observations indieate considérable local variability.

Figure 6

 6 shows a Cassini V1MS image of Saturn at 5 pm which reflects the distribution of clouds and convection ai 2-4 bar. This indicates that convection (both dry and moist) is organized by the winds in the deep troposphère into zonal bands on a finer scale than at the top of the troposphère. This drastic change, évident in the meteorology between tlie visible face of Satum and its expression at just 5 bar, is indicative of the changes in the driving forces as one delves deeper into Satum. We expect that by directly sampling the internai structure down to 10 bai-, below most of the direct radiative heating/cooling of the atmosphère, we wi ll hâve a more direct picture of the processes controlling the dynamics in Satum's atmosphère driven from deep inside tlie planet radier than from solar input/radiative cooling.The primary goals of the meteorology and atmospheric dynamics investigation on Kronos will be to détermine the nature of die deep circulation, differential rotation and convection in Satum's atmosphère by combining observations from die Kronos probes widi remote sensing measurements from the carrier spacecraft. For the probes, this will require in situ détermination of vertical profiles of température, horizontal winds and cloud properties (aérosol particle densities, size distributions and composition, including possible chromophores) throughout the troposphère to a

Fig- 7

 7 Fig-7 Satum 's magnetospheric régions and processes. Bottom inset new radiation belt observed in ENA, top inset électron beams observed by Cassini and mapped to the aurora along field lines; auroral image by HST; orange cône sketches auroral (SKR) radio émissions

Finally,

  Kronos should provide new insights to the inner radiation belt discovered by Cassini inside the D ring[START_REF] Krimigis | Dynamics of Satum's magnetosphere from MIMI during Cassini's orbital insertion[END_REF],2.4 Ring scienceWhereas observed since tlie seventeenth century, Satum's rings aie still one of tlie mostpuzzling structures in our Solar System. To constrain the origin of the rings and their évolution, in situ observations aie needed, foliowing Cassini's success. In addition, Satum's rings are tlie closest example of an astrophysical disk, one of the most fondamental structures in the Uni verse. The outer edge of the rings System (the A-F région) is located on tlie Roche limit itself, allowing for substantial accretion processes. The outer régions of the rings therefore share common properties with a protoplanetary disk.

3. 3

 3 TrajectoryInterpianetary trajeetories to Samm could be supported with Chemical and/or solar electric propulsion (SEP) Systems, combined with gravity assist at Earth and Venus, and for the relay trajectory architecture also at Jupiter. These trajeetories resuit in flight times from -6.3 years to ~17 years, although flight times over 12 years were found not désirable. Inteiplanetary and flyby/orbiter trajectoiy options are different for suppoiting relay or DTE architectures.

3. 4

 4 Direct to Earth communication 3.4.1 Flyby DTE trajectory Using a direct-to-Earth (DTE) communication strategy découplés the probes from the carrier (after release) and thus allows for the most efficient, lowest cost and lowest risk mission architecture possible. Tins, of course, limits the data retum from the probe mission, however, oui-preliminary analysis shows that the DTE capability meets the requirements for the Kronos mission. The probes are delivered on a Type II trajectory to Satum to allow for probe entry at or near the sub-Earth point. Preliminary results suggest a longer flight time associated with DTE due to limits on the Jupiter gravity assist and to the probe entry geomctry requirements. Preliminary launch characteristics for DTE and relay cases are shown in Tables

  potentially provides the longest visibility between the probe and the carrier. This requires the relay spacecraft to be in a low inclination or équatorial orbit at a radial distance > 5 Rs. Zenith atténuation of radio signal is a fonction of probe depth, measured by atmospheric pressure. During tlie probe descent, the transmission of data must be perfonned in real time, with the data upload ending as the probe reaches its deepest point in Satum's atmosphère. The preliminary DTE architecture includes a 200 MHz UHF link from the probes for direct-to-Earth communication to Earth, and an X-band link front tire carrier for orbiter or flyby science and telemetry through tlie DSN. Data rate from the probe is the function of the antenna size, transmitter power, séparation distance, and atmospheric conditions. Data compression potentially can increase the efficiency and tlius the scientific retuni. Technology development in the areas of antenna design, telecom power, and frequency scanning and locking can vastly improve the expected performance.

  Stmcturally, tlie carrier consists of: (a) a spacecraft bus; (b) a propulsion System, including propellant tanks; (c) a power System, consisting of LILT solar arrays and batteries; (d) subsystems, including thermal management, CDH, GN&C, AC S, telecom; (e) and science instruments, including the MWR antennas. For tlie propulsion System both Chemical and SEP were considered.3.5.1 Carrier pawer System tradesAt Satum, the solai' flux is about 15 W/m2, or ~1% of that at Earfh. Low Intcnsity Low Température (LILT) solar eells, developed by ES A for Rosefta, were tested up to -70 W/m2, while Juno's solar panels are designed to operate at Jupiter where the solar flux is -55 W/m2. New, high efflciency LILT cells could achieve a power output per unit area of 4 to 5 W/m2, assuming a conversion efflciency of 30%. At this level of performance, according to a CNES study, a Juno-size solar panel could generate -200 We, wiüi the power System weighing -300 kg (including batteries and conditioning System). Performance could be also increased with the development of concentrators.

3. 5 . 2

 52 Satum probes Satum's atmospheric circulation rcsults in a distinct distribution of latitude bands and zones as seen most recently in the Cassini VIMS spectral data. Thus the Kronos mission should allow for the launch of two probes in different régions. The preliminaty design calls for targeting one of the probes inside and one outside of tlie ±13°latitude band. Silice the scale-height of Satum is about twice tliat of Jupiter, it

Fig. 8 Heritage

 8 Fig. 8 Heritage Galileo descent module and probe (Galileo Probe Décélération Module Final Report, DocNo. 84SDS2020, General Electric Re-entry Systems Operations, 1984; AIAA, "Project Galileo Mission and Spacecraft Design", Proc. 21st Aerospace Science Meeting, Reno, NV, Januaiy10 13, 1983; Proc.

1

 1 Elemental and isotopic measurements of Satum's atmosphère Satum Probe measurements will be selected to meet the measurement requirements that address (1) théories of planetary origin and (2) fundamental knowledge of Satum's atmospheric structure and meteorological processes. The instrument complément will include a mass spectrometer for Chemical and isotopic composition détermination, an atmospheric structure experiment, a Doppler Wind Experiment for retrieval of a vertical profile of zonal winds, and a nephelometer to détermine the structure of the atmosphère to at least 10 Bar and to obtain the radiative balance at this depth. The instruments baselined for the Kronos Probe investigation hâve a solid

  pressure and température profile by measuring décélération of the entry veliicle and performing direct température and pressure measurements during the descent phase.The ASI-ACC will start to operate since the beginning of the entry phase, sensing tlie atmospheric drag experienced by the entry veliicle. Direct pressure and température measurements will be performed by the sensors having access to the atmospheric flow from the earliest portion of the descent until the end of the probe mission at approximately 10 Bar.4.6 Measurements witli a Doppler wind experimentThe primary goal of the Kronos Doppler Wind Experiment (DWE) is to measure a vertical profile of the zonal (east west) winds along the probe descent path[START_REF] Atkinson | Galileo doppler measurements of tlie deep zonal winds at Jupiter[END_REF]. A secondary goal of the DWE is to detect, characterize, and quantify microstructure in the probe descent dynamics, including probe spin, swing, aerodynaniic buffeting and atmospheric turbulence, and to detect régions of wind shear and atmospheric wave phenomena. The Kronos Doppler Wind Experiment (DWE) can be designed to work with a probe DTE architecture or a probe-to-relay architecture. Both options include USO requirements and differ only in the angle of entry and DTE geometry requirements (sce Mission Design). For rclay, the System comprises a probe and a carrier ultras table oscillator (USO) as part of the probe/carrier communication £) Springer 544 package. ïnvolveraent of several observing télescopes in interferometric mode (PRIDE) will significantly improve the accuracy and robustness of the measurements. Tlic proposed experiment benefits fiom the strong héritage of both tlie Galileo and Huygens doppler wind experiments[START_REF] Atkinson | The Galileo probe doppler wind experiment: Measurement of the deep zonal winds on Jupiter[END_REF].4.7 Instruments on the carrier spacecraftFurther définition of the carrier spacecraft will take place during the Kronos Phase A study. This may be a flyby spacecraft or a Satum orbiter. Définition of the orbital instruments will also take place during this period. A key objective of the payload on the carrier spacecraft will be to secure deep, global abundances of H20 and NH3 and their distributions over al latitudes using microwave radiometry. The technology to implement these measurements is mature through the radiometers developed for the Juno mission. The mission objectives for the orbiter/flyby spacecraft are inherited ftom the Juno mission and include;•Measurements ofthe global oxygen (water) and nitrogen (ammonia) abundances. • Measurements of Satum's internai mass distribution (core mass) via gravity science • Measurements of Satum's internai rotation and convection (whether Satum rotâtes as a solid body) via gravity science (détermination of high order gravitational moments; see) • Measurements of Satum's internai magnetic field sufficient to investigate the source location of the field.• Measurements of the propeities of Satum's innennost magnetosphere (inner radiation belt, possible ring-associated cuixents, UV/IR/radio auroras,

4. 7 . 2

 72 Xj'Ka band up/downlink System for gravity science Most of the unceitainty in our knowledge of Satum's core stems from unceitainties in the équation of State and tlie gravity zonal harmonies J4 and J6 and to limited knowledge of how Satum's deep interior rotâtes. Current unceitainties are 1% in J4 £) Springer and 30% in J6. Kronos improves the knowledge of J4 and J6 by over 3 orders of magnitude (J4 -10 ppm accuracy, J6 -100 ppm accuracy). This removes the uncertainty regarding interior rotation and estimâtes the abundance of water, reducing the core mass uncertainty and total mass of heavy éléments to a few Earth masses. Using an existing Linear Ion Trap (LITS) as frequency reference (or an improved follow on) a DSN ground station (DSS25) transmits X and Ka-band radio signais to the spacecraft, then receives tire transponded signais and veiy' accurately measures tlieir frequencies. Simultaneously, water vapor radiometers (WVR) at tire DSN station measure Earth-atmospheric briglrtness températures in the direction of tire antenna beam, tlrereby determining the zénith signal delay caused by the wet component of tire troposphère, and ultinrately tire tropospheric confection (refraction) to the raw Doppler data. The same set of on-board instrumentation on the camer spacecraft will enable nrulti-disciplitrary VLBI tracking experiments witlr ultraprecise characteiization of the State vector of the spacecraft, reaching tens of meters in the latéral direction.4.7.3 Vector magnetometer The instrument must be of sufficient sensitivity to measure perturbations caused by field aligned ciments. The successful operation of the Cassini fluxgate magnetometer at Saturn nrakes it a natural cairdidate for inlieritance for the design of the Kronos magnetonreters. The new génération of digital fluxgate sensors hâve masses of the order -200 g, and therefore nrake very little demand on overall payload mass.

4. 7 . 4 4 . 7 . 6

 74476 Plasmas sensorThe instrument must quickly measure électron and ion energy angular distributions (100 eV-20 keV) fronr an effectively non-rotating platform.4.7.5 Energetic particle sensorTlris instrument must operate at tire higher energies (20 keV-several MeV) while the plasma sensor fonctions at lower energies. Altliough a 1 MeV upper limit is sufficient to address magnetosphere-ionosphere coupling issues, Irigher energies (up to are needed to study the radiation belts. In situ study of the new radiation belt seen by Cassini-MIMI inside the D-ring, would require a small ridealong particle detector on the probe itself. Such a sensor was carried by Galileo probe. Plasma and radio wave receivers Cassini and Sterco spacecraft carry radio and plasma waves spectro-polarimeters which can be programmed to operate according to very diverse setups (spectral and ô Springer 546 temporal résolutions, snapshots or surveys, wavefoim capture, polarization measurement, etc.). In addition, they hâve goniometric capabilities, i.e. they enable the instantaneous détermination of the direction of the brightest point source at each measurement, and thus the possibility to synthesize images of the radio-sources. The next génération of these radio spectro-goniopolarimeters has been miniaturized and could be implemented for resources of ~1 kg and ~1 Wusing development héritage ffom e.g. BepiColombo/MMO.4.8 The Lora (Landers On Rings Array) MicroProbesAil the ring science objectives can be met by two to three very simple probes (MicroProbes), whose sole payload would be a high-resolution caméra, dropped in different locations of the rings (A, B, C by order of priority). The only limit to the resolution is the time for data transfer before impacting into the rings. A Ritchey

Fig

  Fig. 10 Synthetic images of A ring wakes at 60 m resolution (left) and 60 resolution (nght) (H. Salo)

4. 9

 9 Satum atmosphère wind diagnostics by means of radio measurements of the probe motion (PRIDE) This experiment does not require spécial on-board instrumentation and can be conducted with radio signais in both communication scénarios, DTE and probe-tocarrier relay. In particular, it will provide additional input into measurements of tlie wind profile not requiring an a priori assumption on a one-dimensional mode) of tlie wind. DTE could be used for additional measurements by determining an additional component of tlie probe's State vector. Indeed there is no need to limit tlie experiment to zonal component measurements since a méridional component could be measured as well. Even 3D diagnostics of the motion are not out of the question, i.e. the descent velocity can be included in the set of measured values front the experiment. To what extent the measured motion of tlie probe represents the wind dépends on spécifies of the atmosphère fliglit. A combination of Doppler and interferometric measurements can provide unambiguous détermination of the descent trajectory-coordinates and tlieir (at least first) dérivatives.

548 5

 5 Conclusions and summaryMultiple probes to the giant planets are critical for collecting tlie data rcquired for understanding the formation of our solar System, and by extension, of extrasolar Systems. Indeed, our giant planets can be studied in much finer detail than any exoplanet ever will, and they hâve preserved Chemical and isotopic relies of early solar System réservoirs. In exploring the composition of Satum vvith Kronos, we shall go back to the times when a tiny fraction of a molecular cloud collapsed to give birth to a spécial stellar System: the Sun and its planets. Tlie probes will allow an indeptli exploration of Satum and its comparison with results obtained by the Galileo mission at Jupiter. This comparison between the two most massive planets of tlie Solar System is a prerequisite to a detailed understanding of the origins of the Solar System. Kronos shall also investigate Satum's complex atmospheric dynamics, which is a fantastic laboratory for this field of research. It. will measure the planet's magnetic field very close to the planet's atmosphère and thus unfîltered by the rings, with direct conséquences for understanding tlie planet's dynamo and deep rotation.

Figure 2 :

 2 Figure 2: The operation of the MEMS network afterthe touchdown of the probe on the hydrocarbon liquid surface

  Icy inoons of the Outer Solar System are extremely lnteresting planetary bodies since they possess evidence about the origin and évolution of their Systems and tlic Solar System in général. In ternis of surface inorphology, Internai structure as well as their environmcntal uniqueness, Saturn's moons Titan and Enceladus and Jupiter's Europa and Io are the best représentatives In thaï perspective. In tliis study, we propose a seismic experimcnt for the icy moons as a payioad of future missions' landers. This suite of instruments comprises a gas monltoring sensor (micro Gas-Chromatograph) which will operatc along with tlie seismic sensor. Data froin the micro Gas-Chromatograph will give us the opportunily to correlate each tlme tlic rccordcd seismic data of the current gas-relief activity.

  ground vibration. It mainly contai 11s the seismometer and the unit which records tlie signal. The seismometer consists of three sensors placed in the same sealed case. Each sensor can be described as a pendulum tliat moves from its equilibrium position triggered by the ground movement. 'Ilie sensors can measure any ground motion within a frequcncy range of 0.001 Hz to 100 Hz usually, at the north/south, east/west and vertical component in orthogonal System.

Fig. 1

 1 Fig.1shows the basic functional procedure of a seismic instrument on icy moons according to tlie requirements mentioned above. The configuration of tlie System is illustrated in a simplified plot. Ail the components inside tlie dashed line in Fig.1are parts of tlie same physical equipment. tlie seismograph.

Figure 1 .

 1 Figure 1.Functional procedure of the seismic instrument

Figure 2 .

 2 Figure 2. Operstional flow diagram of the micro Gas Chromatograph

  Of A Future Seismic Experiment On Outer Solar System Terms and Conditions of the implémentation of seismic instrumentation on icy moons and the possibly emerged problems and difficulties are listed in the Table II below.

  .1 below.

	CQVIH OEPLOYWENT AOUATOR
	SECONOARï WlRROR
	BAFRE
	RMXATOR/SUNSrtAOE
	PR1MAKY MIRROR
	. ENTRANCC APER1WE
	* ROOr-TOP MKROR
	RR COUJMAIOR
	FIR fOCUS UlRROR
	ANO FP1

Table 2 . 2

 22

: Focal Planes spécifications

  exists, for example, if eventual detector dégradation with time would resuit in a slow systematic drift of the radiation détection efficiency. After personal communications with Dr. D. Jennings, Dr. C. Nixon and Dr. R. Carlson from Goddard/NASA, such a case is not possible. No significant dégradation of the CIRS radiance calibration has been observed over the mission of any of the 21 CIRS detectors.

Table 2 .

 2 7. The abundances of the major trace gases when fitting each sélection and their différence

	Molécule	Grand Average -fit	DS4000 -fit	Différence
	C2H2	2.4x10"6	2.1xl0'6	+14%
	HCN	5.6x10'8	4.7x10'8

+20%

C2H6

Table 2 .

 2

8: CIRSFRANCE nadir sélections used for the présent test. I hâve gathered data from the DS4000

  2.12 -2.15 in black. DS4000 spectra FP4 on December 2006 (Fig.

Table 2 .

 2 10: The abundances of the major trace gases when fitting the January 2010 sélections and their

	différence.
		January 2010 abundances
	Molécule	v2.5	DS4000	Différence
	c2h2	5.60x10-6	5.10 xlO-6	9.8%
	HCN	8.00x10'7	7.90 xl0'7	1.3%
	C3H4	1.22x10'8	1.38x10'8	-11.6%
	c4H2	5.50x10'9	5.20x10'9	5.8%
	co2	1.55 xlO'8	1.39 xlO'8	11.5%
	C2H6	1.22x10'5	1.05x10'5	16.2%

Table 2 .

 2 11: The abundances of the major trace gases when fitting the December 2006 sélections and their

	différence.
		December 2006 abundances
	Molécule	v2.5	DS4000	Différence
	c2h2	3.90x10'6	3.10 xlO"6	25.8%
	HCN	4.70 xl0'7	4.30x10'7	9.3%
	c3h4	1.54x10'8	1.60x10"8	-3.8%
	c4ii2	8.10 xlO"9	8.10 xl0'9	0.0%
	co2	1.60x10"8	1.70x10"8	-5.9%
	c2h6	1.15x10"5	1.15 xlO'5	0.0%

We now focus on the continuum level in both pipelines. The continuum of December 2006 FP3 sélection (Fig.

2

.20) is higher in the DS4000 data than in v2.5 on the order of 45%, 20%

Table 2 .

 2 12. The différences between the continuum level of both pipelines at the December's 2006 sélection within the région around C3H4 and C4H2. DS4000 continuum is higher than the v2.5.

		Différence between
	December 2006 -FP3 continuum level (cm1)
		DS4000 and v2.5
	575 -590	40-50%
	590 -625	20%
	640 -650	13%

Table 2 .

 2 13. The différences between the continuum level of both pipelines at the FP4 December's 2006 sélection.

Table 2 .

 2 14: September 2010, equator, 0.5 cm1. Only slight différences in Signal-to-Noise ratios exists.

	Focal Plane	C1RSDATA version	Number of spectra	Emission angle	Signal-to-Noise ratio
	FP3	v3.2	1503	11.6	60.6
	FP3	DS4000	1503	11.6	60.2
	FP3	DS4000-PhC	1503	11.6	57.8
	FP4	v3.2	1619	7.6	138.0
	FP4	DS4000	1619	7.6	137.8
	FP4	DS4000-PhC	1619	7.6	143.7

Figs. 2.24 and 2.25 below depict the comparison of these three database inferences. The upper plots contain ail the database queries, while the lower panels of each figure show the différence with the DS4000-PhC and the nominal DS4000. It should be noted that the

Table 2 .

 2

	CIRSDATA	C2H2	c2H6	HCN
	v2.5	3.4x10'6	9.2x10'6	2. lxl O'7
	v3.2	3.9x10'6	9.8xl0"6	2.1xl0'7
	DS4000	3.9x10'6	9.8x10^2.1xl0'7
	DS4000-PhC	3.8x106	9.7x10'6	2.0x10'7

15: C2H2, C2Hé, HCN abundances as derived by applying our radiative transfer code in FP3 September 2010 averaged spectra at equator (0.5 cm1).

Table 2 .

 2 16: July 2004, 50°S, 2.5 cm'1. The number of spectra for FP3 and FP4 between the local calibrated database and the deep space calibrated differ significantly. Consequently, the other characteristics also differ.FP1 spectra are the same.

	Focal Plane	CIRSDATA version	Number of spectra	Emission angle	Signal-to-Noise ratio
	FP3	v3.2	982	33.2	117.2
	FP3	DS4000	97	26.7	34.5
	FP3	DS4000-PhC	97	26.7	33.2
	FP4	v3.2	929	37.1	267.3
	FP4	DS4000	929	37.1	268.5
	FP4	DS4000-PhC	929	37.1	267.7
	FP1	v3.2	111	20.8	130.3
	FP1	DS4000	111	20.8	130.3
	FP1	DS4000-PhC	111	20.8	130.3

Table 2 .

 2

	Focal Plane	CIRSDATA version	Number of spectra	Emission angle	Signal-to-Noise ratio
	FP3	v3.2	545	23.0	24.7
	FP3	DS4000	545	23.0	24.2
	FP3	DS4000-PhC	545	23.0	24.2
	FP4	v3.2	356	3.9	30.3
	FP4	DS4000	356	3.9	30.7
	FP4	DS4000-PhC	356	3.9	30.7
	The following Figs. 2.31 and 2.32 show the comparison of these three database inferences for
	the January 2007 50°N query.

17: January 2007, 50°N, 0.5 cm'1. As far as the recorded number of spectra are concemed as well as the émission angles, there are no différence. Any différences in the Signal-to-Noise ratios are insignificant. 1500 Fig. 2.31 : Comparison between the current v3.2 database (blue) and the DS4000 one (red) FP4 averaged spectra of January 2007 at 50°N from C1RSDATA. The deep space calibration éliminâtes the noise in the continuum.

Table 2 .

 2

	CIRSDATA	c2h2	c2h6	C3H4	c4h2	HCN
	v2.5	2.5xl0'6	1.07x10'5	2.0x10'8	9.5x10'9	5.8xl0'7
	v3.2	2.8xl0'6	1.13xl0'5	2.1xl0'8	9.8xl0'9	6-OxlO'7
	DS4000	2.6x10'6	1.05x10'5	1.9x10'8	1.0x10'8	5.8x10'7
	DS4000-PhC	2.8x10'6	1.05x10'5	1.9x10'8	1.0x10'8	6.0x10'7

18:

The abundances of the major molécules as derived by applying our radiative transfer code in FP3 January 2007 averaged spectra at 50°N (0.5 cm'1).

Table 3 .

 3

					Total
	Focal Plane	Year	Month	Latitude	number of	Signal-to-noise	Airmass	Cassini Flyby	ls n
					spectra
	FP3	2008	Feb-Mar-Jul	50°N	145	12.4	1.19	T41-T45	341-347
	FP4	2008	Feb-Mar-Jul	50°N	166	12.8	1.14	T41-T45	341-347
	FP3	2010	Jan	50°N	284	21.3	1.69	T65-T66	005-006
	FP4	2010	Jan	50°N	234	27.6	1.66	T65-T66	005-006
	FP3	2009	Mar-Jul	0°86	18.3	1.75	T51-T59	355
	FP4	2009	Mar-Jul	0°84	39.2	1.71	T51-T59	355
	FP3	2010	Sep	0°1504	50	1.2	T72	014
	FP4	2010	Sep	0°1622	136.9	1.01	T72	014

1-List of CIRS queries used for estimation of the effect of different température profiles in radiative transfer calculations. Ail sélections refer to high resolution spectra (0.5 cm'1).

Table 3 .

 3 2-Retrieved abundances by using both CIRS-and HASI-based profiles for two queries at 50°N.

				Différence			Différence
	Molécule	FEB 2008	FEB 2008	HASI vs	JAN 2010	JAN 2010	HASI vs
		CIRS-based	HASI-based		CIRS-based	HASI-based
				CIRS			CIRS
	C02	1.73xl0'8	1.80x10'8	4%	1.60x10'8	1.55x10'8	-3%
	c2H2	4.70x10'6	5.60x10'6	19%	5.30xl0'6	6.30xl0'6	19%
	c2h4	1.90x10'7	2.00xl0'7	5%	2.30xl0'7	2.30x10'7	0%
	c2h6	1.20x10'5	1.27xl0'5	6%	1.3 5x10'5	1.35xl0'5	0%
	C3H4	1.80x10'8	1.82x10'8	1%	1.18xl0'8	1.17x10'8	-1%
	c4H2	9.50x10'9	9.50x10'9	0%	5.50xl0'9	5.40x10'9	-2%
	HCN	9.50x10'7	1.20x10'6	26%	9.00xl0'7	1.50x10'6	17%

Table 3 .

 3 3-Retrieved abundances by using both CIRS-and HASI-based profiles for two queries at equator.

		MAR-JUL	MAR-JUL	Différence	Différence
				SEP 2010	SEP 2010
	Molécule	2009 CIRS-	2009 HASI-	HASI vs	HASI vs
				CIRS-based	HASI-based
		based	based	CIRS	CIRS
	C02

Table 3 .

 3 4-List of the spectroscopic parameters adopted for the ARTT upgrade. The source the line lists of each molécule of the previous version is also shown.

	Molécule	ARTT 0.3.8	ARTT 1.0.0
	ch4	GEISA 2003	GEISA
	ch3d	GEISA 2003	GEISA
	CO	GEISA 2003	GEISA
	C02	GEISA 2003	GEISA
		HITRAN 2004 and
	HCN		HITRAN 2008
		HITRAN 2008
	hc3n	Jolly et al. (2007)	GEISA
	c2n2	GEISA 2003	GEISA
	c2h2	GEISA 2003	GEISA
			Vander Auwera et al. (2007)
			at 1350-1496 cm-1
	c2h6	Vander Auwera et al. (2007)	GEISA
			at 613-843 cm-1
		Blassetal. (2001) and
	c2h4	Rotger 2006 pers. comm.	GEISA
		Rotger et al. (2008)
	c3h4	G. Graner pers. comm.	G. Graner pers. comm.

Table 3 .

 3 5-The major molécules within the spectroscopic région of comparison

	Molécule	0.3.8	1.0.0	Comments
			HITRAN 2008
				The spectroscopic files are the same but
			(about 14% higher
				the one in 0.3.8 has "almost zéro" fines.
	HCN	HITRAN 2008	abundances with the
				These fines hâve been discarded from the
			new spectroscopic
				latest file in 1.0.0 version.
			files)
				No upgrade has been imported from the
	C2H2	GEISA 2003	GEISA 2009	2003 version to the latest 2009 at the FP3
				spectral région. The émission is the same.
				Change in the émission of v26 propane
				band.
				The GEISA 2009 fines file contains the
	v26 CîHg	GEISA 2003	GEISA 2009	GEISA 2003 ones multiplied by a factor
				of 0.420.
				The CIRS/team site has the same file with
				GEISA 2009 file.
				Différence in fits. The same fit is achieved
				with a value 7.5% higher in 1.0.0 than in
				0.3.8.
				Devi et al. (2010) hâve found that ethane
			GEISA 2009	in V9 fine intensifies are 10-15% lower
		Vander Auwera et	at 613-843 cm'1 and
	c2H6		Vander Auwera et	than in Vander Auwera et al. (2007)
		al. (2007)	al. (2007) at 1350-	which mean that GEISA 2009 values are
			1496 cm'1
				10-15% higher.
				GEISA 2009 fines intensifies are lower by
				a factor of 2.3% than the ones from
				Vander Auwera et al. (2007).
		Blass et al. (2001)		4% lower abundance for fitting the
	c2h4		GEISA 2009
		and Rotger 2006		molécule comparing to the 0.3.8 one.
				No significant change in the émission
				signature of the main isotopologue at 667
	co2	GEISA 2003	GEISA 2009
				cm'1. The abundance of 1.0.0 fit is 2%
				higher than in the previous version.
		Arié and Johns		GEISA 2009 adopts Jolly et al. (2010)
	c4h2	( 1992) for vg and	GEISA 2009	fine lists. No change in fits' abundances
		hot bands

Table 4 .

 4

	1 -Trace gaseous abundances of Titan's stratosphère from Coustenis & Bezard (1995) retrieved from
	Vl/IRIS spectra in 1980.
		1	>i		0	A. B.C		D		E	F	North/
		(l»î. in •> <53*S)	(32*SI		(7-Sl	I7°N)		<30*N)		<50"N)	(TO'NI	souih
	Molécule											4E/1)
	C,H,	2.5 £ 0.3 » 10*	3.0 ± 0.4 X 10'*	2.7 i O.î X I0'<	3.0 i 0.3 x 10'*	2.7	± 0.5 X 10'*	5.3	£ 1.0 X 10'*	6J ± 2.1 X KH	2.2 ± 0.5
	C;H,	1.5 ±0.5X10''	1.5 ± 0.3 x 10'7	I.J ± 0.3 X 10~'	1.5 ± 0.4 x 10-'	1.5 ± 0.5 x HT'	1.0	± 0.3 X 10 '*	1.5 10.5 x |0~*	7 i 3
	C,H,	1.15 ±0.1 x |0"3	1.35 ± 0.1 x IO"5	1.65 ±0.1 x |0->	1.3 ±0.1 x 10'»	1.23 ± 0.1 J x HT'	1.85 ± 0-2 X I0'5	1.63 ± 0.3 x 10'5	l.S ± 0.2
	C,H,	3.3 ± 0.7 X 10''	4 J ± 0.7 x ÎO'7	8.0 ± 0.6 x 10'*	5.0 ± 0.9 x 10'*	6.0 = 1.6 x 10'*	3.3	± 0.2 x I0-'	3.7 ± 0.3 X IO'*	9.5 ± 2
	C,H,	6.0 = 1.5 X 10-'	6.0 2 1.4 X 10*'	9.0 ±1.6 x jer*	3-0 ± 1.7 X 10 1	7.0 ± 3.0 x IO'1	1.0	•i oJ y io	1.2 ± 0.3 x IO"*	1.7 1 0.7
	C<Hj	1.0 1 0.3 X 10'*	\.ï ±0.3X10'*	1.7 ±0.1 X 10'*	1.4 IM* HT*	1.2 £0.4 X 10'*	l.S * 0.2 x 10'*	2.7 ± 0.4 X IO'*	15 3L 5
	HCN	4.7 ± 0.5 X 10-'	7.0 = 0 i x 10'*	2.2 ± 0.2 x IO'7	1.7 ± 0.2 x 10"'	2.4 ± 0.4 x 10*'	1.1	s 0.2 x 10*'	1.5 ± 0.3 x 10'*	23 s 4
	CjNt	<I.S X 10"*	<1.5 X 10"*	<1	X 10'®	<14 x 10-*		<2 x 10'*		± 0.2 x 10"*	2.2 ±0.3 x 10'*	>10
	»1C,N	<1.3 x ht*	<1.3 « 10"*	<1.5 x 10-"	<1.5 x 10**		<2 x 10'*	2.5 ± 0.5 x 10'*	4.5 ± 0.9 x 10'*	>17
	CO,	1.3 ±0.2 x IC*	l.S = 0.2 X 10'*	1.3 ±0.15X10'*	1.4 ±0.2 x 10'*	l.l ± 0.25 X 10'*	1.3	± 0.3 X 10'*	1.3 ± 0.55 x 10'*	1 ± 0.4
					Uazc optical dcplh for <r > 600 cm'1*
		0.83 * 0.06	0.93 = 0.06	1.34 ±0.13	1 00		1.00 ± 0.09		.11 ± 0 IO	2.83 ± 0 40	2.54 i 0.27
	* Relative to the équatorial oplical depth, calculated at wavenumbers 650. 900, 1000. and 1100 cm~:.
	A few months later, in August 1981, Voyager 2 flew by Titan, but at a distance 170
	times greater. The analysis of V2/IRIS spectra probing the atmosphère in the altitude range of
	80-130 km. V2 results confirm the ones from VI and the température structure which was
	retrieved from V2/IRIS showed a decrease from equator to northem latitudes of about a 10 K.
	No significant temporal variations were observed within this time interval of 9 months
	between the two encounters. No significant latitudinal variations were recorded for acetylene,
	ethane and propane within the error bars. Like in the Vl/IRIS spectra, HCN showed a
	significant increase at higher northem latitudes (50°N) by a factor of 2. C4H2 and C3H4 also
	presented the same enhancement. Only upper limits were derived for C2H4, HC3N, C2N2 and
	CH3D. The results are listed in Table 4.2 below (Letoumeur & Coustenis, 1993).

Table 4 .

 4

		2 -Trace gaseous abundances of Titan's stratosphère from Letoumer & Coustenis (1992), retrieved from
	V2/IRIS spectra in 1981.
			Sample A	Sample B	Sample C	Sample D
		Abundances	( -B + C + D)	50 N	28 N	6 N
		Acctvlcnc (C\HT	2.6! j ,x 10"	2.6'! îx 10 "	2.6! J < x 10 •'	2.7! J j x 10 "
		Ethylène (C:Hj	<1 x 10"	-	-
		Elhunc	i.4*;;ixio-?	1.92-ï.ÜkIO-5	i.4:SSxio-'	l.2+itxl0-5
		Methyl acctylcnc (C,H4)	6.0: Cx 10 v	1.2^xIO"	6.0!;, x 10"	6.0'Cx 10-'*
		Propane (C,H*)	5dx I0-'	2.2+1** x 10"	5.0" x 10 "7	7,0' t! x 10 '
		Diacctvlcnc (C4H.)	I.4tr*l0"	2.81v?x 10"**	1.4! Six 10 "	1.3! ?K x 10 "
		Hydrogen cyanide (HCN)	2.8 + ! vx 10-7	6.0t^xl0 7	2.8! 14 x 10"7	2.2+ j ,x 10
		Carbon dioxide (CO:)	l.ltîlfxlO "	i.r l",x 10 *	1.0'S 7 x 10 *	1.5! iî y 10 *
		Cyanogcn (C>N;)	<2 x 10"	-	-	-
		Cvano acctylcnc (HC,N)	<3* 10"	-	-	-
		Monodeuteratcd mcthanc	<2x10"	-	-	-
		(CH,D)
	4.1.2	Infrared Space Observatory (ISO)
	Almost two Titan seasons after the Voyager close fly-by, Titan's stratosphère was again
	observed by the Infrared Space Observatory (ISO) in 1997, close to northem autumnal
	equinox. The ISO spectrometers performed disk-averaged observations within the 7-30 pm
	région with a grating spectral resolution (>106), an order of magnitude higher than
	Voyager/IRIS.
		The nominal thermal profile retrieved from ISO is warmer than any of the IRIS-
	inferred thermal profiles at pressures lower than 1 mbar. The ISO higher resolution
	measurements helped to separate the spectroscopic signatures of C4H2 and C3H4 (at around
	630 cm'1) as well as of the HC3N and CO2 (contributing at around 660 cm'1), which were
	blended in the IRIS data. The abundances of the trace constituents retrieved from ISO spectra
	are listed in Table 4.3 below with the Voyager 1 équatorial abundances for comparison. The
	latter are the outcome of a recalculation in 2003 of the VI values from Coustenis et al. (1995)
	by using updated spectroscopic Unes for the contributing molécules. It should be noted that
	the VI values are latitude-depended and the ISO results are disk-average abundances.
	Additionally, the altitudes which are probed should be taken into considération as well as the
	fact that ISO data incorporate also a limb contribution. The équatorial VI retrievals are
	therefore the doser for an actual comparison with the ISO ones (Coustenis et al., 2003) as
	indicated at the Table 4.3.

Table 4 .

 4 3 -Trace gaseous abundances of Titan's stratosphère fromCoustenis et al. (2003), retrieved from ISO spectra taken in 1997. The Voyager équatorial values hâve been recalculated by adopting updated spectroscopic parameters.

	Molécule or	Abondance
	isotopic ratio	Voyager equator	ISO disk-average
	CA	2.85*1, X 10~6	55 ± 05 x 10 6
	CA	15t°0i X 10 "7	12 ± 03 X 10~7
	CA	x 10"s	2.0 ± 0.8 x 10~s
	ca	6.51% x 10-9	12 ± 0.4 x 10~8
	C3H4 (aliéné)		<2.0 x 10 9
	ÇA	7.0 ± 4.0 x 10-7	2.0 ± 1.0 x KT7
	CA	15 ± 0.7 x HT9	2.0 ± 05 X 10"9
	CA (benzene)		4.0 ± 3.0 x 10 10
	HCN	1.9S1H x 10"7	3.0 ± 05 x nr7
	CO,	145îgj| x 10"*	2.0 ±02x 10-8
	H20 at 400 km		8 x 10"9
	CHjD	I.1Î& X 10*	6.7!f9 x 10 #
	D/H	15tii x 10~4	8.7±fJ x 106
	HCjN	<1.0 x 10~9	5.0 ± 35 x 10-10
	ca	<1.0 x 10"9

Table 4

 4

		. 4 -Trace gaseous abundances of Titan's stratosphère from ground-based observations. To compare them with the Cassini data, we take into account only the
	observations pertaining to Titan's center and in the stratosphère, essentially at altitudes between 1 and 10 mbar (100-200 km roughly) depending on the molécule (see
	contribution functions and text for more details).
	Species	DATE OBSERV.	Ls	ABUNDANCE	(km) ALTITUDE	DE LATITU	INSTRUMENT	REFERENCE	POWER RESOLVING	" cm	-i	COMMENTS
		10/1/1997	194	4.92±0.43 x 10'6	stratosphère	disk average	ISO	Coustenis et al. 2003	1650-2000	>		reanalysis of Coustenis et al. 2003 values
				7.00±4.50 x 10'6	300							DATA: reanalysis of Geballe et al.
	c2h2									3225	2003 observations. The
		21/11/2001	257	5.60±3.30 x 10'6	200	disk	KECK 11	Kim et al. 2005	2.5 x 104			tempareture profile used in this
						average	NIRSPEC			3472	work is consistent with Yelle
												(1991) and Strobel et al. (1992),
				4.00±2.30 x 10*6	100							but not with Lellouch et al. (1990)
												The abundance has been adjusted
	c2h4	04/1974-04/1975	282-295	1.10±0.90 x 10'7	stratosphère	disk average	Kitt Peak National Observatory	Gillett, 1975 & Orton, 1992	2.0 x 10'2			lowered by a factor of 35% (see Coustenis et al. 2007; 2010) since the original value was derived by
												using old spectroscopic data
		10/1/1997	194	8.48±1.00 x 10'8	stratosphère	disk average	ISO	Coustenis et al. 2003	1650-2000			reanalysis of Coustenis et al. 2003 values
		10/1/1997	194	1.67±0.42 x 10'9	stratosphère	disk average	ISO	Coustenis et al. 2003	1650-2000			reanalysis of Coustenis et al. 2003 values
												v4 band. Reanalysis of Geballe et
	c4h2									3225	al. 2003 observations, T-profile:
		21/11/2001	257	< 1.00 x 10'7	stratosphère	disk average	KECK II NIRSPEC	Kim et al. 2005	2.5 x 104	3472	consistent with Yelle 1991 and Strobel et al. (1992) but not with
												Lellouch et al. (1990). Below 1
												mbar the VI /ISO profile adapted

Table 4 .

 4 5 -Characteristics of the CIRS nadir spectral sélections from Titan flybys TB-T44 at high resolution of 0.53 cm'1 (adapted from Table1of[START_REF] Coustenis | Titan trace gaseous composition from CIRS at the end of the Cassini-Huygens prime mission[END_REF].

	Mean		FP3			FP4
	Latitude
		Spectra	airmass	S/N	Spectra	airmass	S/N ratio
	70°N		1.21	34.4	2011	1.32	102.6
	60°N		1.25	48.0	2147	1.25	91
	50°N		1.12	48.7	4302	1.09	117
	42°N		1.08	84.5	4098	1.12	133.7
	33 °N		1.12	116.8	8222	1.14	242.6
	25 °N		1.09	140.3	7942	1.09	277.6
	15°N		1.07	114.9	9299	1.06	340.2
	5°N	11981	1.03	173.8	10899	1.04	383.7
	5°S	10873	1.04	161.7	10903	1.07	392.3
	15°S		1.15	92.6	3567	1.25	229.6
	25°S		1.13	62.9	3457	1.07	206.8
	33°S		1.13	85.0	2949	1.24	197.4
	42°S		1.24	65.7	3332	1.28	205.8
	50°S		1.15	60.7	1968	1.16	149.4
	60°S		1.28	60.8	1845	1.26	145.1
	70°S	778	1.45	39.1	2543	1.41	173.0

Table 4 .

 4 6 -Characteristics of the CIRS nadir spectral sélections from Titan flybys TB-T44 at high resolution at 2.54 cm'1 (adopted from Table1of[START_REF] Coustenis | Titan trace gaseous composition from CIRS at the end of the Cassini-Huygens prime mission[END_REF].

	Mean		FP3			FP4
	Latitude	Spectra	airmass	S/N	Spectra	airmass	S/N
	70°N	528	1.27	59.2	530	1.24	135.7
	60°N		1.21	84.9	1617	1.03	201.2
	50°N		1.11	203.9	6471	1.12	382.9
	42°N		1.13	227.3	7170	1.15	480.5
	33°N		1.12	337.2	11154	1.13	721.7
	25°N	12614	1.11	444.9	10317	1.13	861.1
	15°N	13018	1.07	479.3	14509	1.08	1123.8
	5°N	15639	1.09	534.8	15741	1.09	1231.8
	5°S	19953	1.05	575.5	19726	1.04	1357.3
	15°S	18953	1.05	524.3	19140	1.05	1338
	25°S		1.11	323	8266	1.10	861.2
	33°S		1.08	296.8	5791	1.12	706
	42°S		1.12	244.9	4722	1.12	613.1
	50°S		1.10	176.8	3035	1.12	475.4
	60°S		1.16	121.6	1270	1.16	313.3
	70°S		1.33	126.8	1421	1.28	330.2

Table 4 .

 4 7 -Titan flybys and FP3 data acquisition characteristics from March 2007 to September 2011 averaged in northem latitudinal bins (adapted from Bampasidis et al. 2012a).

				Total
				number of	Signal-		Cassini
	Year	Month	Latitude	spectra	to-noise	Airmass	Flyby	Ls (°)
		Mar	50°N	357	18.6	1.02	T26-T27	329
	II	Dec	50°N	335	19.9	1.25	T38-T39	338-339
		Feb-Mar-Jul	50°N	145	12.4	1.19	T41-T45	341-347
		Mar	50°N	517	25	1.35	T51	355
	II	Apr-May	50°N	275	15.3	1.01	T52-T55	355-357
	II	Jun	50°N	476	25.6	1.88	T56-T57	358
		J an	50°N	284	21.3	1.69	T65-T66	005-006
		Jul	30°N	551	34	1.17	T15-T17	320
		May	28°N	934	44.7	1.06	T30-T31	331-332
	tl	Dec	31°N	905	45.4	1.14	T38-T39	338-339
		Mar	30°N	992	42.3	1.01	T42	342
	It	Dec	36°N	262	20.5	1.03	T48-T49	351-352
		Dec	30°N	748	44.5	1.46	T63-T64	004-005
		Jun	35°N	599	37.4	1.24	T69-T70	010
		Sep	32°N	771	46.8	1.49	T78	025
		Jan	1°N	1291	56.2	1.01	T10	314
	"	Jul	1°N	2118	71.8	1.01	T15-T16	320
	2007	Apr-May	1°N	189	23.9	1.21	T28-T29	330-331
	II	Jun	1°N	1497	62.5	1.06	T32-T33	332-333
	II	Aug	1°N	152	19.3	1.01	T35	335
	2008	Jan	1°N	589	39.4	1.07	T40	340
	II	May	2°N	597	47	1.55	T43-T44	344-345
	2009	Mar-Jul	1°N	86	18.3	1.75	T51-T59	355
	II	Oct	1°N	72	14.9	1.19	T62	002
	II	Dec	1°N	283	25.6	1.01	T63-T64	004
	2010	Sep	3°N	1504	58.9	1.02	T72	014
		May	11°N	921	50	1.20	T76	021

Table 4 .

 4 8 -Titan flybys and FP3 data acquisition characteristics from February 2006 to January 2012 averaged in Southern latitudinal bins (adapted by Bampasidis et al. 2012a).

				Total number	Signal-to-
	Year	Month	Latitude	of spectra	noise	Airmass	Cassini Flyby	ls n
	2006	Mar	1°S	1515	61.3	1.04	T12	316
	2007	Oct	1°S	93	17	1.28	T19-T20	323-324
	H	Jul	rs	232	23.3	1.01	T34	334
	»l	Nov	1°S	813	44.5	1.01	T37	338
	"	Dec	3°S	276	25.3	1.01	T38-T39	338-339
	2008	Feb	rs	426	39.8	1.46	T41	341
	2010	Apr-May	1°S	466	33.0	1.01	T67	008
	II	Jun	1°S	391	30.6	1.02	T69-T70	010
	lt	Jul	rs	985	49.8	1.08	T71	011
	2011	Dec	2°S	1047	52.1	1.19	T79	028
	2006	Feb	30°S	666	37.7	1.88	TU	315
	lt	May	30°S	536	35.3	1.33	T14	318
	2008	Nov	38°S	1055	39.9	1.02	T46-T47	350-351
	2009	Dec	24°S	961	40	1.08	T63-T64	004-005
	2010	May	24°S	911	44.9	1.30	T68	009
	2012	Jan	30°S	1980	56.1	1.49	T80-T81	029
	2006	Oct	50°S	568	30.9	1.18	T19-T20	323
	2007	Jan	50°S	925	37.0	1.04	T23-T24	327
	"	Mar-May	50°S	341	28.2	1.64	T26-T31	329-332
	"	Jul-Aug	50°S	647	39.6	1.70	T34-T35	334
	II	Dec	50°S	467	31.7	1.53	T38-T39	338-339
	2008	Jul	50°S	34	9.6	1.28	T41	341
	2009	Mar	50°S	198	17.2	1.09	T51	355
	II	May	50°S	1288	45.2	1.25	T54-T55	357
	2010	Apr	50°S	124	16.0	1.69	T67	008

Table 4 .

 4 9 -Titan flybys and FP4 data acquisition characteristics from March 2007 to September 2011 averaged in northem latitudinal bins (adapted by Bampasidis et al. 2012a).

				Total number	Signal-		Cassini
	Year	Month	Latitude	of spectra	to-noise	Airmass	Flyby	Ls O
	2007	Mar	50°N	1376	64.6	1.03	T26-T27	329
	II	Dec	50°N	342	34.4	1.26	T38-T39	338-339
	2008	Feb-Mar-Jul	50°N	166	12.8	1.14	T41-T45	341-347
	2009	Mar	50°N	417	30	1.28	T51	355
	"	Apr-May	50°N	1481	49.9	1.01	T52-T55	355-357
	It	Jun	50°N	575	35.8	1.65	T56-T57	358
	2010	Jan	50°N	234	27.6	1.67	T65-T66	005-006
	2006	Jul	30°N	567	67.1	1.17	T15-T16	320
	2007	May	30°N	1789	125	1.06	T30-T31	331-332
	"	Dec	30°N	657	81.8	1.18	T38-T39	338-339
	2008	Mar	30°N	1013	86.5	1.02	T42	342
	"	Dec	35°N	394	74.1	1.03	T48-T49	351-352
	2009	Dec	30°N	764	83.9	1.36	T63-T64	004-005
	2010	Jun	36°N	503	69.6	1.25	T69-T70	010
	2011	Sep	28°N	1107	106	1.39	T78	025
	2006	Jan	2°N	869	107.3	1.01	T10	314
	fl	Jul	1°N	1840	152.0	1.01	T15-T16	320
	2007	Apr-May	3°N	93	37.5	1.20	T28-T29	330-331
	II	Jul	1°N	198	51.9	1.02	T34	334
	fl	Aug	1°N	146	45.5	1.02	T35	335
	2008	Jan	1°N	449	80.3	1.09	T40	340
	tl	Feb	1°N	439	86.2	1.59	T41	341
	II	May	3°N	64	32.9	1.70	T43-T44	344-345
	2009	Dec	1°N	322	64.8	1.03	T63-T64	004
	2010	Apr-May	1°N	233	54.1	1.01	T67	008
	2011	May	7°N	765	99.8	1.18	T76	021

Table 4 .

 4 10 -Titan flybys and FP4 data acquisition characteristics from February 2006 to January 2012 averaged in Southern latitudinal bins (adapted by Bampasidis et al. 2012a).

				Total number	Signal-		Cassini
	Year	Month	Latitude	of spectra	to-noise	Airmass	Flyby	Ls O
	2006	Mar	1°S	1515	142.1	1.09	T12	316
	II	Oct	rs	96	39.4	1.26	T19-T20	323-324
	2007	Jun	1°S	2545	195.1	1.09	T32-T33	332-333
	II	Nov	1°S	869	106.7	1.01	T37	338
	II	Dec	rs	382	70.9	1.02	T38-T39	338-339
	2009	Mar-Jul	1°S	84	39.2	1.71	T51-T59	355
	II	Oct	1°S	60	26.6	1.09	T62	002
	2010	Jun	rs	367	69.4	1.01	T69-T70	010
	II	Jul	rs	913	110.8	1.13	T71	011
	"	Sep	3°S	1622	136.9	1.01	T72	014
	2011	Dec	10°S	1866	157	1.24	T79	028
	2006	Feb	30°S	666	96	1.21	Tl 1	315
	II	May	30°S	551	85.5	1.26	T14	318
	2008	Nov	31°S	1055	106.4	1.01	T46-T47	350-351
	2009	Dec	28°S	980	103.4	1.12	T63-T64	004-005
	2010	May	24°S	1028	114.3	1.21	T68	009
	2012	Jan	30°S	2083	133.8	1.15	T80-T81	029
	2006	Oct	50°S	546	78.8	1.16	T19-T20	323
	2007	Jan	50°S	842	95.8	1.07	T23-T24	327
	II	Mar-May	50°S	320	63.1	1.42	T26-T31	329
	II	Jul-Aug	50°S	538	86.8	1.68	T34-T35	334
	II	Dec	50°S	456	77.5	1.59	T38-T39	338-339
	2008	Jul	50°S	51	25.6	1.62	T41	341
	2009	Mar	50°S	584	73.4	1.13	T51	355
	II	May	50°S	892	91.0	1.33	T54-T55	357
	2010	Apr	50°S	525	35.4	1.71	T67	008

Table 4 .

 4 11 -Titan flybys and medium FP3 resolution data acquisition characteristics from January 2008 to June 2010 averaged in 50°N, 0°and 50°S at medium resolution (2.53 cm'1).

				Total number	Signal-		Cassini
	Year	Month	Latitude	of spectra	to-noise	Airmass	Flyby	Ls(°)
	2008	Jan	50°N	459	60.2	1.26	T40	340
	2009	Apr	50°N	1193	87.1	1.08	T52-T53	355-356
	2010	May	50°N	160	46.7	1.84	T68	009
	2010	May	0°N	396	87.5	1.12	T68	009
	2010	May-Jun	50°S	484	77.1	1.74	T68-T70	009-010
	Table 4. 12 -Titan flybys and medium FP4 resolution data acquisition characteristics from January 2008 to June
	2010 averaged in 50°N, 0°and 50°S at medium resolution (2.53 cm"1).
				Total number	Signal-		Cassini
	Year	Month	Latitude	of spectra	to-noise	Airmass	Flyby	Ls(°)
	2008	Jan	50°N	495	106.9	1.27	T40	340
	2009	Apr	50°N	4447	227.0	1.01	T52-T53	355-356
	2010	May	50°N	163	67.2	1.84	T68	009
	2010	May	0°N	390	192.7	1.13	T68	009
	2010	May-Jun	50°S	473	173.9	1.78	T68-T70	009-010

Table 4 .

 4 13 -Titan flybys and medium FP1 resolution data acquisition characteristics from December 2004 to

	December 2008.
		Total number		Cassini
	Latitude	of spectra	Airmass	Flyby
	80°S-45°S	1700	1.22	T40
	45°S-10°S	3800	1.27	T52-T53
	0-30°N	7000	1.21	T68

Table 4 .

 4

	c4h2	628
	c3h4	633
	13cl8o16o	643
	13c16o2	648.5
	h13cI2c2n	658.7
	hc3n	663.2
	co2	667.4
	c6h6	674
	c2hd	678
	h13cn	706
	hcI5n	711.5
	HCN	712.7
	13c12ch2	728.5
	c2h2	729.5
	c3h8	748
	c2h6	822
	c13c12h6	822
	c2h4	949.5

14 -List of the currently detected molécules and isotopologues in Titan's atmosphère in the FP3 région. molécule cm'1

Table 4 .

 4 15 -Candidate molécules that are expected to be detected in CIRS FP3 and FP4 spectra

	molécule	cm'
	DCN	554
	13c17o16o	646
	12c18o2	657
	12c18o17o	660
	l2c18o16c	662.5
	hcc13cn	663.2
	hc13ccn	663.2
	12c17o16o	664.9

  the model underestimates the INMS value which is suggestive of the contribution of dynamical effects that are not taken into account in the 1-D model. Overall, the model is quite consistent with the INMS data, although it présents a small underestimation of the CfH(, INMS abundance at some northern latitudes. The next step is to include in our analysis the INMS measurements for the whole latitude range, and we hâve been in discussions over this with Prof. H. Waite, PI ofINMS.

							53.21N
							-47.92N
							43.12N
		Density Vertical Distribution	38.165N	Mixing Ratio Vertical Distribution
	1400						23.21N	1400-
							17.761 N
							2.54N
	1200							1200
							-27S
								1000
								800
								600-
								h
								400
	100	1000	104	105	106	107	108	10'°10-9	108	10'7	10'6
			DENSITY (cm-1)			MIXING RATIO

Table 4 .

 4  averages over 2009 at high resolution at 50°N.

		Total number	Cassini
	Focal Plane	of spectra	Airmass

  model, while butane and acetylene can be exposed as sédiments. It should be noted however, that Cordier et al. model does not take into account the temporal variation of lakes' Chemical composition.

Table 5 .

 5 1 -Lakes15 on Titan's surface with mean diameter greater than 200 km.

	Diameter

Figure 5.3-Titan's North Pôle projection from Cassini/SAR, showing the extent of Titan's large north polar lakes. Jingpo Mare is the large unlabeled long narrow feature eastwards of Kraken Mare (crédit: Turtle et al., 2009). 15 http://planetarynames.wr.usgs.gov/ 5.4 Numerical analysis in Cassini/SAR images 5.4.1

The context

Table 5 .

 5 2-Number of Streams formed the 3rd order basins of the island and Bifurcation Ratio

			Stream Number	Bifurcation
	3rd order
	Basin	T1	2nd	3rd	TOTAL	Ratio
		N,	n2	n3	SNU	Rb
	I	12	4	1	17	3.00
	II	13	6	1	20	2.17
	III	8	2	1	11	4.00
	IV	8	2	1	11	4.00
	V	11	2	1	14	5.50
	TOTAL	52	16	5	73

Table 5 .

 5 3-Lengths of the streams formed the 3rd order basins of the island.

	3rd	Stream Length
	order
	Basin	lst Order 2nd Order	3rd Order	TOTAL
		SL, (m)	EL2 (m)	£L3 (m)	ELU (m)
	I	38,425.3	49,287.3	25,733.8	113,446.5
	II	76,962.7	165,799.5	45,986.8	288,749.0
	III	32,136.7	58,416.1	6,799.0	97,351.8
	IV	17,857.3	49,628.5	8,456.5	75,942.3
	V	48,650.1	36,025.7	18,827.3	103,503.1
	TOTAL 214,032.2 359,157.2	105,803.4 678,992.8

Table 5 .

 5 4 below, contains the number of the streams of each order of the rest of the island,

Table 5 .

 5 4-Streams of the island, which they do not belong to a 3rd order basin.

		North 2nd	South 2nd
				Total Number	% of the
	Stream Order	order	order	of the streams of the island	Total number of streams
		Basins	Basins
	1	37	37	126	59
	2	9	12	37	57
	3	0	0	5	-
	Total	46	49	168	57

Table 5 .

 5 5-Average length of the streams formed the 3rd order basins of Mayda Insula.

	3rd order	Average Stream Length
	Basin	1st Order	2nd Order	3rd Order
		S<L]> (m)	£<L2> (m)	£<L3> (m)
	I	3,202.1	12,321.8	25,733.8
	II	5,920.2	27,633.3	45,986.8
	III	4,017.1	29,208.0	6,799.0
	IV	2,232.2	24,814.3	8,456.5
	V	4,422.7	18,012.9	18,827.3

Table 5 .

 5 6-Stream Length Ratio Area and Drainage Density of the 3rd order basins of the island.

	Stream	Stream length	Area of	Drainage
	3rd	Ratio	Basin	Density
	order Basin	RL(2,1)Rl(3,2)	A (xlO3 km2)	D (km'1)
	I	3.8	2.1	1.35	0.08
	II	4.7	1.7	2.96	0.10
	III	7.3	0.2	0.76	0.13
	IV	11.1	0.3	0.72	0.11
	V	4.1	1.0	1.05	0.10
	TOTAL			6.84

Table 5 . 7

 57

	Angles (deg) of the main Mountain Ridges of
				Région A -Island
	41.7	-5.6	6.6	-7.3	30.6	2.7	-47.8 71.4	19.4
	39.5	-6.1 -12.2 -19.6 -10.0 42.9 -42.5 -2.2	22.8
	41.4 20.5 -16.4 38.6 -30.1 39.0 39.9 -14.2 16.6
	14.8	-8.7	26.4 27.3 -54.4 -0.5	59.4 -21.5 -25.9
	-27.2 41.0 -11.7 46.3	52.1 -10.1	-0.1	34.8 21.1
	12.1 -19.7 -19.1 -19.1 -30.9 -47.4 -18.7 -3.2	62.1
	44.6	-6.5	15.9 -22.4 -23.9 -3.7

-Angles of mountain ridges according the N-S direction, Région A -Mayda Insula. The positive angles are looking to the East and the négatives to the West.

Table 5 .

 5 8 -Angles of mountain ridges according the N-S direction, Région B -Western Mountains. The positive angles are looking to the East and the négatives to the West.

	Angles (deg) of the main Mountain Ridges of
		Région B -Western Mountains
	20.8	13.7 -20.6 -5.8	20.8 53.1 -47.5 31.6 -30.5
	7.5	30.7 43.7 -27.1	8.1	-19.0 -34.8 -5.2 -8.7
	65.3 32.8 53.3 -10.6 70.2 47.0	14.5 25.7 24.4
	30.4 -45.0 -2.7	36.2 -54.7 -2.9 -24.0 51.3 28.7
	-58.2 -37.9 -27.0 12.8	38.3	0.0	-22.1 25.0 20.8
	43.9 72.2 35.8 32.6	5.3	-17.3 -2.3

Table 5 .

 5 9 -Angles of mountain ridges according the N-S direction, Région C -SW Mountains. The positive angles are looking to the East and the négatives to the West.

	Angles (deg) of the main Mountain
			Ridges of
		Région C-SW Mountains
	-3.5	-39.9	30.3	16.7	0.0
	16.1	8.3	47.1	12.8	16.7
	51.4	54.1	13.8	0.0	22.1
	40.9	-9.8	-57.1	11.0	-46.0
	33.3	-39.9	44.1	35.3	-33.7
	9.5

Table 5 .

 5 10 -2nd Order Stream Angles Orientation according to N-S direction. The positive angles are looking to the East and the négatives to the West.

	2nd Order Stream Angles (deg) to N-S
			direction
	83.7	-67.1	-56.7	-72.5
	55.7	67.9	-39.8	-72.5
	22.9	-42.1	4.5	-88.2
	27.3	-72.3	11.3	31.9
	55.6	-59.6	61.0	-26.7
	Histogram of 2nd Order Streams orientation

  . RADAR opérâtes only during close flybys of Titan or other planetary objects of interest. Figure5.13 illustrâtes the observational sequence during nominal Titan flybys, when the closest approach (C/A) is at an altitude of 1000 km or lower.

	300	78		30			16	6	0
	min	min	min	min	min	min min
	Titan Disc Scan	Calibration		Titan Disc Scan	Nadir Pointing	Side Looking	REVERSE
	Radiometry Only (S/C rotate)	L-R Altimetry (Scatterometry), Radiometry	H-R Altimetry, Radiometry	Imaging, Radiometry	SEQUENCE
	100000 km
				241

  TSPR filter application in a SAR image we needed a 32-bit version of the image. Ail SAR data of the Cassini-Huygens mission are Using the Integrated Software for Imagers and Spectrometers (ISIS)17 that US Geological Survey has developed, Cassini/SAR images can be retrieved. The software packages, the installation procedure and an assistance forum are available online as well. ISIS inputs the JPL Cassini RADAR science data products and extracts 32-bit images which can be processed by GIS software applications and other numerical computing software.A 32-bit version of Cassini/SAR image PIA0863018 was provided by Dr. A. Hayes(Fig. 5.14). We hâve chosen a SAR image with lakes for a test case, since the contrast between the lakes and the surroundings is significant. This image was taken during the Cassini Tl6 Titan flyby onJuly 21, 2006. It is centered near 80°N, 92°W and its dimensions

	available in NASA Planetary Data System (PDS) except for the most recent ones, which are
	in raw format. NASA/PDS database contains a pair of Short Burst Data Record (SBDR) and
	Long Burst Data Record (LBDR) files for each flyby. The only différence between the

datasets formats is whether or not two data fields are included: the sampled écho data and the altimeter profile. The altimeter profile is an intermediate processing resuit between sampled écho data and a final altitude estimate. LBDRs include the écho data but not the altimeter profile16.

are 420 km horizontally and 150 km vertically. The dark patches are lakes, filled with organics. This data is an 860x6200 array and is in a band sequential ISIS 3 format with an oblique cylindrical projection. The image covers an area 750x3100 in pixel size and the pixel scale is set to 175.558 meters per pixel. The actual SAR resolution is around 350 meters per pixel and the image contains some interpolation.

  , 1.4% in the stratosphère and 5% on the surface(Niemann et al., 2010). Interactions among Satum's magnetosphere and Titan's ionosphère and atmosphère, involving solar EUV and UV radiation, energetic ions and électrons, energetic particles ffom the Satumian magnetosphere and galactic cosmic rays, photo-dissociate both nitrogen and methane and trigger a complex photochemistry.

The Cassini Ion and Neutral Mass Spectrometer (INMS) is extensively studying in situ Titan's upper atmosphère during low altitude Cassini flybys. INMS showed the presence of many organic species at détectable levels, at very high altitudes (1100 -1300 km) establishing that this active photochemistry produces complex organic species and nitriles. Extrapolation among the INMS measurements (limited to mass up to 100 Daltons) and Cassini Cosmic

  long-term remote and in situ observations of these icy moons, comparative studies operate supplementary to our perspective of these worlds. In our paper,[START_REF] Solomonidou | Water Océans of Europa and Other Moons: Implications For Life in Other Solar Systems[END_REF], we study the surface aspects of Jupiter's Ganymede and Europa and Satum's Titan and Enceladus and try to find which of their similarities are connected with astrobiological implications. In this section, I hâve expanded the comparison among the Satumian moons Titan and Enceladus and the Jovial Ganymede and Europa by considering

	also Callisto.

Callisto

Since we lack

6.4

.1 Ganymede, the biggest satellite of the Solar System Ganymede is the biggest satellite in our Solar System with a diameter of 5,272 km, a mass of 148.2 x 1021 kg and a density of 1940 kg/m3. It orbits Jupiter at 1.07 million km and its surface has a mean température of about 110 K. It is the only moon that possesses a magnetosphere

  Europa Jupiter System Mission) was a joint NASA and ESA mission study, which would focus on the exploration of Jupiter and its moons. It was a L-scale (large) mission consisting of two orbiters, the Jupiter Ganymede Orbiter (JGO) and the Jupiter Europa Orbiter (JEO). JGO would be developed by ESA and it was a continuation of the LAPLACE proposai[START_REF] Blanc | LAPLACE: A mission to Europa and the Jupiter System for ESA's Cosmic Vision Programme[END_REF], submitted in response to ESA's Cosmic Vision 2015-2025 Call mission {EJSM, NASA/ESA Final Report, 30 January 2009). It would focus mainly on Ganymede and Callisto and the Jupiter System. JEO, on the other hand, would hâve been developed by NASA and it would focus mainly on Europa and explore its habitability. Both orbiters'

	7.1.1	The Europa Jupiter System Mission (EJSM)
	EJSM (

Table 7

 7

		Radiator
		Thruster
	3m Articulated HGA	Cluster (8)
	Telecom
	Electronics
		Electronics
	Ox	Châssis (2)
	Tank
	Electronics Châssis
	Magnetometer	Star Tracker (2)
	Boom

.1 below describes the main payload proposed for this mission {EJSM, NASA/ESA Final Report, 30 January 2009). 23 http://juno.wisc.edu/juno-mission.html; http://missionjuno.swri.edu/ Figure 7. 1 -NASA's Jupiter Europa Orbiter (crédit: NASA/EJSM Final Report) Figure 7.2-ESA's Jupiter Ganymede Orbiter (crédit: ESA)

Table 7 .

 7 1 EJSM payload (EJSM Final Report). However, due to budget cuts, NASA could not support in the immédiate future the development of such a large-scale mission as JEO and terminated that study. ESA has decided to continue with the same objectives and some enhancement with respect to the JGO, but with a unique spacecraft dedicated to study the icy moons of Jupiter : Ganymede, Callisto and Europa. In this context, the ESA Jupiter Icy Moons Explorer mission (JUICE) was proposed as a reformulation of JGO. The mission has won the compétition for the first L-class mission of ESA's Cosmic Vision 2015-2025 Program. It has been selected for the L1 slot of ESA's Cosmic Vision science program on May 2, 2012. Two other proposais were competing:

	Orbiters				Science Payload
			Instrumentation	Main objectives
	JGO/JEO	Laser Altimeter		Surface topography -cohérence with gravity, tides
	JGO/JEO	Radio Science		Interior State of moons, presence of a deep océan and other gravity
					anomalies
	T_,^.TT,_	t	_	_ ,	Structure of the subsurface & identify warm ice and/or anomalies within the
	JGO/JEO	Ice Penetrating Radar	.
	JGO/JEO	Visible-IR		Composition of non-ice components on Ganymede & Callisto;
		Spectrometer		State & crystalinity of surface ices
		Ultraviolet
	JGO/JEO				Composition & dynamics of the atmosphères of Ganymede & Callisto
		Spectrometer
		Ion and Neutral Mass
	JEO	Spectrometer		Composition of sputtered products from Europa
	JGO/JEO	Thermal Instrument	Map température anomalies and thermal inertia of surface materials on
					Ganymede
	JEO	Narrow Angle		Local-scale géologie processes on Europa, Ganymede & Callisto; Io volcano
		Caméra		monitoring; Jupiter cloud dynamics & structure
	JGO/JEO	Wide and Medium	Global morphology of Ganymede; Global to régional scale morphology of
		Angle Caméra
	JGO/JEO	Magnetometer
	JGO/JEO	Plasma and Particles
		Submillimeter Wave
	JGO
		Sounder

Callisto

Ganymede's intrinsic magnetic field and its interaction with the Jovian field Interaction between Ganymede & Callisto and the space environment to constrain induction responses Dynamics of Jupiter's' stratosphère; Vertical profiles of wind speed and température EJSM was in compétition with a mission concept for the exploration of the Kronian System which focused on Titan (the Titan Satum System Mission or TSSM) for a slot among the Lclass Cosmic Vision Program missions. After the initial studies, the Europa Jupiter System Mission (EJSM) mission was prioritized by NASA and ESA to be launched first in order to explore the Jupiter System. The Titan Satum System Mission (TSSM) mission, which will be described later in this Chapter, was to be developed and launched later for the Satumian System.

Table 7 .

 7 5 -Comparison between TSSM Montgolfière and TAE.

		TSSM Montgolfière	TAE balloon
	Gas	Titan gas	Hélium
	Diameter	10.5 m	4.6 m
		Multi-Mission Radioisotope	Advanced Stirling Radioisotop
	Power
		Thermoelectrical Generator (MMRTG)	Generator (ASRGs)
	Scheduled mission duration	6 months	3-6 months
	Low altitude orbit	10km	8 km
	Communication link	Orbiter intervation	Direct
			No magnetometer
			Imaging spectrometer probes from
	Payload	see Table 7.3
			4.6 to 5.6 pm
			(TSSM balloon: 1 -5.6 pm)
	Total data volume	up to 1.3 Tb	up to 1 Gb

Table 7 .

 7 6 -Main constituents of lakes on Titan and the hydrogen bonding strength of each one.

		m.r. pôles
	Main Solvent	(90 K)	H-bond
		(Cordier et al., 2009; 2012)
	c2h6	7.64x10''	poor
	n2	4.90xl0'3	no
	ch4	9.69xl0'2	poor
	c3h8	7.42x10'2	poor
	Main Solutés
	HCN	2.09x10'2	strong
	C4H10	1.21xl0'2	poor
	c2h2	1.15xl0'2	poor
	CH3CN	9.89xl0'4	poor

Table 7 .

 7 7-The requirements of the seismic instrument

			Seismic Sensors
	Power	Mass	Velocity	Operational
			Bandwidth	température
	0.32 W	0.200g

Table 7 .

 7 8 -Engineering constraints and conditions for a future space seismic experiment on icy moons MEMS technology can be implemented in future seismic experiments on icy moons[START_REF] Bampasidis | Evolution of minor trace gases and isotopic ratios in Titan's stratosphère using CIRS/Cassini spectra[END_REF]. A MEMS device is in experimental State in the Arizona State University, which is in collaboration with the MET Tech (USA) and the Moscow Institute of Physics and Technology (Russia).

	Constraints and	Possible Problems and Difficulties
	Conditions	Science	Engineering	Précautions
	Fluctuations in		Malfunction of the
		Accuracy affected		Spécial thermal shield
	température		sensors
				Spécial installation
			Loose the equilibrium
	Surface characteristics	Sensors' deformation		and stabilization
			position
				needed
			Insufficient solar power
	Dense atmosphère	Chemical Corrosion		Shield case
			to recharge batteries
	Stability of the	First time in space		Robotic installation of
			Orientation lost
	apparatus	exploration		the seismometer
	Maintenance	No recordings	Impossible to repair	Autonomous System
				Orbiter and large
		Radio frequency	A permanent link
	Data transmission			ground
		signais	needed
				radiotélescopes
				Seismic network
	Exact recordings	3D plotting	Different devices
				needed
		Low Signal to Noise		A shallow hole
	Ground/wind noise		Stable structure
		ratio		needed for installation
				Radioisotope
	Power	Continuous operation	Batteries	thermoelectric
				Generator (RTG)

Table 1 .

 1 9 -The Arizona State University MEMS device spécifications(Hongyu Yu, pers. comm)

	Spécification	Value
	Sensor Type	Molecular Electronic Transducers
	Proof Mass	None
	Size [mm]

Table 8 .

 8 1 -List of outreach events

				Number of
	Event	Location	Date
				visitors
		EUROPE
	Observation of the Venus transit
		Athens, Greece	6/6/2012	100
	Univ. of Athens
	Exhibition
		Evora,
	Antikythera Mechanism		19-23/09/2011	1000
		Portugal
	University of Evora
	Exhibition and Lecture
	Antikythera Mechanism and the Sun	Athens,
			02-05/09/2011	3000
	Stadium of Peace and Friendship	Greece
	2nd Festival for the child
	Lectures and Observation
		Salamis,
	of the starry sky		09/07/2011	300
		Greece
	Salamis island
	Exhibition and Lecture
	Antikythera Mechanism and	Heretria and
	Astrophysics for pupils	Gianitsa	26-28/06/2011	150
	Hellenic Physical Society	Greece
	Summer School
	Exhibition and Lecture
	Antikythera Mechanism
		Athens,	08/06/2011-
	Hellenic Foundation for Culture			2200
		Greece	30/09/2011
	Science and Technology of the Ancient
	World
	Exhibition and Lecture
		Agrinio,
	Antikythera Mechanism		16/05/2011	800
		Greece
	Papastrateion Centre
	Exhibition and Lecture
		Nicosia and
	Antikythera Mechanism		27/04/2011-
		Limassol		800
	European University Cyprus		14/05/2011
		Cyprus
	Cyprus Science Festival
	Exhibition and Lecture
	Antikythera Mechanism and	Naousa,
			11-13/03/2011	200
	Astrophysics for pupils	Greece
	Hellenic Physical Society
	Exhibition and Lecture	Portaria,
	Antikythera Mechanism and	Thessalia,	11-13/02/2011	200
	Astrophysics for pupils	Greece

Table 8 .

 8

	2 -Group formation
	Members in Groups	5	4	3
	Number of Groups	4	4	1

Table 8

 8

	.3-Main points of the tasks' worksheet
	Activity	Tasks

Table 8 .

 8 4 -Results of the online questionnaire

	Question	Very well	Well	Moderate	Little	Dislike
	General Satisfaction	49	42	9	0	0
	The concept	40	39	21	0	0
	Exploit Previous knowledge
		30	17	20	16	17
	of Science
	Instructions	37	49	14	0	0
	Interesting task	46	48	6	0	0
	Experiments helpful	57	17	14	9	3
	Teacher as a facilitator	94	3	0	3	0
	Ability to solve similar
		11	54	26	9	0
	problems in the future
	Satisfaction of their
	participation in the NASA	92	8	0	0	0
	contest
	Attend Astronomy lessons in
		42	32	19	4	3
	the future
	Astronomy issues in Science
		44	32	14	5	5
	disciplines
	Group coopération	43	26	14	14	3
	Ail the group member help
		32	20	31	6	11
	understanding
	No help from rest the
		38	14	11	17	20
	members

Table 2 Titan

 2 Flybys and FP3 Data Acquisition Cliaracteristks from 2006 February to 2012 Januaty Averaged in Soulhem latitudinal Bins

	TkE Astrûphvsic al Journal, 760:144 (8pp), 2012 Decembcr 1				Bampasidis et al.
				Table 1
	Titan Flybys and FP3 Data Acquisition Characterixtics frora 2007 Marc]) to 2011 Septeinber Averaged in Northern Latitudinal Bins
	Ycar	Month	Latitude	Total Nuraber of Spectra	Sig n al-to-noise	Airraass	Cassini Flyby	Ls(°)
	2007	Mar	50°N	357	18.6	1.02	T26-T27	329
	"	Dec	50"N	335	19.9	1.25	T38-T39	338-339
	2008	Feb-Mar-Jul	S0°N	145	12.4	1.19	T41-T45	341-347
	2009	Mar	50"N	517	25	1.35	T51	355
		Apr-May	50°N	275	15.3	1.01	T52-T55	355-357
		Jun	50°N	476	25.6	1.88	T56-T57	358
	2010	J an	50*N	284	21.3	1.69	T65-T66	005-006
	2006	Jul	30°N	551	34	1.17	T15-T17	320
	2007	May	28°N	934	44.7	1.06	T30-T31	331-332
		Dec	31°N	905	45.4	1.14	T38-T39	338-339
	2008	Mar	30°N	992	42.3	1.01	T42	342
		Dec	36°N	262	20.5	1.03	T48-T49	351-352
	2009	Dec	30°N	748	44.5	1.46	T63-T64	004-005
	2010	Jun	35°N	599	37.4	1.24	T69-T70	010
	2011	Sep	32°N	771	46.8	1.49	T7S	025
	2006	Jan	1<>N	1291	56.2	1.01	T10	314
		Jul	FN	2118	71.8	1.01	T15-T16	320
	2007	Apr-May	1°N	189	23.9	1.21	'128-129	330-331
		Jun	1#N	1497	62.5	1.06	T32-T33	332-333
		Aug	1°N	152	19.3	1.01	T35	335
	2008	Jan	1°N	589	39.4	1.07	T40	340
		May	2°N	597	47	1.55	T43-T44	344-345
	2009	Mar-Jul	1°N	86	18.3	1.75	T51-T59	355
		Oct	FN	72	14.9	1.19	T62	002
		Dec	FN	283	25.6	1.01	T63-T64	004
	2010	Sep	3»N	1504	58.9	1.02	122	014
	2011	May	11°N	921	50	1.20	T76	021
	Year	Month	Latitude	Total Numbcr of Spectra	Signal-to-noise	Aimiass	Cassini Flyby	Ls(°)
	2006	Mar	FS	1515	61.3	1.04	Tl 2	316
	2007	Oct	1°S	93	17	128	T19-T20	323-324
		Jul	FS	232	23.3	1.01	T34	334
	44	Nov	FS	813	44.5	1.01	T37	338
		Dec	3°S	276	25.3	1.01	T38-T39	338-339
	2008	Feb	FS	426	39.8	1.46	T41	341
	2010	Apr-May	FS	466	33.0	1.01	T67	008
		Jun	FS	391	30.6	1.02	T69-T70	010
		Jul	FS	985	49.8	1.08	T71	011
	2011	Dec	2°S	1047	52.1	1.19	T79	028
	2006	Feb	30'-S	666	37.7	1.88	Tll	315
		May	SO^S	536	35.3	1.33	T14	318
	2008	Nov	38°S	1055	39.9	1.02	T46-T47	350-351
	2009	Dec	24'S	961	40	1.08	T63-T64	004-005
	2010	May	24°S	911	44.9	1.30	T68	009
	2012	Jan	30°S	1980	56.1	1.49	T80-T81	029
	2006	Oct	50°S	568	30.9	1.18	T19-T20	323
	2007	Jan	50°S	925	37.0	1.04	T23-T24	327
	"	Mar-May	so^s	341	282	1.64	T26-T31	329-332
	"	Jul-Aug	50°S	647	39.6	1.70	T34-T35	334
	"	Dec	50°S	467	31.7	1.53	T38-T39	333-339
	2008	Jul	50°S	34	9.6	128	T41	311
	2009	Mar	50°S	198	172	1.09	T51	355
	"	May	50°S	1288	452	125	T54-T55	357
	2010	Apr	50°S	124	16.0	1.69	T67	008

Table 3

 3 Titan Flybys and FP4 Data Acquisition Characteristics from 2007 Match to 2011 September Aventged in Northern Latitudinal Bins

	Year	Month	Latitude	Total Number of Spectra	Signal-to-noise	Aimiass	Cassini Flyby	Ls(°)
		Mar	50°N	1376	64.6	1.03	T26-T27	329
		Dec	50"N	342	34.4	1.26	T38-T39	338-339
		Feb-Mar-Jul	50°N	166	12.8	1.14	T41-T45	341-347
		Mar	50°N	417	30	1.28	T51	355
		Apr-Mav	50°N	1481	49.9	1.01	T52-T55	355-357
		Jun	50°N	575	35.8	1.65	T56-T57	358
		Jan	50°N	234	27.6	1.67	T65-T66	005-006
		Jul	30°N	567	67.1	1.17	T15-T16	320
		May	30°N	1789	125	1.06	T30--T31	331-332
		Dec	30°N	657	81.8	1.18	T3S-T39	338-339
		Mar	30°N	1013	865	1.02	T42	342
		Dec	35°N	394	74.1	1.03	T48-T49	351-352
		Dec	30°N	764	83.9	1.36	T63-T64	004-005
		Jun	36°N	503	69.6	1.25	T69-T70	010
		Sep	28°N	1107	106	1.39	T78	025
		Jan	2°N	869	1073	1.01	T10	314
		Jul	1°N	1840	152.0	1.01	T15-T16	320
		Apr-May	3°N	93	375	1.20	T28-T29	330-331
		Jul	1°N	198	51.9	1.02	T34	334
		Aug	1°N	146	45.5	1.02	T35	335
		Jan	TN	449	80.3	1.09	T40	340
		Feb	1°N	439	862	1.59	T4I	341
		May	3°N	64	32.9	1.70	T43-T44	344-345
		Dec	1°N	322	64.8	1.03	T63-T64	004
		Apr-May	1°N	233	54.1	1.01	T67	008
		May	7°N	765	99.8	1.18	T76	021

Table 4 Titan

 4 Flybys and FP4 Data Acquisition Characteristics from 2006 Februaiy to 2012 January Averaged in Southern Latitudinal Bins

	Year	Month	Latitude	Tbtal Number of Spectra	Signal to-noise	Airmass	Cassini Flyby	U o
		Mar	1°S	1515	142.1	1.09	T12	316
		Oct	1°S	96	39.4	1.26	T19-T20	323-324
		Jun	rs	2545	195.1	1.09	T32-T33	332-333
		Nov	1°S	869	106.7	1.01	'07	338
		Dec	rs	382	70.9	1.02	T38-T39	338-339
		Mar-Jul	1°S	84	39.2	1.71	T51-T59	355
		Oct	rs	60	26.6	1.09	T62	002
		Jun	rs	367	69.4	1.01	T69-T70	010
		Jul	rs	913	110.8	1.13	T71	011
		Sep	3°S	1622	136.9	1.01	T72	014
		Dec	10*S	1866	157	124	T79	028
		Feb	30°S	666	96	121	Tl 1	315
		May	30°S	551	85.5	126	T14	318
		Nov	31'S	1055	106.4	1.01	T46-T47	350-351
		Dec	28°S	980	103.4	1.12	T63-T64	004-005
		May	24^8	1028	114.3	121	T68	009
		Jan	30°S	2083	133.8	1.15	T80-T81	029
		Oct	50"S	546	78.8	1.16	T19-T20	323
		Jan	50°S	842	95.8	1.07	723-124	327
		Mar-May	50"S	320	63.1	1.42	T26-T31	329
		Jul-Aug	50°S	538	86.8	1.68	T34-T35	334
		Dec	50°S	456	77.5	1.59	T38-T39	338-339
		Jul	50°S	51	25.6	1.62	T41	341
		Mar	50" S	584	73.4	1.13	751	355
	"	May	50°S	892	91.0	133	T54-T55	357
		Apr	50°S	525	35.4	1.71	T67	008

  ,*50?3.jQ g.

									2006	2007	2008	2009	2010	2011	2012
														Year
									w.				ISfSTT	I l"TT"T-| -<d)
									r	_	^r'c ri	'} l-, f ' T 3CTN
									-					À -. ï -
										11 rT :		1 "'f fcrftc	ci '
									l , i , , i , . ; . I i , , ,*i ,	. t / py, y	. 30 fS .
			2007	2008	2009	2010	2011	2012	2006	2007	2008	2009	2010	2011	2012
					Year									Year
	2006	2007	2008	2009	2010	2011		2006	2007		2008	2009	2010	2011	2012
					Year									Year
	§ 0.3	.-JL, i%$TO oV	(9) ==Hl| { 30ÜN 'L iilcja 3trs	S 1.5 "					'.so^tvi ' h -1 f Ot t i A IMU i '1 . 50°$ .	' ' ' I ' 30r rr fi	N	(h) = ; 30°S-
	2006	2007	2008	2009	2010	2011	2012
					Year

Table 1

 1

  Fig.12a. Comparison of the results from this work at 33"N for major hydrocarbons (points in green boxes) with the vertical profiles inferred from lirabdata (bluelines. 15'S,

				A Coustenis et aL/lcams 307 (3010) -161-476			47.3
	Molécule	Mixing ratios
		70°S	50°S	33*S	15*S	5*N	33"N	50°N	70*'N
	Mixing ratios of trace gnses in Titan's stratosphère
	CjHj	2.45E-06	2 35E-06	3.05E-06	2.97E-06	3.25E-06	330E-06	3.10E-06	5.10E-06
	c2hd	1.10E-09	7.00E-10	1.40E-09	7.50E-10	7.40E-10	1.00E-09	1.25E-09	2.00E-09
	c2h"	1.38 E-07	1.05E-07	1.40E-07	1.20E-07	1.35E-07	1.16E-07	2.05E-07	5.50E-07
	CaH6	6.07E-06	6.50E-06	6.65E-06	730E-06	8.00E-06	8.80E-06	8.50E-06	1.15E-05
	c3h4	3.98E-09	4.00E-09	4.20E-09	4.80E-09	5.90E-09	7.10E-09	1.23E-08	2.40E-08
	c3h.	4.45E-07	3.50E-07	5.20E-07	4.50E-07	5.00E-07	6.00E-07	6.30E-07	6.95E-07
	C"Ha	7.80E-10	9.00E-10	1.04E-09	1.12E-09	1.32E-09	1.77E-09	5.60E-09	2.30E-08
	CsH*	<1.00E-10	<1.00E-10	1.20E-10	2.20E-10	3.00E-10	1.90E-10	6.00E-10	420E-09
	HCN	2.95 E-08	3.60E-08	3.85E-08	6.70E-08	1.13E-07	2.25E-07	3.62E-07	9.70E-07
	HC3N	3.10E-10	4.20E-10	2.90E-10	2.80E-10	5.00E-10	4.50E-10	5.10E-09	4.50E-08
	COj	9.90E-09	1.10E-08	1.08E-08	1.10E-08	1.11E-08	1.15E-08	1.25E-08	1.40E-08

Table 1

 1

  general, image segmentation is a classification problem. A classification problem can be formalized as a pair (0,0), where O dénotés a set of objects and C a collection of disjoint subsets Ci Ci that partitions O. The problem is to détermine the subset CjCC to which a given object oeO belongs. The supervised method of minimum Euclidean distance uses the mean values (or vectors) of each mernber and calculâtes the Euclidean distance from each dassified object to the nearest class segmenting the image into different régions of interest or different labels (Richards, 1999). In our case we choose chree régions of interest: dark lakes (black label), granular lakes (dark gray label) and local background (light gray label). The TSPR filter can be used as an intermediate step before the extraction of the régions of interest.

As Euclidean distance Itéré, we define the distance de =

  . In the case of Mars, two major régions are known to display morphotectonic features; the Tharsis volcanic plateau, which was possibly formed after crustal deformation in Please cite this article as; Solomonidou, A., et al., Morphotectonic features on Titan and their possible origin. Planetary and Space Science (2012), http://dx.doi.Org/10.1016/j.pss.2012.05.003 A. Solomonidou et ai/ Planetary and Space Science l finiJ mm

Table 1

 1 Comparison between silicate and ice properties.

	Planetary and Space

Table 2b

 2b

Table 3

 3 Please cite this article as: Solomonidou, A., et al., Morphotectonic features on Titan and their possible origin. Planetary and Space

	suinmarizes tlieir major
	observations by Cassini.

A.

Solomonidou et aL [ Planetary and Space Science i (Mil} m m

Table 3

 3 Major fault, fracture and canyon formations on Titan.

	Location	Chara cter lza tlon	Proposed mechanism	FJyby/Tlme	Instrument
					(reference)
	15*S, 155"W	Conjugate-like faults	i. Large scale tectonic modification of bedrock material.

Table 4

 4

  proposed oscilla tions in the thermal State of the ice-I shell of the Satumian satellites,

	which may cause repeated épisodes of extensional and compressional
	tectonLsm. Similarly on Earth, a current widely accepted internai
	évolution model suggests that expansion and contraction processes
	are due to internai thermal runaway cycles and can be important in
	controlling geotectonic mechanisms (Kice and Fairbridge, 1975;

  A Solomonidou et ai / Planetary and Space Science fini,) m m comments that hâve immensely helped in improving the original manuscript. MH and AC acknowledge financial support from the

	French Agence Nationale de la Recherche (ANR Project: CH4®Ti-
	tan, PI: Athéna Coustenis).
	Anezina Solomonidou lias been co-financed by the European
	Union (European Social Fund-ESF) and Greek national funds
	through the Operational Program "Education and Lifelong Learn-
	ing" of the National Strategie Reference Framework (NSRF)-
	Research Funding Program: Heradeitus II. Investing in knowledge
	society through the European Social Fund.

  Gcorgios Bampasidis, Konstantinos Kyriakopoulos, Emmanuel Bratsolis. Mathieu Hirtzig.Aihena Coustenis 4Xenophon Moussas nents endorse théories suggesting that even though Titan is

	258	Anezma Solomonidou,
		Titan is tlte second largest inoon in the Solar System after
		Jupiter's Ganymede, with a radius of 2,575 km (Lîndal et
		al., 1983) and spherical geometry. Titan has a unique atmos
		phère, in that it is dense and consists mainly of N2 (98.4%),
		as on Earth. CH4 (1.4%), H2 (0.1%) and traces of argon,
		ethane, acetylene, propane and more complex hydrocarbons
		and nitriles, as well as condensâtes and organic aérosols
		(Coustenis & Taylor, 2008) constitute the rest of tlte at
		mosphère. The identification of such atmospheric compo-

  Geoi^ios Bampasidis, Kunstantinos Kyr'takopotikts, Emmanuel Bratsolis, Mathieu ilirtzig, Athéna Cousienù & Xenophon Moussas

	Anaina Solomonidau,
	(scale 0-8), was set to be 8 values and the total amount of
	erupted material volume of 2,800 km3. Taking into considér
	ation the amount of the erupted material, the size of the vol
	canic structure and the hazards that could affect the satellite,
	we can assume that super volcanoes could be hosted in Titan's
	geological history. The candidate cryovolcanic figures
	Ganesa Macula, Thi Regio and Hotei Regio could resemble
	the supervolcanic structures seen on Earth.
	In this study, we focus on the volcanological structures
	like the ones seen in TUi Regio and Hotei Regio that resem
	ble terrestrial volcanic terrains and characteristic s. Toi Regio
	is a massive (1,500 km)
	flow fïeld-like figure
	that could possibly hâve
	formed after accumula
	tion of cryolava flows
	erupted	at	different
	times, following	the
	area's topography. On
	Earth, a massive édifice
	resembling T\ri Regio,
	emerges in the Tblarosa
	Basin in south-central
	New	Mexico,	USA.
	Carrizozo	flow	field
	(Fig. 14) is 75 km long
	and covers 328 km2. The
	volume of eruptive ma
	terial	was	4.3	km3
	(Bleacher étal., 2008).
	The field was probably
	formed from periodic	Fig. 14. (up) Carrizozo flow field,New
	déposition of eruptive	Mexico, USA (USGS). (down) One of
	material spewingfrom a	the largest candidate cryovolcanic
	source located 27	m	flows on Titan, TUi Regio

  Is liquid water présent on Enceladus, either in a subsurface océan, in the plume vent régions, or elsewhere? How extensive and long lived is the water, if présent, and what is its chemistry? How does the liquid réservoir communicate with the surface? How thick is the ice crust and how uniform is that thickness? What energy sources are available for life? Is life présent there now?

	The standards for life that Enceladus' possible océan does not. fall into are the
	sunlight, the oxygen compounds, and the organics produced on a surficial-crust
	environment. Environments and ecosystems that do not meet the aforementioned
	prerequisites exist and evolve on Earth as vvell. Such an environment is located
	deep inside South Africa's surface, where sulfur-reducing bacteria consume
	hydrogen and sulfate, produced by radioactive decay (McKay et al., 2008; Muyzer
	and Stams, 2008). In addition, other analogous to Enceladus ecosystems are
	found within the magmatic volcanic rocks, which are produced through the activ-
	ity of metasomatism. The metasomatosis of the volcanic rocks under the reaction
	of water produces methanogens of hydrogen on which the primary productivity
	is based on (McKay et al., 2008). Cassini's data showed that methane is présent
	in the plumes. The terrestrial comparison regarding the ecosystems suggests that
	plume's methane may be biological in origin (McKay et al., 2008).
	Spécifie questions relevant to the goal of understanding habitability on
	Enceladus include:

  6. Future Exploration of the Kronian SatellitesCould Titan be providing us with hints as to the future of our own planet? Indeed, in addition to the past, Titan appears to be an analogue, albeit with different working materials, of the future State of the Earth when surface conditions preclude stable equatorial/midlatitude océans. If we are to focus on the Earth and its climate, as well as on its organic chemistry, we need in the future to concentrate on another object in the solar System that sustains an active hydrologie cycle with surface liquids, meteorology, and climate change. The Cassini-Huygens mission has firmly established that Titan is such a body, in vvhich the active working fluid of the hydrologie cycle is methane. The cycle is active but different from the Earth because Titan lacks a surface methane océan. It possesses, however, methane lakes and seas, fluvial érosion, rounded pebbles, and liquid methane in the soil at the

	Huygens site.

  . Additionally, a scénario that consists of a thick icy crust of almost 15 km suggests a stratigraphie model with a 100 km deep océan, which may be extremely deep, aboutlO times deeper than the deepest point of Earth's océans, the Challenger Deep on the Mariana Trench, which is 11 km below sea level and the lowest élévation of the surface of the Earth's crust. The potential océan of Europa would contain an amount of water twice the entire terrestrial surface hydrological System.

Table 1

 1 Composition measurements in Satum's deep atmosphère and their conséquences Another means to better understand the formation of the Solar System is through a detennination of the planet's core mass and global composition: we do not yet know whether Satum possesses more or less heavy éléments than Jupiter, and we do not know which planet has tire largest core. This impacts directly on the foimation models of these planets and of the Solia System as a whole.Satum is also fascinating for its intriguing meteorology. Contrary to Jupiter, whose rotation axis is almost perpendicular to its orbit, Satum has an inclination of 20°(similar to that of the Earth), plus rings which lead to marked seasonal variations. Perhaps as a resuit of this, Satum undergoes the most intense stomis of the Solar System, with planet-wide events that can last for months. Its zonal wind pattern is similar to Jupiter, but appears to vaiy significantly, as shown by the mysterious slowing of Satum's équatorial jet from 450 m/s in the Voyager era, down to 250 m/s at tire arrivai of the Cassini spacecraft. Two key éléments are missing to model the causes of these variations: Satum's deep water abundance as a powerful source of convective potential energy, and Satum's deep rotation field.

	Species	Conséquence
	He
		Crucial for accurate understanding of the thermal évolutions
		of Satum and Jupiter
	Ne	Test prédiction of Ne capture in He droplets. Refine H He
		phase séparation diagram
	CH4	Fine dét ermination crucial to understand the formation of the planet
	nh3 nr,sh	Key to décidé between models of planetesimal delivery and planet
		formation. Important for Satum's meteorology
	H2S NH4SH	Key for planetesimal delivety, with possibility that the abundance
		is linked to that of rocks deep inside. Important for Satum's meteorology
	h2o	(by radiometry); Key to understand the planet's structure, formation,
		and meteorology
	Ar, Kr, Xe	Key to décidé between models of planetesimal delivery and planet
		formation. Link with the compositions of the Sun and protosolax disk
	CO, PH3, AsH3, GeH4	Disequilibrium species are important to understand convection in Satum's
		deep atmosphère.. Help to further test planetesimal delivery models
	D/H	Test models that predict it should be similar to Jupiter and to the
		protosolar value
		Test models that predict value similar to Eaxth
		Important to understand whether N was delivered as N2 or as NH3.
		Test models of planetesimals delivery
		Origin of gas, Test évaporation processes in the early solar System
		Origin ofgas, Test évaporation processes of these gases in planetesimals
		& Springer

Table 2

 2

	Preliminary launch
	characteristics for DTE option	EWES case for DTE trajectory
	(shortest cruise for DTE)
		Launch vehicle	Atlas V-551
		Departure condition	Fmf=4.1 fcm/s
		Launch C3	C3=17 km2/s2
		Launch mass	4,665 kg
		Gravitv assists	Earth-Venus-Venus-Earth
		Arrivai velocity	Vü,r=5.8 km/s
		Cruise time	11 years
		Satum arrivai mass

Table 3

 3 Mb from the two probes. DTE potential at low frequencies will be significantly (by an order of magnitude) improved with the next génération radio télescopes LOFAR (Low Frequency Array) and SKA (Square Kilométré Array.

	Preliminary launch
	characteristics for relay options	Launch characteristics
		2015 EEJS Relay Trajectory
		Launch vehicle	Atlas V-551
		Departure date	12/7/2015
		Departure condition	Fin{=5.2 km/s
		Launch C3	C3=29.5 km2/s2
		Gravity assists	Earth-Earth-Jupiter
		DSM	685 m/s
		Arrivai date	3/30/2022
		Arrivai velocity	vm£=9.4 km/s
		Omise time	6.3 years
		Dry mass	3,073 kg
		Référencé 2017 EEJS Relay Trajectory
		Launch year	January 2017
		C3	C3=2S km2/s2
		DSM	840 m/s
		Arrivai velocity	v1d{=5.8 km/s
		Cruise time	7 vears
		Dry mass	2,935 kg
			ô Springer
		538

Table 5

 5 Ring science operations , with nodes located around 1(L km from Satum's center. The velocity in the Satum's inertial haine would be about ~27 km/s. With tliis approach geometry, tlie

	Altitude above rings (km)	Résolution (m/pixel)	Science objective
	2,000	10	M Lcrostructure/aggregates
	200	1	Ring tliickness
	20	0.1	Particle size distribution

plane

Table 6

 6 Timeline for tlie Lora probe

	Tune before impact	Altitude above	Resolution	Science Objective
	(or ring-plane Crossing)	rings (km)	(m/pixel)
	-1 h 9 min 21 s	2,000	10	Microstrjcfare/aggregates
	-6 min 56 s	200	1	Ring tliickness
	-41 s	20	0.1	Particle size distribution
				ô Springer

TABLE I

 I

		.	Seismic Instruments Spécifications
				Seismic Sensors
	Power		Mass	Veloclty	Opcratlonal
				Bandwldtli	température
	0.32 W	0.200g each	O.OOl-lOOHz	below 100K.

  Table III lists the main benefits and advances in short ternis of mounting seismic equipment in the surface of icy mooas.

TABLE III

 III

				Benefits of space seismic experiments
			Features
				Icy moons	Engineering
			instruments and	Jupiter/Saturn
			applications
	.	The Benefits Of Future Seismic Instrumentation
		Benefits of space seismic experiments
	Features
		Icy moons	Engineering
	Completed scientific experiment	Régional scale geology	New science-Planetary Seisntology
	Small size	Idéal as a payload for a	Low power
		space mission lander	consomption
	Continuons	Improve modeling for the threc-dimensional	Service for bulky data
	operation	internai structure	needed
	Evolution of	Earthquake prédiction	No spécial software
	seismic	Tidal effects caused by	needs to be developed

January 2010 FP4 (Fig.2.15) spectra there is no such différence. In both FP4 files, the v2.5 data présent ripples below 1050 cm'1 as it is shown in Figs.2.14 and 2.15.

HITRAN website: http://www.cfa.harvard.edu/hitran/ GEISA website: http://ether.ipsl.jussieu.fr/etherTypo/?id=950

HIgh-resolution TRansmission molecular AbsorptioN database(Rothman et al., 2005)

v26 band of propane: first resolved détection, latitude independent, constant-to-height

Température (K)

Submitted to National Research Council of US in response to the Planetary Science Decadal Survey Call in 2009.

Dwarf Planet in 2006.

http://planetary.s3.amazonaws.com/assets/resources/ESA/ESA-SPC_20120417_selection-L 1 -mission.pdf http://sci.esa.int/juice http://media.egu2012.eu/media/filer_public/2012/04/19/eguJuice_press.pdf

http://sci2.esa.int/cosmic-vision/JUICE_Yellow_Book_Issuel .pdf

33http://www. lesia.obspm.fr/cosmicvision/tssm/tssm-public/ http://sci.esa.int/science-e/www/area/index.cfm?fareaid=106 http ://opfm .j pl .nasa.gov/titanri skreduction/

http://eed.gsfc.nasa.gov/562/mems accelerometer. htm

January 2008 (flyby T40), April 2009 (flybys T52-T53). and May 2010 (flyby T68). They contain459,1193, and 160 spectra each in FP3 and a similar or higher number in the FP4. Each sélection at high or medium spectral resolution was restricted to émission angles lower tlian 65°for optimum radiative transfer treatment.

'Ita Astrophysical Journal, 760.144 (8pp), 2012 December 1Bampasidis et al.

'l'UE Astrophysical Journal, 760:144 (8pp), 2012 December 1Bampasidis et al.

* Aaeeikdvtot| ttûv Ta&avcùv evepytbv yefcjXoyiKtîjv îtepiox&v arouç ÊopiHpôpouç ïov Kpôvcru Tuâva koa EyvÉÀaôo: Mc>.éxr| yia tt^v KpuüT]tpaicrtEiôii|xa pu ir| XP^ctîi wiCpuÔptûv tpaoparoccbv 6e6opÉvci>v

Heltenic Journal o/'Geosciences, vol.[START_REF] Cowlmg | The magnetic field of sunspots[END_REF][257][258][259][260][261][262][263][264][265][266][267][268]

Prockter, L. M., and 14 colleagues. (1998). Dark Terrain on Ganymede: Geological Mapping and Interprétation of Galileo Regio at HighResolution. Icarus, 135,

20Ne/z2Ne 3£As/3tAi Kr, Xe isotopes Détermine extent of hélium sédimentation in Satum's interior.

ExpAitron (2009) 23:947 976

& Springer

Acknowledgements

For the completion of this work I am indebted to my supervisors Dr. Athéna Coustenis, Prof.

Xenophon Moussas and As. Prof. Panayota Preka-Papadema not only for their invaluable guidance and insight but also for their constant encouragement over the last years.

We thank Nicolas Gorius, Marcia Segura, and Florence Henry for help with the data acquisition and processing during this project. We also acknowledge support from the CNES Cassini program.

Acknowledgments

We gratefully recognize the contribution and valuable help of Alexandre Laurent and Zhi-Fang Xu in the CIRS data préparation and Processing. We also thank GCM modelers (A. Crespin and S. Lebonnois) for providing listings of their vertical distributions at our averaged latitudes. We are grateful to the CIRS team as a whole for valuable assistance with the science and the data préparation and Processing described here. In particular, we would like to thank Jacqueline Mondellini for the data processing in Meudon and the planning and processing team at Goddard.

Acknowledgments

Valeria Cottini is supported by the NASA Postdoctoral Program. Thanks to S Hôrst, E. Wilson and S. Atreya for providing their pho tochemical water profiles for comparison, and to Paul Romani for water chemistry discussions. The US-based authors were funded by the NASA Cassini Mission during the period in which this work was perfonned. N. Teanby was supported by the Leverhulme Trust and the UK Science and Technology Facilities CounciL Acknowledgments The authors would like to gratefully acknowledge the assistance of Dr. Alexander Hayes in providing us witli raw SAR Cassini data and with a thorough and constructive review of the manuscript and subséquent fruitful discussions. We are also grateful to another anonymous referee for helpful comments and suggestions. A. Solomonidou is supported by the "HRACLEITOS II" project, co-fïnanced by Greece and the European Union. of this article.) Please cite this article as; Solomonidou, A., et al., Morphotectonic features on Titan and their possible origin. Planetary and Space Science (2012), http://dx.doi.Org/10.1016/j.pss.2012.05.003 Acknowledgments The authors are grateful to Julie Castillo, Ellen Stofan and an anonymous referee for extremely constructive and extensive Please cite this article as: Solomonidou, A., et al., Morphotectonic features oii Titan and their possible origin. Planetary and Space Science (2012), http://dx.doi.Org/10.1016/j.pss.2012.05.003 Acknowledgment The authors would like to thank Mr Dimosthenis Kouritis for ffuitful discussions conceming the concept of the proposed experiment and Mr Theodosios Chatzistergos who provided valuable assistance with the figures. Anezina Solomonidou is financially supported by the European Union (European Social Fund -ESF) and Greek national funds through the Operational Program "Education and Lifelong Leaming" of the National Strategie Reference Framework (NSRF) -Research Funding Program: Heracleitus IL Investing in knowledge society through the European Social Fund.

Scientist for a day 8.

Appendices

Appendix A

The inversion approach of the Radiative Transfer Equation solutionThe inversion approach

The inversion approach is a method which allows the calculation of the values of atmospheric parameters from the recorded infrared intensities. In brief, we describe the whole System following a reversai sequence, beginning from the results of a physical process and seeking for the embedded parameters which cause the recorded results [START_REF] Tarantola | Inverse Problem Theory and Methods for Model Parameter Estimation: Society for Industrial Mathematics[END_REF]. At first, we should seek for the température structure of the atmosphère. Titan's atmospheric température is considered as a fonction of the vertical distance from the satellite's center and the planetgraphical coordinates.

The équation (4) follows the generic form of the homogeneous Fredholm intégral équation of the first kind in one -linear -dimension which can be formally illustrated as: g(y) = J K(y,x)f(x)dx (5) a The fonction g(y) corresponds to the recorded radiance in our case and can be considered as a known terni as well as the continuous Kernel fonction K(y,x) which is the partial dérivative of the transmittance in respect to the altitude. On the other hand, the /(x) fonction illustrâtes the Planck fonction, the source of the émission.

The term inverse problem is common in Astronomy and in other experimental branches of Science and Engineering such as image and signal processing, geophysics, remote sensing etc, when researchers try to détermine the parameters of a physical System by using the interaction data of its sample and the radiation emitted by a known source (e.g. [START_REF] Rawlinson | Seismic tomography: A window into deep Earth[END_REF][START_REF] Bertero | Introduction to Inverse Problems in Imaging[END_REF]. The Kernel fonction K(y,x) in ( 5) is at least known in principle while g(y) is only known in a determined set of points y,, where i=l,...,m is a finite set of quantities with a certain accuracy [START_REF] Hansen | Numerical tools for analysis and solution of Fredholm intégral équations of the first kind[END_REF]. [START_REF] Scott | A direct method of computation of the transmission function of an inhomogeneous gaseous medium -I: Description of the method[END_REF]. The program has been firstly applied in 1977 in the Laboratoire de Météorologie Dynamique (LMD) studies [START_REF] Scott | A fast line-by-line method for atmospheric absorption computations -The Automatized Atmospheric Absorption Atlas[END_REF] and several upgrades hâve been done up to date.

Atmospheric Radiative Transfer code for

The first large upgrade was taken place in 1986 by Bruno Bezard, who modified the code initially for applying it in the Satumian atmosphère and then for the Titan's case. This release of the code calculated the transmittance and the atmospheric radiance of Titan's atmospheric constituées based on spectroscopic parameters adapted from the GEISA 1977 database (Gestion et Etude des Informations Spectroscopiques Atmosphériques: Management and Study of Spectroscopic Information, [START_REF] Chedin | Une banque de données pour l'etude des phenomenes de transfert radiatif dans les atmosphères planétaires: la banque GEISA[END_REF].

From 1986 to 2005, several modifications hâve been imported into the code, mainly adapting the new GEISA release updates [START_REF] Husson | Management and Study of Spectroscopic Information -The GEISA program[END_REF]1994). In 2006, two separate upgrades hâve been performed. GEISA 97 spectroscopic parameters (Jacquinet- [START_REF] Jacquinet-Husson | The GEISA System in 1996: towards an operational tool for the second génération vertical sounders radiance simulation[END_REF]1999) were adapted in mid-2006, while at the end of the year the new spectroscopic files from the GEISA 2003 version (Jacquinet- [START_REF] Jacquinet-Husson | The 2003 édition of the GEISA/IASI spectroscopic database[END_REF] replaced the old ones, except for the HCN, HC3N, C2H4, C3H4 and C4H2. For the HCN the GEISA 97 lists were kept (Jacquinet- [START_REF] Jacquinet-Husson | The 1997 spectroscopic GEISA databank[END_REF]. On the other hand, new spectroscopic files from Jolly et al.

for the V5 stretching band at 663.2 cm'1 of HC3N replaced the previous ones (Jolly et al., 2007). C2H4 lines were initially taken from Blass et al. [START_REF] Blass | Absolute intensities in the v7 band of ethylene: tunable laser measurements used to calibrate FTS broadband spectra[END_REF] and M. Rotger 2006 (pers. comm.) [START_REF] Rotger | Line positions and intensities in the vl2 band of ethylene near : An experimental and theoretical study[END_REF]. The latter data were added later in the upgraded version of GEISA 2003atlas (Jacquinet-Husson et al., 2008). The vg band data of C3H4 was replaced from the files provided by G. Graner (pers. comm.). Only C4H2 vg and vg+ V9-V9 hot

The folder "ARTT-src-*.*.*-yyyy-mm-dd" contains some files too with data needed for the installation. The most important files from these is a file called "config.txt" which contains the installation path and the compiler options and a makefile with the installation procedure. If we execute the command "make" in the folder containing the source files then we can see the options for the installation of ARTT.

After installing ARTT a folder called "bin" is created, in the installation folder, which contains the exécutable files. Another folder called data is created in the path specified within the file "config.txt". This folder contains the following sub-folders:

• ARTT-work-*.*.*: If we execute the "make" command in this folder then a list with ail the available actions is seen. In order to generate a model of Titan's atmosphère a template file is needed with the name "Name_of_the_model.mod"

• bandes-*.*.*: This folder contains the files "c3h8.dat" and "continuum.dat"

• cirs: In this folder we can find the data ffom CIRS used for the création of the plot based on space data. This is a copy of the cirs folder of the source distribution.

B.3 ARTT Upgrade procedure

To add a new molécule in our model, we should proceed through the following procedure.

The sources databases for the spectroscopic files we search are the following:

Gestion et Etude des Informations Spectroscopiques Atmosphériques (GEISA) (Jacquinet- Lopes et al., 2010). Furthermore, these features can be modified under the influence of exogenetic processes: the résultant niorphotectonic structures are mainly mountains, ridges, faults, and structures of probable cryovolcanic origin as we will argue further down. Titan, despite its small size, displays surface features that resetnble the structure of terrestrial volcanic fields albeit they are much more extensive. For example, Hotei Regio (see Section 4.2) covers an area of 140,000 km2, which is an order of magnitude larger than Harrat Khaybar (14,000 km2) in Saudi Arabia, which represents one of the most extensive volcanic fields on Earth [START_REF] Camp | The Arabian continental atkali basait province: Part D. Evolution of Harrats Khaybar, Ithnayn, and Kura[END_REF].

Silicate and icy tectonism

A surface feature largely owes its shape to the material composing the planetary body's crust and to the forces tirât formed it. The response of the trustai material to the applied stresses, defines to a large extent the main topographie terrain of an area, along with the atmospheric conditions. Even thougli geological features such as mountains, faults and rifts on Titan présent similar Visual characteristics, the type of material that builds the features plays an important rôle. Indeed, the properties of the source material such as viscosity, elastirity and density in addition to the geological forces control the structural characteristics of the feature such as height, expansion and gradient slope.

On Titan, the surface should probabty be composed mainly of mixtures of water and other ices, organics-tholins, nitriles (e.g Soclerblom et al" 2007a), while, most likely, its interior is composed of rock and high-pressure ice (Tobie et al, 2005). Since the wellknown tectonic features of Earth are dosely linked to silicate geology, we must first assess the similarities and différences between water ices and silicates, so that our comparative study can be based on reasonable assumpdons. The icy crust of the outer System satellites possibly reacts in a brittle fashion to the application of stresses, similaiiy to the Earth's rocky upper crust (Collins et al., 2009). Both on Titan and Earth this reaction changes in proportion with dejxh.

However, while water ice and silicate rock exhibit similar frictional strength (Beeman et al., 1988), wlien ductile yielding becomes important, ice is about ten times weaker than silicate rock (Melosh and Ninuno, 2011). The major différences and similarities between water ice and silicates are noted in Collins et al. (2009) and summarized in Table 1.

Table 1 shows that silicate materials, when compared to water ice, exhibit higher viscosity, Young modulus i.e. the ratio of linear stress to linear strain and melting température, but display lower density. As a resuit, the homologous température, on which rheology dépends, is reached at greater depths in silicate environments while silicate magma éruptions are statistically more possible to occur than emptions of ice (Collins et al., 2009) Our présent knowledge of the Chemical and isotopic composition of the Jupiter and Satum atmospheric Chemical and isotopic composition is given in Table 4 together with the required précision of a Satum Probe Gas Analysis System (GAS).

The précisions listed represent a substantial improvement ffom the Galileo GPMS measurements. This performance will not only establish a full comparison with the Jupiter values, but will provide a robust data set for comparison with other planetary values. For example, différences in Xe isotopic fractionation between solar and meteoritic and between solar and terrestrial of <1% and <3% respectively per Dalton are resolvable with this précision. The static mass spectrometer design described