Président Claudio Codella

Rapporteur Serena

Viti Rapporteur

Valentine Wa K E L A M Examinateur

Examinateur Claudine Kahane

Cecilia Ceccarelli

Directeur De Thèse

Bertrand Lefloch

T H Èse

Prepar Ée

Introduction xxi

1.1 Before the 80's, the collapse was treated in the spherical symmetry approximation (left panel). With the discovery of the outflow phenomena and circumstellar disks, it was clear that the spherical approximation was not valid 1.3 Fragmentation is a multi-scale process from spiral arms of a galaxy down to protostars whose characteristic lengths spanning from 10 kpc to 100 AU. Taken from [START_REF] Schulz | From Dust To Stars Studies of the Formation and Early Evolution of Stars,P r a w i svii[END_REF]. . . . . . . . . . . . . 1.4 Column density map of Aquila derived from SPIRE/PACS data from the Gould Belt Survey carried out by Herschel Space Observatory. The candidate Class 0 protostars and bound prestellar cores identified in Aquila by [START_REF] Bontemps | [END_REF]a n d [ Könyves et al., 2010]a r es h o w na sg r e e ns t a r sa n db l u et r i a n g l e s ,r e s p e ctively. Note the good correspondence between the spatial distribution of the cores and protostars and the regions where the filaments are unstable to gravitational collapse. For more details see [André et al., 2010]. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 1.5 Evolutionary sequence and spectral energy distributions of lowmass YSOs as proposed by André (1994). The four classes 0, I, II, and III correspond to successive stages of evolution. Adapted from [Feigelson and Montmerle, 1999] 1.6 (top) Interferometer map of the molecular jet in the HH211 outflow. The CO J=2-1 line (white lines) [Gueth and Guilloteau, 1999]i so v e r l a i do naf a l s ec o l o ri m a g eo fH 2 [McCaughrean et al., 1994]. The red lines delineate the 230 GHz continuum emission from dust concentration around the protostar. (bottom) In color scale we see the near infrared H 2 emission, in contours blue and red there are the blue-and red-shifted emission of SiO , the green contours mark the dust emission at 870 micron surrounding the protostar. [Lee et al., 2007]. . . . . . . . . . . . . . . . . . . 1.7 Gray scale of the 2.12 µmS ( 1 )l i n ee m i s s i o nf r o mm o l e c u l a rh ydrogen. The cross marks the position of the IRAS 23011+6126 central source. Taken from [Ladd and Hodapp, 1997]. . . . . . . . 1.8 Overabundance factors with respect to dense core values for a number of molecules towards the outflows from L1448-mm and IRAS 04166+2706. Data from Tafalla et al. (2010). Taken from [START_REF] Tafalla | IAU Symposium,V o l .2 8 0o fIAU Symposium[END_REF]. . . . . . . . . . . . . . . . . . . . . . . . . .

1.9 left:C O2 →1 emission from the L1157 outflow. right:M a p so f integrated intensity for different molecular species towards the blue (southern) lobe of the L1157 outflow. Note the different extension and peak position of each species. Taken from [Bachiller and Perez Gutierrez, 1997]. . . . . . . . . . . . . . . . . . . . . . . . . . . . 1.10 HH 1/2. These are among the brightest HH objects in the sky, they form bow shocks in a highly collimated bipolar outflow with dimension of 2.5 arsec (0.34 pc), symmetrically opposite from a young star which is ejecting material along its polar axis. . . . . .

LIST OF FIGURES ix

1.11 Several types of outflows are observed in low-mass YSOs. High velocity and very collimated jets are mainly observed in Class 0 sources. The figure indicates the trend that outflow collimation decreases with the age. Similarly the flow velocity (several hundreds km s -1 )d e c r e a s e sf r o mt h ej e tt ot h eo u t e re d g e so fad i s k wind (≈ 5k ms -1 ). The process that originates the outflows is still not fully understood, but it seems that magnetic fields and rotation trigger the jet launching mechanism. Taken from [START_REF] Schulz | From Dust To Stars Studies of the Formation and Early Evolution of Stars,P r a w i svii[END_REF].................................. 1.12 Observable molecular outflow properties predicted by the four leading broad classes of models: 1) turbulent jet [Canto and Raga, 1991]; 2) jet bow shock [Raga and Cabrit, 1993]; 3) wide-angle wind [Shu et al., 1991]; and 4) circulation models [Fiege and Henriksen, 1996b]. Taken from [START_REF] Arce | Protostars and Planets V[END_REF]. . . . . . . . . . . . 1.13 The postshock temperature structure of a fast shock. Taken from [Hollenbach and McKee, 1989]

3.1

The Cepheus OB3 active star formation region overlaid on the map of visual extinction, obtained from 2MASS [Lombardi and Alves, 2001]. Large circles denote the clouds associated with young stars. The meaning of the different symbols are as follows: Filled triangles are T Tauri stars, filled squares are Herbig Ae/Be stars, filled circles are weak-line T Tauri stars, open squares are photometric candidate and possible PMS members, X are H α emission stars, crosses are T Tauri candidates. For details see [START_REF] Kun | Star Forming Regions in Cepheus[END_REF]. .

Cep OB3b region. In the image the molecular clouds Cep A to

Cep F, the dark cloud Lynds 1211, as well as the most prominent associated young stars are labeled. In contours the distribution of the visual extinction [Dobashi et al., 2005]. The lowest contour of the extinction is at A V =1mag,andtheincremen tis0.7mag. F or details see [START_REF] Kun | Star Forming Regions in Cepheus[END_REF]. . . . . . . . . . . . . . . . . . . . . .

Spectral energy distribution of Cep E-mm (IRAS 23011+6126

) source observed and modeled. Taken from [Crimier et al., 2010]. . et al., 2001]. Taken from [Velusamy et al., 2011]. (right) Superposition of the H 2 1-0 S(1) 2.12 µm emission (gray scale) on the 13 CO and 222 GHz continuum in contours. The dashed contours toward the north correspond to the redshifted lobe (v LSR =-9 . 5t o-2 . 5k ms -1 ) ,w h i l et h es o l i dc o ntours to the south correspond to the blueshifted lobe (v LSR =-20.5 to -13.5 km s -1 ). Note that the 222 GHz continuum trace a double component. Adapted from [Moro-Martín et al., 2001] )anditextendsto red velocities up to reach a second component that peak at about 50 km s -1 .T h emiddle panel shows emission comes from the central source, the spectrum display a main component centered at ambient velocity and weak emission comes from the blue-shifted velocities. The bottom panel exhibit the spectrum comes from the south lobe. It shows again a bright component at ambient velocity and a second component that peak around -125 km s -1 .T h e s e second components at high velocities are interpreted as evidence of the presence of jets. [Lefloch et al., 1996]. . . . . . . . . . . . .

(left)

H 2 1-0S ( 1 )c o n t o u rm a ps h o w st h eb i p o l a rs t r u c t u r eo f Cep E. There are also two H 2 knots at the west of the IRAS source region labeled as W1 and W2. [Eisloffel et al., 1996]i n t e r p r e t these knots as a second independent outflow and, therefore, the presence of a second outflow source can be inferred. (right) Cep E could be a double or multiple system. In gray scale, the 2.12 µm S(1) emission from H 2 overlaid to contours of redshifted [-6 to 0 km s -1 ] (dotted) and blueshifted [-14 to -20 km s -1 ]( s o l i d ) 12 CO J= 2-1 emission. Both maps are centered on the IRAS source. For details see [Ladd and Hodapp, 1997]. . . . . . . . . . . . . . . . .

3.11

The figure shows the 450, 850 and 1300 µm images of the IRAS 23011+6126 region. The dust emission trace the quadrupolar outflow observed by [Ladd and Hodapp, 1997](Seetext). T ak enfrom [Chini et al., 2001] Clearly the emission at positive velocities comes from two components, the first one around 25 and 45 km s -1 and the second one between 50 and 70 km s -1 .A tn e g a t i v ev e l o c i t i e s ,t h ee m i ssion similarly is mainly from one component between -25 and -50 km s -1 and a second component from -50 to -90 km s -1 .I nt h i s transition the EHV emission is not observed. . . . . . . . . . . . .

4.3

The high-velocity jet emission from CepE CO (2-1) emission map at 1" resolution obtained with the PdBI (Lefloch et al. 2011). Central position is placed at 23h03m13.0s +61 • 42 '21" (J2000). The red circle represents the IRAM beam at 3mm, the green one at 1.3mm and the blue one at 0.8mm. . . . . . . . . . . . . . . . . .

(left)

A spectra with fluctuations in the continuum baseline, it is different to zero along the spectra. (right) The same spectra of the left but corrected, it has been removed a baseline in order to obtain a continuum close to zero. . . . . . . . . . . . . . . . . . .

4.5

Examples of spikes present in the spectra. Note that the intensity of these spikes are 4 times brighter than the emission from the lines.

4.6 Example of a ghost line: CS (5-4) spectra (black line) with a contamination in the wing by CO . For comparison we overplot the CS 4.9 Three kinematical components were found: the ambient gas at ∼ -10.9 km s -1 ,t h ee n t r a i n e dg a sb e t w e e n-9 0→ -12 and -9 → +30 and the jet, detected at -125 km s -1 (south jet) and at around +50 At the left (0,0) position, and at the right the (-12,-18) position. In both pannels the spectra of lowest transition were observed with IRAM-30m telescope and transition CO (4 -3) were observed with CSO for the (0,0) offset and with JCMT for the offset (-12,-18). Intensities are in units of antenna temperature (K). A vertical offset is added to each emission line to help the comparison. The vertical dotted lines show the velocity range where appear the EHV component associated with the jet. (See the text). . . . . . . . . . 179 5.2 SiO jet emission from both; the source and bowshock positions. A vertical offset is added to each emission line to help the comparison.

The vertical dotted lines show the velocity range components. . . 180 5.3 HCO + jet emission from both; the source and bowshock positions.

A vertical offset is added to each emission line to help the comparison. The vertical dotted lines show the velocity range components. 181 5.4 HCN jet emission from both; the source and bowshock positions. A vertical offset is added to each emission line to help the comparison.

The vertical dotted lines show the velocity range components. . . 181 5.5 SO jet emission from both; the source and bowshock positions. A vertical offset is added to each emission line to help the comparison. The vertical dotted lines show the velocity range components. . . 183 5.6 (left) CO, SiO, HCN, HCO + and SO spectra at (0,0) offset with detected jet emission at high velocities. Vertical dash line marks the ambient cloud velocity (∼ -10.9 km s -1 ), and vertical dotted lines mark the velocities at 30, 50 and 80 km s -1 .I n t h e r a n g e 30 to 45 km s -1 we have the moderate velocity component of the north jet (North-MV), and between 50 to 80 km s -1 we have the high velocity component (North-HV). (right)Thesamemolecular transitions at (-12,-18) offset. Unlike the spectra on the left, only CO and HCO + spectra show clearly detected jet emission at South-EHV. Vertical dotted lines mark the velocity limits between -100 and -150 km s List of Tables [Fuente et al., 2012]. When they are on the pre-main-sequence they are known as Herbig Ae/Be stars.

Theses young stellar objects (YSOs) are located in less complex regions than high mass star forming regions and they are in larger number closer to the Sun (< 1k p c ) . A tl e a s ti ns o m ec a s e s ,i ti sp o s s i b l et os t u d yt h ep h y s i c a l and chemical structures at similar spatial scales as those achievable in low mass protostars. Once on the main sequence, the IM stars are among the sources that dominate the FUV interstellar radiation field that regulate the different phases of interstellar medium and therefore the overall Galaxy star formation process [Habing, 1968]; [Gondhalekar and Wilson, 1975]. In addition, the study of IM-YSOs is fundamental to establish the initial conditions of the planetformation process around Herbig Ae/Be stars, the precursor of Vega-type debris disks [START_REF] Schütz | [END_REF].

Even though the IM-YSOs are important in the study of star formation, very little is known about the formation and first evolutionary stages of IM protostars.

Unbiased spectral surveys are a powerful tool to characterize the chemical composition of an astrophysical object, and the only way to obtain a complete census of the chemical species. A spectral survey provides also multiple lines from the same molecule, giving the possibility of a multi-frequency analysis and modeling. Besides, different lines from transitions with different excitation conditions (upper level energies and Einstein coefficients) are excited at different temperatures and densities. Therefore, they probe different physical regions in the line of sight.

Finally, through line profiles, we can obtain kinematic information,a n d identify structures along the line of sight, as multiple sources, outflows, jets or cavities,e . g . [ Caux et al., 2011].

Several unbiased spectral surveys,c a r r i e do u ti nt h em i l l i m e t e ra n ds u bmillimeter bands accessible from the ground, were obtained mainly toward highmass protostars e.g. Orion KL [Margulès et al., 2009], [Tercero et al., 2011] and [Tercero et al., 2012], AFGL 2591, IRAS 20126+4104 [Plume et al., 2007], IRAS 23385+6053 [Thompson and Macdonald, 2003], NGC 6334 I, I(N) [Thorwirth et al., 2007] and compact HII regions as W3 IRS5, IRS4 [Kim et al., 2006], G34.3+0.15 [Kim et al., 2000], G5.89-0.39 [Thompson and MacDonald, 1999]. These targets are mainly chosen because they have bright molecular lines and a rich chemistry. On the contrary, only two solar-type protostars have been studied in the mm/submillimeter window: IRAS 16293-242216293- ([Blake et al., 1994]]; [van Dishoeck et al., 1995]; [Caux et al., 2011]a n dL 1 5 2 7 [ Sakai et al., 2008]. A complete list of these surveys can be found in the review by [Herbst and van Dishoeck, 2009].

The outflow phenomena (jets, winds and bipolar molecular outflows), are an inherent phase in the process of star formation observed in YSOs of all range of masses at millimeter wavelengths. In the last years, biased observational studies were carried out towards young protostellar outflows, e.g. [Bachiller, 1996], [Bachiller et al., 2001], [Jiménez-Serra et al., 2004]. The latest works present unbiased surveys towards L1157-B1 with Herschel satellite as part of the key program CHESS [Lefloch et al., 2012], [Codella et al., 2012], the most studied outflow to date.

Recently, the 500-2000 GHz range was covered by the Herschel satellite as part of the key programs CHESS [Ceccarelli and CHESS Consortium, 2010]a n d HEXOS [Bergin et al., 2010]t o w a r d sl o w -t oh i g h -m a s ss t a rf o r m i n gr e g i o n s .

Up to date only few individual IM protostars have been studied in detail ( [Hogerheijde et al., 1999]; [Neri et al., 2007]; [Schreyer et al., 2002]; [Fuente et al., 2007]; [Crimier et al., 2010]; [Lefloch et al., 2011]). However, there are not systematic studies in IM range as in the case of of low-and high-mass protostars.

Given the lack of a systematic study of the mm/submm spectrum of IM protostars, during my thesis I carried out an unbiased spectral survey towards Cep E and its molecular outflow.

This source was chosen because it is a nearby (distance 730 pc) young, Class 0c a n d i d a t e ,w h o s ea g ei se s t i m a t e dt ob ei nt h er a n g eo f2 0-4 0 0×10 3 years [Fro ebrich, 2005]. It has a remarkably bright outflow (10 L ⊙ )a n dab o l o m e t r i c luminosity of ∼ 100 L ⊙ [Lefloch et al., 2011]. This thesis manuscript is composed of six chapters. In chapter 1,Ig i v ea n introduction of the process of star formation of low-, intermediate-and high-mass regimes and a general review on outflows, observationally and theoretically. In chapter 2,Ipresen tsomefundamen talconceptsofradioastronom y . Inchapter 3, I present a detailed description of the source Cep E and the studies carried out to date. In chapter 4,Ip r e s e n tt h es pe c t r a ls u r v e yo v e r v i e w ,t o g e t h e rw i t ht h e line catalogue. In chapter 5,Ipresen ttheresultsofthejetcomponen tanalysis. In chapter 6 Iw i l lc o n c l u d et h et h e s i s .

Chapter 1

Star formation and ejection phenomena

The star formation is not a process that only occurred in the past, but it is an active process that takes place nowadays in our Galaxy even, in regions close enough, that the birth of stars can be examined in detail.

The molecular clouds are the sites where the star formation occurs. Gas in these regions is so cold (∼ 10 K) that the bulk of radiation is emitted at radio and mm wavelengths. Protostars are also relatively cold (T≤1000K) and emit in the mm to IR wavelengths.

The radio/mm astronomers use single dishes telescopes to map extended areas of the sky covered by these clouds. For a detailed study of individual regions they use interferometers, especially for observe the distribution of matter around newly formed stars.

After the 80's instrumental advances in IR and FIR detections provided the first high-sensitivity maps at infrared wavelengths. The launch of IRAS (Infrared Astronomical Satellite)i n1 9 8 3a l l o w e dt os c a nm o r et h a n9 6p e r c e n to ft h es k y at mid-and far-infrared wavelengths and provided the first census ever of protostars. Years later in 1995, ISO (Infrared Space Observatory), allowed to carry out higher sensitivity, and higher spectral resolution observations on a larger wavelength range than IRAS. More detailed studies have been possible thanks to the launching of Spitzer Space telescope in 2003 and recently in 2009 with Herschel

Molecular clouds

Space Observatory.

Molecular clouds

Molecular clouds get their name from the fact that they are made mainly of molecular hydrogen (H 2 ). Although molecular hydrogen is their main constituent, molecular clouds are detected and mapped using the second most abundant molecule, carbon monoxide (CO) because the upper energy level of the first rotational level (J = 1) of the molecule is only 5.5 K and then readily exited at cloud temperatures. On the contrary, the H 2 first rotational transition lies at about 500 K, making impossible to detect it at the ∼ 10 K molecular cloud temperature. Thus, large-scale molecular cloud maps are usually carried out with the J=1 -0r o t a t i o n a lt r a n s i t i o nl i n eo ft h eC Om o l e c u l e ,a taw a v e l e n g t ho f2 . 6m i l l i m eters. The 2.6 mm transition is most readily excited by gas with number density around 10 3 cm -3 . However, other tracers are available to explore denser regions, for example the 3.1 mm line of CS or the 1.3 cm of NH 3 are transitions excited near to 10 4 cm -3 .T h i si sb e c a u s eo ft h e i rl a r g e rc r i t i c a ld e n s i t y .

Molecular clouds come in all sizes. In the following we briefly review the main classes of molecular clouds.

Giant molecular clouds

Large clouds, called, Giant Molecular Clouds (GMC), have masses in the range of 10 5 -1 0 6 M ⊙ ,s i z e so ft e n so fp c ,t e m p e r a t u r e sa b o u t1 0-3 0K ,a n dn u m b e r densities in the order of 10 3 -10 5 cm -3 . The 80 percent of the molecular hydrogen in our Galaxy resides in GMCs. A well-known example is the Orion molecular cloud at 450 pc of distance. Massive stars (M * > 9M ⊙ )a r ef o r m e di nt h i st y p e of clouds. Recents studies suggest that Infrared Dark Clouds (IRDCs) are the cold condensations inside which massive stars form e.g., [Beuther and Steinacker, 2007], [Kauffmann and Pillai, 2010].

Dark clouds

Dark clouds1 have masses and radii in the range of 10 3 -1 0 4 M ⊙ and 2 -5 pc, respectively and they do not show high mass star formation. This kind of clouds form exclusively low and intermediate mass stars (M * ≤ 8M ⊙ ). The Taurus cloud at a distance of 140 pc, is a good example of this type of cloud.

Molecular transitions sensitive to higher densities like NH 3 (n crit 10 4 cm -3 ) [Myers, 1985]h a v eb e e nu s e dt oe x a m i n et h ed e n s e rr e g i o n sn o wc a l l e ddense cores.T h e s ed e n s ec o r e sm a yo rm a yn o tc o n t a i ni n f r a r e ds o u r c e si nt h e i rc e n t e r .

These cold (≤ 10 K) and dense (≥ 10 5 cm -3 )c o n d e n s a t i o n sa r eb e l i e v e dt o be the first step towards the formation of low-mass stars and are called prestellar cores (e.g. [START_REF] Di Francesco | Protostars and Planets V[END_REF], [START_REF] Caselli | IAU Symposium,V o l .2 8 0o fIAU Symposium[END_REF]).

Turbulence and the evidence for magnetic field in molecular clouds

The observation of supersonic motions in molecular clouds (e.g., [START_REF] Zuckerman | [END_REF]) raised the question of how these motions could be supported. Supersonic motions are expected to dissipate their energy quickly in highly radiative shocks because of the very short cooling time of molecular gas or metal-rich atomic gas ( [Goldreich and Kwan, 1974]).

Closely related was the issue of the support of molecular clouds (MCs) against gravitational collapse, since it was soon realized that the observed motions could not be understood as a gravitational collapse ( [START_REF] Zuckerman | [END_REF]), although MCs contain many Jeans masses.

Theoreticians therefore formulated the hypothesis that MCs were primarily magnetically supported ( [Mouschovias, 1976], [McKee and Zweibel, 1995] [START_REF] Shu | [END_REF]). They interpreted the observed motions as long-wavelength hydromagnetic waves ( [Falgarone and Puget, 1986]). In this scenario star formation is controlled by ambipolar diffusion:

) i.e., E Gravitational ∼ E Kinetic ∼ E Magnetic >> E THermal ([
The overall picture of low-mass star formation while the electrons and ions remain linked to the field lines, the neutral particles (mostly molecular hydrogen) do not feel the magnetic field and pass through, allowing collapse.

It was also shown that the properties of the observed "turbulence" (e.g. [Larson, 1981]) could be understood if the motions were sub-Alfvenic ( [START_REF] Mouschovias | [END_REF], [Xie, 1997]).

Studies based on the Zeeman effect in lines of molecules such as OH and CN, which probe regions of dense gas, resulted in measures of the magnetic field of a few to a few tens of µGf o rr e g i o n sw i t hd e n s i t i e so f∼ 10 4 cm -3 ([Troland and Crutcher, 2008]) and lower, ∼0.5 µG, in dense regions (e.g. [Falgarone et al., 2008]; [Crutcher et al., 2010]).

The next section will describe, in a general way, the process that leads to low-mass and high-mass star formation. The multiple observational studies towards region of star formation demonstrate that a variety of stars of different masses from 0.1 to <100 M ⊙ continuously form. But what is the sequence of events that brings a molecular cloud condensing into a star? and what are the right conditions to trigger the star formation process?

Ab r e a k t h r o u g hi nt e r m so ft h e o r e t i c a lm od e l i n go ft h es t e l l a rf o r m a t i o np r ocess was the work by Larson in 1969. Using numerical models, Larson was able to simulate, for the first time, all the evolution process, from a fragment of interstellar cloud of gas to a pre-main-sequence object. Larson simplified the problem and modeled the evolution of a molecular cloud fragment with spherical symmetry without rotation or magnetic field and initially homogeneous and isothermal, for different values of its mass. The initial conditions for the models were determined in a way that the initial molecular spheres were just at the point to become unstable, according to the criterion of Jeans.

Asecondmajormilestonew asac hiev edb ySh u(1977),whodev elopeda"selfsimilar collapse" theory. Shu found a semi-analytical solution for a sphere in isothermal collapse. He described the collapse like non-homologous, i.e., in his model the central parts of the sphere collapse earlier than the outer parts introducing pressure gradients. In practice, the inner parts collapse first and the outer parts remain in hydrostatic equilibrium until they are reached by the outward propagating expansion wave that moves at the speed of sound. This type of collapse is known like inside-out collapse.

However, in the 1980s, it was soon realized that this simplistic view had fundamental problems and was not in agreement with the new observed phenomena like circumstellar disks, outflows, and multiple star formation. The spherical symmetry was then abandoned (Fig. 1.1). The mass ejection phenomenon was first discovered through detection of the Doppler shift in the CO line towards the Orion Nebula in the BN-KL region [START_REF] Kwan | [END_REF], where, highvelocity gas is streaming away from the infrared star. In the vicinity there is also copious near-infrared emission from H 2 ,r e s u l t i n gf r o mc o l l i s i o n a le x c i t a t i o n of the molecule in shock fronts, and interstellar masers, small regions of strongly beamed radiation from molecules such as H 2 O.

Several stages have been proposed (e.g., [Andre et al., 1993], [Lada and Wilk-The overall picture of low-mass star formation ing, 1984]) to describe the process of the birth of a star, but any classification scheme necessarily has a degree of arbitrariness. The following sections present the main stages involved in the low-mass star formation process (Fig. 1.2). Stars are formed by the gravitational contraction of interstellar molecular gas, but as the mass of a molecular cloud is much larger than a single star, it is obvious that a star does not originate from the collapse of a cloud as a whole. There is a hierarchical and multiscale process [Efremov and Elmegreen, 1998], called fragmentation,i nw h i c ht h ep a r e n tc l o u db r e a k su pi n t os u b s t r u c t u r e sl i k e filaments or subclouds, which might themselves break into smaller and denser structures, like cores (Fig. 1.3).

Fragmentation and pre-collapse phase

Starless cores represent the initial condition of (low-mass) star formation, prior to protostars, and if they are also gravitationally-bound (cf. [START_REF] Andre | Protostars and Planets IV[END_REF]; Fragmentation is a multi-scale pro cess from spiral arms of a galaxy down to protostars whose characteristic lengths spanning from 10 kpc to 100 AU. Taken from [START_REF] Schulz | From Dust To Stars Studies of the Formation and Early Evolution of Stars,P r a w i svii[END_REF] [di Francesco et al., 2007]) they are known as prestellar cores [START_REF] Caselli | IAU Symposium,V o l .2 8 0o fIAU Symposium[END_REF]. [START_REF] Caselli | [END_REF], [Benson and Myers, 1989], [START_REF] Caselli | IAU Symposium,V o l .2 8 0o fIAU Symposium[END_REF]) and in (sub)millimeter continuum (e.g. [Motte et al., 1998], [Commercon et al., 2012]). The density peaks at about ∼ 10 5 -10 6 cm -3 following a flat inner profile e.g., [Ward-Thompson et al., 1994], [Motte et al., 1998]. It is believed that prestellar cores result from the gravitational fragmentation of filaments in the cold ISM (Fig. 1.4)( e . g . , [ Könyves et al., 2010]; [START_REF] Bontemps | [END_REF]; [Men'shchikov et al., 2010]; [Ward-Thompson et al., 2010]).

The dense cores are made of a few M ⊙ ,h a v ed i a m e t e r so fa b o u t1 0 4 AU, T ≤10 K and n(H 2 ) ≥ 10 4 cm -3 . Abundances of C-bearing molecules drops towards the center while N-bearing species seem to survive (NH 3 enhanced) [Bacmann et al., 2002], [Tafalla et al., 2002], [Jørgensen et al., 2004]. This is consistent with the selective freeze out of molecules onto cold dust grains as the gas becomes centrally concentrated during core formation [START_REF] Tafalla | [END_REF]. CO freezeout is particularly important as the depletion of CO from the gas and the low temperatures favor the extreme molecular deuteration [START_REF] Ceccarelli | Protostars and Planets V[END_REF], [Bacmann et al., 2002].

Despite the last decade advances in the comprehension of the pre-stellar phase, there are still questions to answer. Why are clouds so fragmented but still so close to Virial Equilibrium? ( [START_REF] Sadavoy | [END_REF], [START_REF] Dib | [END_REF]) Why do they Recent results obtained with Herschel Space Observatoty (HSO) suggest that core formation occurs in two main steps: (1) Filaments form first into cold ISM, probably as a result of the dissipation of MHD turbulence . (2) The densest filaments fragment into prestellar cores via gravitational instability. This is consistent with the view that filaments form as a result of turbulent compression in low-velocity shocks. [START_REF] Padoan | [END_REF]. However this is still widely debated.

First phase of the collapse

If the clumps in molecular clouds are the sites of star formation, what conditions must exist for collapse to occur? and what are the physics that induces and maintains collapse?

Classically, the main force counteracting gravitational collapse is thermal pressure. This sets a minimum mass that a cloud with density n, and temperature T must has to overcome pressure and to collapse. This minimum mass is called Jean Mass, M J :

M J =1.6
T 10K

1.5 10 5 cm -3 n 0.5 (1.1) The Jean mass M J associated with the average conditions of a clump is a few M ⊙ ,m uc hsmalle rthanthemassofac lumpM cl ∼ 10 4 M ⊙ [Kramer et al., 1998]. Therefore, the rate of star formation should be higher than observed (1-4 M ⊙ yr -1 ). If star formation takes place in molecular clouds, then the overall rate of star formation in the Galaxy is expected to be roughly the total mass in molecular clouds (2 × 10 9 M ⊙ )d i v i d e db yt h e i rf r e e -f a l lt i m e( 4× 10 6 yr -1 )b u tt h i sg i v e s ah i g h e rr a t e( 5 ×10 2 M ⊙ /yr) of star formation than observed. Thus some other source of support, besides thermal pressure is required to prevent the molecular clouds from collapse. This conclusion is supported by the fact that the widths of spectral lines in molecular clouds are far broader than one would expect from thermal effects alone.

These alternative mechanisms of support might be, magnetic fields (e.g. [Chandrasekhar and Fermi, 1953], [Shu et al., 1987], [START_REF] Mouschovias | NATO ASIC Proc. 210: Physical Processes in Interstellar Clouds[END_REF], rotation (Field 1978), or turbulence, e.g. [Larson, 1981], [Ballesteros-Paredes et al., The overall picture of low-mass star formation 2007]and [START_REF] Ostriker | The Evolving ISM in the Milky Way and Nearby Galaxies[END_REF]. As said in section 1.1.3 turbulence and magnetic fields play a key role in counteracting the gravitational force. Three cases are possible, as follows:

1. Low-mass cores are magnetically subcritical,n a m e l ym a g n e t i ce ff e c t s prevent the collapse and support the molecular clouds. The onset of the collapse began in the densest regions, controlled by ambipolar diffusion. The characteristic timescale of ambipolar diffusion (τ AD ), i.e., the time required for neutrals to move a distance R agains the ions is given by:

τ AD 2.5 × 10 6 yr n H 10 2 cm -3 2 B 3µG -2 R 1pc 2 χ 10 -6 (1.2)
where χ is the ionization fraction, B is the magnetic field and n H is the hydrogen density.

This ambipolar diffusion timescale is several times longer (10 7 yr [Tassis and Mouschovias, 2004]) than the dynamical contraction timescale or free-fall timescale (i.e., collapse governed by self-gravity in the absence of thermal pressure), given by

t ff = 3π 32Gρ c 1/2 (1.3)
where G = 6.674×10 -8 cm 3 g -1 s -2 is the gravitational constant and ρ c the initial density of the core, assumed to be a uniform sphere.

t ff 0.5 × 10 6 yr n 10 4 H 2 cm -3 3/2
(1.4) Thus, the magnetic flux of the cloud is not be lost during the dynamical contraction (free-fall).

2. Star formation controlled by turbulence.S u p e r s o n i c t u r b u l e n c e g e n e rates a complicated shock pattern. Randomly produced shock-compressed regions of high density can occasionally reach the point of instability to collapse. The time scale is shorter than the first one, about 10 6 yr.

3. Cores that are intrinsically stable are induced to collapse by a external event that trigger the collapse, such as a ionization front, a supernova shock wave, or a cloud-cloud collision. The time scale is the shock crossing time, which could be as short as 10 5 yr.

Finally, once magnetic forces have become dynamically unimportant, turbulence has weakened, and rotational effects have become unimportant relative to gravity in the core of the molecular cloud, and the mass of the core exceeds the Jeans mass, the gravitational collapse proceeds, the first seed of the star has been formed.

Protostellar collapse can be divided in two phases (e.g., [START_REF] Larson | [END_REF]). ( 1) Isothermal phase; initially, when the collapsing fragment is optically thin to the infrared radiation of the grains, the temperature stay stable, around 10K. This phase is the result of the balance between the heating of the gas by compression and cosmic rays and, the grain and molecular cooling. The gravitational energy is all radiated away, so the thermal energy is well below gravitational energy. ( 2) Adiabatic phase begins once the gas collapses to the point to reach a density high enough to become optically thick. Now the gravitational energy heat the gas and the pressure increases rapidly, and the collapse slow down. However, once the central temperature has reached about 2000 K, the molecular hydrogen dissociates and the collapse resumes, because much of the gravitational energy is used to dissociated hydrogen molecule and not to increase the thermal pressure. The collapse continues until the dissociation is complete at the center and an amount of material comes into hydrostatic equilibrium, forming a protostellar core surrounded by an envelope of gas and dust.

Accretion phase

This phase involves the collapse of the remaining envelope onto the protostellar core. The principal source of energy of the protostar is the kinetic energy of the material infalling into the accretion shock. The accretion shock is generated by the supersonically infalling material striking the disk orbiting the protostar [START_REF] Hollenbach | From Gas to Stars to Dust,V o l .7 3o fAstronomical Society of the Pacific Conference Series[END_REF].

The presence of a central source of energy produces the heating of the inner

The overall picture of low-mass star formation regions of the envelope. Therefore, the molecules that were frozen onto the grains in form of ice mantles evaporate back to the gas phase forming a region with a gradient in its chemical composition called Hot corino , [ Bottinelli et al., 2004], [Bottinelli et al., 2007]( f o rt h el o w -m a s sc a s e ,t h ec o u n t e r p a r tf o rt h eh i g h -m a s s protostars is called Hot core).

Knowing the density and temperature profiles it is possible to understand which are the processes involved in the star forming process. Dust and density profiles are consistent with free fall motions in Class 0 protostars (e.g., [Chandler et al., 1998], [START_REF] Ceccarelli | [END_REF], [Jørgensen et al., 2002], [Schöier et al., 2002]). Through observation of different species we can constrain models to know the chemical composition across the envelope to understand the physical status of the gas and the history of the protostar.

YSOs classification

During the collapse of the core, the dusty, cool and optically-thick collapsing layers absorb the radiation from the accretion (UV range) and re-radiate in the infrared region of the spectrum, so that the protostar becomes a bright observable infrared source. This is why young stellar objects (YSOs) have been classified observationally by their near infrared-millimeter spectral energy distributions [START_REF] Lada | Star Forming Regions,V ol. 115 of IAU Symposium[END_REF], [Andre et al., 1993]a s ,C l a s s0 ,C l a s sI ,C l a s sI Ia n dC l a s sI I I ,a ss u m m a r i z e di n Figure 1.5.

The youngest observed YSOs are the Class 0 sources, with an age of about 10 4 yr [START_REF] Andre | Protostars and Planets IV[END_REF], [Evans et al., 2009]. This age is based on the statistical detection rate. This kind of sources show (1) Indirect evidence of a central YSO, as indicated by, e.g., the detection of compact cm radio continuum emission or the presence of a collimated CO outflow. (2) Centrally peaked but extended submmillimeter continuum emission tracing the presence of a circumstellar dust envelope.

(3) High ratio of submmillimeter to bolometric luminosity L submm /L bol >> 0.5%, where L submm is measure longward than 350µm. This ratio means that Class 0 protostars are still surrounded by a dense circumstellar envelope more massive than the central protostellar core.

The overall picture of low-mass star formation envelope and thus they are observable in the visible. They correspond to the T-Tauri stars. The youngest members of these classes drive outflows, and all drive strong winds with Ṁ ≈ 10 -7 M ⊙ yr -1 and v w ≈ 200 km/s. When Class II sources are unobscured, they can be placed on the Hertzsprung-Russell diagram and compared with theoretical evolutionary tracks.

Outflow phase

The outflow phenomenon is associated to the star formation process, and it is observed in YSOs of nearly all masses e.g. [Lada, 1985], [START_REF] Shepherd | Galactic Star Formation Across the Stellar Mass Spectrum,V o l .2 8 7o fAstronomical Society of the Pacific Conference Series[END_REF], from below the brown dwarf limit (e.g. [START_REF] Whelan | IAU Symposium,V o l .2 7 5o fIAU Symposium[END_REF]) to the precursor of the ultra-compact HII regions (e.g. [Codella et al., 2004]), and in environments as different as isolated globules ([Fro ebrich et al., 2005]) and cluster-forming regions (e.g. [Zhang et al., 2001]).

As a protostar accretes, the bipolar outflow phase starts, making the collapse possible as it takes away the excess of angular momentum of the accreting material. In fact, even the relatively small-scale molecular cloud cores have far too much angular momentum to be able to collapse to stellar dimensions. Various physical effects can contribute to the solution of this problem but all these mechanisms are still heavily debated, as well as the role of the magnetic fields. One is the fragmentation of the core into a binary or multiple system, where much of the angular momentum goes into orbital motions. Another one is the collapse of the cloud onto a central stellar object surrounded by an orbiting disk, which contains most of the angular momentum. Whatever, it is clear that the removal of angular momentum by outflows is a key process.

In fact, outflows from Class 0 sources are particularly powerful [Saraceno et al., 1996]i nc o m p a r i s o nw i t ht h em o r ee v o l v e dp r o t o s t a r s .

Observationally and theoretically there are several explanations proposed, and still debated about the origin of the outflows. One possibility is that they are formed from the ambient gas entrained progressively by the ejected jet powered by the central protostar. However, the physics of the driving mechanism of the jet is not yet well understood. (To read more about outflows, see section 1.5.1)

The overall picture of the intermediate and high-mass star formation

The current picture of low-mass star formation is quite well established and understood, both, observationally and theoretically. There were attempts to extend the same paradigm to the intermediate and high mass star formation. However, it was found that protostars of higher masses likely evolve differently from those of low mass. There are still a lot of uncertainties on how intermediate and high-mass protostars form.

Global properties

High-mass forming stars should be easy to observe because they are very bright, but observations are difficult because massive star forming sites are far away, they are relatively few in number and the evolutionary time scales are shorter than those of low-mass protostars. Typically, they are found at distances of some kiloparsecs from the Sun limiting the spatial resolution [START_REF] Qiu | [END_REF]. Massive star formation occurs in dense clusters (e.g. [START_REF] Beuther | [END_REF]), so at their large distances its difficult to resolve the individual features of their formation. In addition, massive stars destroy their natal environment via HII regions.

The intermediate-mass (IM) protostars are objects between 2 and 8 M ⊙ .T h e Class II and III sources are classified as Herbig Ae/Be stars (Herbig, 1960). We can differentiate two groups of IM protostars: those with masses between 2 and 5 M (the so called Herbig Ae stars), and between 5 and 8 M (Herbig Be stars). The formation of Herbig Ae stars shows similarities with low-mass or T Tauri stars. In fact, in these objects, the mass accretion phase of the protostar occurs prior to (and independently from) its contraction towards the main sequence. The relatively slow evolution during some tenths of million years simplify the study of these objects, since they soon disperse the parent cloud and become visible. In contrast, Herbig Be protostars are usually found in clusters, highly embedded in massive and chemically complex molecular clouds [Alonso-Albi and [START_REF] Alonso-Albi | Proceedings of the Annual meeting of the French Society of Astronomy and Astrophysics[END_REF].

The overall picture of the intermediate and high-mass star formation

Formation schemes

Theoretically, the origin of the problem in the high-mass star formation is that the Kelvin-Helmholtz time scale is shorter than the infall time scale. This means that the protostar reaches the main sequence while accretion is still going on.

The resulting nuclear burning produces a huge luminosity and a corresponding large radiation pressure force on dust grains in the collapsing dense core and it can be large enough to becomes more important than gravity force and reverse the infall of matter, once a the star reaches a mass of roughly 20 -40 M ⊙ [Wolfire and Cassinelli, 1987] for typical Galactic metallicities. However, accretion can still be accomplished through a high accretion rate generated by the turbulent environment [McKee and Tan, 2003] and also through disks, e.g. [Krumholz et al., 2005a]w i t ho u t fl o w st h a tg r e a t l yr e d u c et h er a d i a t i v ep r e s s u r e [ Krumholz et al., 2005b].

The difficulties associated with forming massive stars via accretion resulted in new theories, like that massive stars might form by collision and coalescence of lower-mass stars in the clusters [START_REF] References Bonnell | Hot Star Workshop III: The Earliest Phases of Massive Star Birth,V o l .2 6 7o fAstronomical Society of the Pacific Conference Series[END_REF].

It was earlier suggested that massive stars form in magnetically supercritical molecular cloud cores [Shu et al., 1987], while low-mass stars form in subcritical cores by ambipolar diffusion. Although magnetic fields are undoubtedly important in both mass ranges, there is no strong observational evidence that support this differentiation [Beuther et al., 2007]. Rather, it is becoming apparent from observations [Beuther et al., 2007], [Plume et al., 1997]t h a tm a s s i v es t a r sf o r m in a highly turbulent environment from high-mass, high-density (n ≈ 10 6 cm -3 ) cores, while low-mass stars form from relatively quiescent lower-mass cores of density n ≈ 10 5 cm -3 .

The protostellar evolution time is roughly the same, between high-mass stars (≈ 10 5 yr) and low-mass stars (3×≈10 5 yr); thus Ṁ for accretion of high mass stars must be significantly greater than that for low-mass stars. This requirement is apparently consistent with the formation of high-mass stars in very turbulent high-density regions; this initial condition naturally provides a high mass accretion rate. Bipolar outflows have been observed in the high and intermedia-mass proto-stars. High-angular resolution observations show that, in general, intermediatemass outflows are intrinsically more energetic than those driven by low-mass YSOs ( [START_REF] Beltrán | [END_REF]). They appear collimated, even at low velocities, and have properties that do not differ significantly from those of low-mass stars ( [Beltrán et al., 2002]). Their formation, like in the low-mass case, seems to be due to the magnetic interaction near the star-disk interface with the radiation pressure. Circumstellar disks around intermediate-mass protostars have been observed at high angular resolution in the continuum and in the CO rotational lines [Fuente et al., 2006]. [START_REF] Alonso-Albi | [END_REF]c o m p l e t e das e a r c hf o rc i r c u mstellar disks in a sample of six HBe stars: R Mon, MWC 1080, MWC 137, MWC 297, Z CMa, and LkH 215.

Chemical properties

The high-mass stars (>9M ⊙ )a r ee x t r e m e l yi m p o r t a n tf o rt h eg a l a c t i cc h e m i c a l evolution and the physics of the interstellar medium. Stars in this mass range can evolve into Type II supernovae which produce and release a significant amount of heavy elements that later will be the cooling agents in the interstellar gas, promoting the formation of later generations of stars.

The chemical features of the high-mass protostars are also notable. Unlike the hot corinos, their low mass counterparts, the hot cores around massive protostars have lower abundances in O-and N-bearing complex organic molecules compared to that of methanol and formaldehyde by about one order of magnitude [Bottinelli et al., 2007]. Additionally, the deuteration, large in low-mass hot corinos, is observed and found to be lower in high-mass star forming environments, like Orion.

Even though it is clear that high, intermediate and low-mass star formation differs, there is not a cutting border between them.

Besides, there is not a sharp change in the slope of the initial mass function in the high-mass region compared to 1-10 M ⊙ region, although the initial mass function (IMF) is difficult to determinate.

Observations of outflows

In the middle of the 1980's, radio astronomers carried out the first large-scale molecular outflow surveys towards a great number of IRAS sources in extensive regions of the sky (e.g., [Margulis and Lada, 1985], [START_REF] Fukui | [END_REF], and [START_REF] Armstrong | The Physics and Chemistry of Interstellar Molecular Clouds -mm and Sub-mm Observations in Astrophysics,V o l .3 3 1o fLecture Notes in Physics[END_REF]) This was important in estimating the energy deposited by molecular outflows in the parental clouds, and concluded that outflows have a significant impact on the dynamical state of their host cloud.

The YSO outflows strongly interact with the surrounding molecular cloud, injecting momentum and energy at scales ranging from a few AU up to few parsecs away from their powering sources. These ejection phenomena may reveal themselves in different ways.

In the optical wavelengths, it has been observed nebulosities known as Herbig-Haro (HH) objects. They are often seen in pairs or in a chain of knots with a highest surface brightness. These HH flows are observed to be very collimated and sometimes connected by fainter emission, have length-to-width ratios of 10:1 or more, and are normally referred to as HH jets [Reipurth and Bally, 2001]. As these collimated jets penetrate the ambient medium, they transfer momentum and accelerate the surroundings.

Beyond the collimated jet the flow continues up to the bright and more extended bow-shaped HH objects, that can be located at distances of about 0.1 to 1pc [START_REF] Arce | Protostars and Planets V[END_REF] from the central driving source (Fig. 1.10).

The optical emission comes from the shocks arising from the interaction of ah i g h -v e l o c i t yfl o wo fg a se j e c t e db yay o u n gs t e l l a ro b j e c ta n dt h ea m b i e n t medium. These shocks are traced by the collisionally excited Hα,whileforbidden line emission, such as [O I], [N II], [S II], and [Fe II] in the near-IR, trace the sites behind the shock when the gas cools and recombines.

Radio interferometric observations can disentangle the morphology and dynamics of the different elements that are present i.e., to resolve the detailed structure of the low-velocity gas around the optical jet.

For example, in the HH 111 object using PdB observations [Lefloch et al., 2007]o b t a i nav e l o c i t yc h a n n e lm a p so fC O( 1-0 )e m i s s i o nt h a ts h o w st h e emission that comes from the cavities traced by low-velocity gas.

Molecular CO transitions are the most common tracer of outflowing molecular gas (Fig. 1.6). The CO line profile can trace a low-velocity outflow cavity or even show extremely high velocities features that can reveal the presence of a jet (e.g. [Tafalla et al., 2010]). In some cases CO emission can reveal molecular bullets along the axis of the jet probably formed in the shocks resulting from time-variable ejections in the powering source [Raga et al., 2002].

It is useful to observe different molecular lines besides CO because they trace different components of the gas, due to their different excitation conditions, and [Gueth and Guilloteau, 1999]i so v e r l a i do naf a l s ec o l o ri m a g eo fH 2 [McCaughrean et al., 1994]. The red lines delineate the 230 GHz continuum emission from dust concentration around the protostar. (bottom) In color scale we see the near infrared H 2 emission, in contours blue and red there are the blue-and red-shifted emission of SiO , the green contours mark the dust emission at 870 micron surrounding the protostar. [Lee et al., 2007] Observations of outflows also because the molecular abundances can be greatly enhanced behind shocks. Examples of such molecules are SiO, HCN, SO, OH, and HCO + .O t h e rm o l e c u l e s like the molecular hydrogen (H 2 )c a nt r a c es h o c k -e x c i t e dr e g i o n si ft h e r ei sa enough quantity of warm (>200 K) molecular gas present in the surrounding star-jet system (Fig. 1.6). The observations of different chemical species can provide the physical conditions in different regions of the outflow.

The importance of the outflow phenomenon

The outflow phase is thought to play an important role into the star-formation process for two main reasons. First, outflows serve as a way to get rid of the excess angular momentum transferred from the circumstellar disk onto the forming star. Second, they can clear out cavities in their natal molecular clouds, and may eventually cause cloud disruption [Maret et al., 2009].

Finally, the involved kinetic energies (E flow ∼ tens of 10 44 ergs) and momenta (P flow ∼ tens of M ⊙ km s -1 ) are very large and their formation rates so high that the outflows from young stars are an extremely important source of mechanical energy that have major effects on the dynamics and structure of their parent cloud [Reipurth et al., 2000].

H 2 emission

The propagation of a supersonic protostellar wind through its surrounding medium happens primarily via shock waves that compress and heat the gas, which in turn cools down through line emission at different wavelengths. In molecular gas, shock cooling is mainly through line emission. Molecular hydrogen is by far the most abundant specie in the gas phase, and although H 2 emits only through quadrupole transitions with low radiative rates, it often represents the main gas coolant in flows from young protostars [Kaufman and Neufeld, 1996].

In general, the detection of H 2 emission is evidence of the presence of C-shocks, because a shock at high velocity should destroy the hydrogen molecules.

H 2 -shocked emission in outflows has been widely studied in the past (e.g. the Infrared Space Observatory) mainly through its ro-vibrational and rotational emission in the mid-IR, (the pure rotational lines of H 2 S(0) up to S(8) e.g., 2010). Taken from [START_REF] Tafalla | IAU Symposium,V o l .2 8 0o fIAU Symposium[END_REF] it allows to investigate the properties of the region closest to the origin of the outflow.

Subarcsecond angular resolution observations reveal that the centimeter emission traces collimated partially ionized jets extending typically to distances 10-100 AU from the exciting source, corresponding to material ejected from the protostar with dynamical ages of the order of a few years or less.

Outflow chemistry

In addition to the physical interaction of the outflowing material with the surrounding one, chemical changes are seen in the gas that has been accelerated by supersonic wind, and has been shock-processed.

Outflows drive shock waves that compress and heat the gas triggering important chemical processes, such as molecular dissociations, ice sputtering, and dust grain disruption.

Different molecular lines are observed to trace different components of the gas. While the CO line wings trace the bulk of the outflowing material, molecules such C 3 H 2 ,N 2 H + and DCO + trace the quiescent medium [Bachiller et al., 2001]. Other molecules like the SiO, CH 3 OH and H 2 CO lines trace the high-density shocked gas zones.

The SiO molecule is one of the molecules that show a high enhancement in abundance (up to 10 6 ) with respect to the ambient medium, often found close to bowshocks and along the axes of some collimated outflows, where there is the maximum interaction of the outflow with the ambient gas. Very likely, this is due to the sputtering or vaporization of Si atoms from dust mantles and grains and the grain-grain collisions [Guillet et al., 2009]. SiO lines usually present broad wings, and together with CO, reaches the highest terminal velocities among all molecular species.

Other molecules whose abundances are enhanced in outflows are methanol (CH 3 OH) and formaldehyde (H 2 CO). The abundance enhancement is likely due to their release from dust grain mantles.

Molecules larger than CH 3 OH have also been recently identified in L1157-B1. [Arce et al., 2008] detected for the first time in an outflow HCOOCH 3 (methyl formate), CH 3 CN (methyl cyanide), HCOOH (formic acid), and C 2 H 5 OH (ethanol). The detection of these molecules shows that outflow shocks can lead to a very rich chemistry, that in low-mass sources was previously restricted to the hot corino regions in the vicinity of some protostars like IRAS 16293-2422 [Cazaux et al., 2003].

Water is also a sensitive outflow tracer through its thermal and maser emission in outflows. The maser emission is so bright that it can be easily detected across the Galaxy, and the abundance of water required to produce the emission must have been significantly enhanced by shocks (e.g., [Kaufman and Neufeld, 1996], [Cernicharo et al., 1996]). Early observations with ISO revealed thermal water emission from a number of bipolar outflows (e.g., [Nisini et al., 1999], [START_REF] Giannini | The Promise of the Herschel Space Observatory,V o l .4 6 0[END_REF]). Instruments like ISO, SWAS and Odin telescopes had a limited angular resolution, that only allowed to study global averages of the water emission over the whole outflow, or at least over each of their lobes. In the past two years the Herschel Space Observatory has started to provide water data with high spatial resolution thanks to its Heterodyne Instrument for the Far Infrared (HIFI), with velocity resolutions below 1 km s -1 .M u c hw o r ks t i l ln e e d st ob e done to understand the excitation conditions of the H 2 Oe m i t t i n gg a s ,a n dt h i s understanding is critical to extract the full potential of the water data.

Sulfur-bearing molecules are believed to provide a useful chemical clock and allow to study the evolution of outflows. For example, the H 2 S abundance, once the molecule is ejected to the gas phase by the effect of shocks, will decrease rapidly after about 10 4 yr due to oxidation with O and OH, which produces SO and SO 2 .T h e r e f o r e ,S O / H 2 Sa n dS O 2 /H 2 Sr a t i o sm a yb ep a r t i c u l a r l yu s e f u lt o eL 1 1 5 7o u t fl o w . right:M a p so fi n t e g r a t e di n t e n s i t yf o rd i ff e r e n t molecular species towards the blue (southern) lobe of the L1157 outflow. Note the different extension and peak position of each species. Taken from [Bachiller and Perez Gutierrez, 1997].

estimate the relative age of shocks in an outflows (e.g., [Wakelam et al., 2005], [Buckle and Fuller, 2003]).

The amount of the abundance enhancement depends on the species under consideration, from factors of a few for species like HCO + to several orders of magnitude for species like CH 3 OH and SiO (Fig. 1.8). Not all the outflows show the same chemical richness. The ones that shows more chemical activity have been called "chemical active"o u t fl o w s [ Bachiller et al., 2001].

The exact cause of the chemical richness is however unclear. One possibility is that the youngest outflows interact with denser envelopes, and therefore give rise to stronger shocks. Another possibility is that early chemical activity is associated with the higher energy known to characterize the earliest outflow phases (e.g. [Saraceno et al., 1996]; [Bontemps et al., 1996]). Among the group of chemically active outflows is the L1157 outflow.

To summarize, millimeter spectral line surveys are the best tool to explore the chemical composition of the vicinity of YSOs. The observational study of molecular outflows at different scales provide a huge amount of information, not only about the physics and chemistry of the jets and outflows themselves, but also about the young stellar objects powering them and the molecular environment of the systems. Observations towards different regions allow to distinguish variations in the excitation conditions and/or chemical composition along the jet/outflow.

Herbig-Haro objects

The Herbig-Haro (HH) are nebulous knots, observed in optical images, which are distinguished by their spectrum with characteristic emission lines, especially some forbidden transition of oxygen, sulfur and neon ions.

They were discovered in the 50s by George Herbig and Guillermo Haro, believing then, that they were nebulae that contained a star in formation. Several years after their discovery, Herbig suggested that HH objects could be associated in some way with winds from young stars (Herbig 1958).

The optical spectrum of the HH objects can be reproduced if the emission is modeled as high-density gas cooling after having pass through a shock wave, i.e, the observed emission comes from the recombination zone which is located behind the shock wave. By comparing detailed observations with the theoretical predictions, it is found that it is better reproduced the HH spectra using bow shock models.

The gas that comes just at the apex of a bow shock have a larger velocity with respect to the wings, generating strong forbidden lines (e.g [OIII] 5007 to [SII] 6312 to [NeIII] 3869 Å). The gas that passes through the wings collides with a slower velocity gas producing only low excitation lines (e.g., the Lyman of H 2 molecule). This spectrum is probably produced by the interaction between a strong wind or high velocity jet (100 -400 km s -1 )a n dt h es u r r o u n d i n gg a s .

In addition, it has been observed that HH objects are often (but not always) in pairs, one redshifted and the other blueshifted, symmetrically arranged with Observations of outflows respect to the position of the powering star. However, in general, the redshifted HH object is much weaker than the blue counterpart. This indicates that the blueshifted part is closer to the surface of the cloud respect to the observer, so it suffers less extinction in the visible.

There are about 400 known Herbig-Haro objects [Reipurth, 2000].

On the other hand, the jet phenomenon itself is directly related to HH objects. The jets appear as compact chains of aligned knots, sometimes connected by weak emission and have a length/width ratio of 10:1 or more. In some cases both the collimated jet as well as the shock fronts are visible, sometimes only the shock is seen.

HH jets are always found to originate very close to their power source (less than 10 2 to 10 3 AU from their source), and they usually extend less than 0.2 pc in length. These jets present velocities from about 100 km s -1 up to 600 km s -1 .

There are two main theories for the origin of the structure of knots and the nature of the jets with a low excitation spectrum. The first mentions a series of oblique shocks traveling through the jet and perturbed by hydrodynamic, magnetohydrodynamics or thermal instabilities. The second theory considers that jets are not a continuous phenomenon.

Emission lines observed in the jets have properties of the ionized gas features, exited by shocks. Optical jets are usually monopolar and radial velocities observed are usually blueshifted by several hundred km s -1 .T h i si si n t e r p r e t e da si nt h e case of HH objects, as a consequence of the extinction effect.

Finally, the most important evidence relating HH objects with outflowing material comes from measurements of their proper motions. The velocities of the measured proper motions can reach 200 to 300 km s -1 . In the case of dual systems of HH objects, they have symmetrical motions, moving in the opposite direction with respect to the position of the central source successfully.

To conclude, the gas ejected by young stars can b e detected and studied in various ways, according to the spatial scale to be studied: very close to the star the radio continuum (thermal jets) are detected, at intermediate distances molecular emission lines in radio and infrared (bipolar outflows), and far from the star as it is seen in the visible like HH objects.

Theory of the outflow phenomenon 1.5.1 Origin and structure of the outflows

There are many theories about how outflows and jets are originated, and which is the relation between them, but it is still matter of discussion. However, the observed correlation between the mass outflow with the accretion rate into the disk seem to favor some magnetohydrodinamic (MHD) jet launching mechanism [Cabrit et al., 1990]. In the case of the jet, it can be generated from the interface between the star's magnetosphere and the disk (the called X-wind model) [Shang et al., 2002]orfromawiderangeofdiskradii(theD-windmodel) [START_REF] Pudritz | Protostars and Planets V[END_REF].

Actually, it is not clear yet what physical mechanism produces the outflows, and the fundamental stellar or protostellar wind that should sweep up the fast moving molecular gas is very hard to detect. However, there are three empirical correlations that can be summarized from many observations:

1. Progressively with the age, the outflow activity declines. It exists a tight or intermediated-mass system, but, it seems that massive stars are very inefficient accreting material and that is reflected in the higher outflow rates.

3. The degree of collimation is directly determinate by the outflow rate. This correlation suggest that the apparent decrease in outflows energetics with the age reflects a decay in the mass accretion and infall rate [Bontemps et al., 1996].

Outflow Models

Several outflow models have been proposed to explain how molecular outflows from protostars are formed. Outflow models can be separated in four principal groups: Wind driven shells, jet driven bowshocks, jet driven turbulence flows and circulation flows. In the first three groups the molecular outflow is considered as ambient material entrained by a wide angle wind or accelerated by a collimated jet. Conversely the last one considers that the outflows is formed by deflected infalling gas.

1. Wind driven shell models.I nt h i sm o d e l ,aw i d e -a n g l er a d i a lw i n db l o w s into the surrounding ambient molecular cloud, forming a thin swept-up shell that corresponds to the outflow shell (e.g. [Shu et al., 1991], [Matzner and McKee, 1999]). In these models it is often assumed that the ambient material has a toroidal form. These kind of models are attractive because they explain old outflows of large lateral extents and low collimation.

2. Jet bowshock model. In the jet driven bowshock model, a high-velocity and well-collimated internal jet entrains the ambient molecular gas through a bowshock. [Raga and Cabrit, 1993] As the jet hits into the ambient material, ap a i ro fs h oc k s ,aj e ts h oc ka n dabo w s h oc k ,a p pe a ri nt h eh e a do ft h ej e t .

Between the shocks the gas pressure is enough high that material is ejected sideways out of the jet beam, interacting with the ambient gas through a broader bow shock surface, producing an outflow shell around the jet.

3. Turbulent jet model.I nt h ej e t -d r i v e nt u r b u l e n tm o d e l ,t h ei n t e r a c t i o nb etween the jet and the surroundings environment produces Kelvin-Helmholtz

Theory of the outflow phenomenon instabilities along the jet/environmental boundary that lead to the formation of a turbulent viscous mixing layer, through which the cloud molecular gas is entrained. The mixing layer grows both into the jet and into the environment, and eventually the whole flow becomes turbulent [Canto and Raga, 1991].

4. Circulation models.I nt h ec i r c u l a t i o nm o d e l ,t h em o l e c u l a ro u t fl o wi sn o t formed by entrained material by a jet or stellar wind. Instead, it is formed by infalling matter that it is deflected away from the protostar in a central torus of high MHD pressure through a quadripolar circulation pattern around the protostar, and accelerated above escape speeds by local heating (e.g., [Fiege and Henriksen, 1996a], [Fiege and Henriksen, 1996b]). The molecular outflow can have some contribution by entrainment from winds or a jet, but their contribution are limited to the polar regions [Lery et al., 2002].

The current observations at high angular resolution allow to confront the outflow models in detail. Comparing the jet-diven bowshock models vs the wideangle wind models, there are features that one model can reproduce and that the other can not. However, there are some models that try to unified both pictures, in which both the jet and the wind play a role in driving the outflow e.g. [Shang et al., 2006].

The jet-driven bow shock models are characterized by an outflow velocity that increases toward the position of the internal leading bowshock, giving the appearance of a peak curve in the PV (position velocity) diagram, (Fig. 1.12). These models are capable to produce the observed outflow width for high collimated outflows, basically because the shocked gas in the bowshock working surface limits the transverse momentum, that can be released to the surroundings. However, these models have problems to reproduce the observed width of poorly collimated outflows. Also, there are some difficulties to produce the right momenta. In addition, the jet-driven model may need jet wandering to produce the observed outflow width and momentum for most of the outflows.

In the wide-angle wind models the outflow velocity also rises near to the source, but, in this case, they exhibit a lobe PV structure (Fig. 1.12)M o s to f reduced to a value that is subsonic with respect to the local sound speed (see Fig. 1.14). Downstream behind the shock front, the gas temperature falls, first quickly and then slowly, in an "relaxation region". It is the radiation generated in this cooling zone that provides the information about the shocks in star forming regions.

Following [Hollenbach and McKee, 1989], the figure 1.13 shows the typical temperature structure of a J-shock. First, the preshock gas experiences the radiation field from the shock front, the gas molecules are ionized and dissociated collisionally in the immediate postshock environment (T ≈ 10 5 K) and UV photons arise affecting both preshock and postshock gas. Downstream, at a hydrogen column density ≈ 10 17 cm -2 from the shock front the gas temperature cools until T ≈ 10 4 K and the recombination starts. Temporarily an equilibrium is reached between recombinations and photoionizations by the UV photons produced upstream. This is the so-called "recombination plateau" where the Lyman continuum photons are absorbed, maintaining a T ≈ 10 4 K. Once the ionizing photons are absorbed (at N ≈ 10 18-19 ), the gas cools, recombines, and the molecular chemistry begins.

First to appear is H 2 ,i n i t i a l l yp r o d u c e db yt h eH -formed out of the residual electron supply. Later, H 2 formation continues on grain surfaces. The H 2 molecules may be vibrationally excited when injected back into the gas phase. If the density is at least 10 5 cm -3 , they decay collisionally. Under these conditions, the formation of H 2 becomes the major heating source for the gas, stabilizing

Theory of the outflow phenomenon

the gas temperature at a second plateau of about 500 K until all the hydrogen becomes molecular.

Other molecules also form, through purely gas-phase processes. The low ionization level and prolonged temperature elevation of the plateau result in neutralneutral reactions being efficient.

An e t w o r ko fs u c hr e a c t i o n si sa c t i v a t e dw h i c hy i e l d sav a r i e t yo fs p e c i e s , including the important coolants: CO, OH and H 2 O.

In addition to these gas-phase processes, grain cores and their mantles are affected by the shock. Millimeter and infrared observations of these shocked regions reveal the abundance enhancement of species like SiO or Fe [Nisini et al., 2007], which are usually highly depleted from the gas in quiescent regions.

This can be explained by models in which dust is destroyed (vaporization and sputtering) in the shock waves initiated by jets and outflows.

They are indeed expected to exist at the apex of bow shocks around YSOs where the observed strong lines of Fe II can be explained with J shocks modelling [Hartigan et al., 2004].

The models of [Guillet et al., 2009] show that dust destruction occurs in Jtype shocks at levels comparable to that of C-type shocks of the same velocity, and that this destruction is dominated by vaporization in grain-grain collisions.

C-shock

In C shocks, the fluid variables are everywhere continuous, they change smoothly throughout a broad transition region by dissipative effects.

Ct y p es h o c k sf o r mw h e nt h ei o n i z e df r a c t i o ni sl o ws ot h e i rc o u p l i n gw i t h neutral is weak. The electron and the ion flow velocities and temperatures can differ appreciably from those of the neutrals. The charged particles that are upstream function as a magnetic precursor. The radiative dissipation process of this precursor can be so important that the neutral shock becomes a continuous transition. Then, the structure of these shock waves is dependent on several factors. ( 1) Ionization fraction through the ion-neutral collisions. ( 2) The rates of cooling through the various molecular coolants of the gas; and (3) the abundances of those coolants. [Kaufman and Neufeld, 1996] For typical ionization fractions in dense clouds x(e) ≤ 10 -6 ,t h eC -s h o c k s occur at V shock ≤ 40-50 km s -1 , causing peak temperatures of typically 2000-3000 K [Draine and McKee, 1993].

These shocks are characterized by bright molecular emission as first emphasized by [Hollenbach and McKee, 1979]. Molecules generally survive their passage, although the chemical composition of the gas can be significantly altered. On the other hand, the heating and compression of the region trigger different microscopic processes such as molecular dissociation, endothermic reactions, ice sublimation, and dust grain disruption which do not operate in the unperturbed gas.

Grain-grain destructive processes (sputtering) in C-type shocks have been extensively studied in order to confront shock models with line measurements of tracers such as SiO, H 2 Oa n dC H 3 OH in molecular flows around YSOs (e.g., [START_REF] Jiménez-Serra | [END_REF]; [Gusdorf et al., 2008].)

Recently, [Guillet et al., 2011]modeledtheimpactofthegrain-graincollisions in C-type shocks propagating in dense clouds (n H ≥ 10 4 cm -3 )o nt h ee v o l u t i o n of the dust distribution, the shock dynamics and the release of chemical species into the gas phase such as SiO. They found that dust destruction by grain-grain collision (vaporization) dominates over that by gas-grain collisions (sputtering) at high densities (n H ≥ 10 5 cm -3 ), but is negligible at low densities (n H ≤ 10 4 cm -3 ).

Chapter 2 Radio observations

At present, more than 150 gaseous molecules have been detected in the interstellar medium (ISM), including alcohols, ketons, ethers, aldehydes, and others. They are mainly observed via their rotational transition, e.g. CO, CS, HCO + and/or ro-vibrational transitions, e.g. H + 3 ,H 2 .D e p e n d i n go nt h em a s so ft h em o l e c u l e and its structure, these transitions lie in the radio (cm) to far infrared (>100µm) wavelengths. Heavier molecules have the rotational transitions at longer wavelengths. The transitions can be observed either as emission or absorption lines, depending on the continuum at the transition wavelength. Indeed, the same transition can be in emission in some regions and in absorption in others.

These line spectra can be observed using millimeter-wave single-dish telescopes on the ground or, in order to increase resolution, interferometers and, to avoid the atmosphere effects, space-borne telescopes.

The first steps of Radioastronomy

In 1931, Karl Jansky was conducting experiments for Bell Laboratories related to the production of radio-wavelength noise from thunderstorms, in the first transatlantic radio telephone between the U.S. and London.

He built an antenna, (2.1), designed to receive radio waves at a frequency of 20.5 MHz (wavelength about 14.5 meters). It was mounted on a turntable that allowed it to rotate in any direction, earning it the name "Jansky's merry-go-

The first steps of Radioastronomy round". By rotating the antenna, one could find what the direction was to any radio signal.

During the course of his investigation, Jansky discovered that the noise originated from three sources: local thunderstorms, distant thunderstorms and a constant noise of unknown cause. Over time, he realized the noise gradually changes of position, as if it came from a fixed direction in the sky. He also noted that the signal was detected 4 minutes earlier each day, as the stars do, due to the translational motion of the Earth. This fact made Jansky suppose that the signal did not have a terrestrial origin.

By 1935, he had correctly concluded that much of the signal he was measuring originated in the plane of the Milky Way, with the strongest emission coming from the constellation Sagittarius, which lies in the direction of the center of our Galaxy. Jansky's pioneering work represented the birth of radio astronomy,a whole new field of observational study.

During the years of the Second World War radio physics had made great progress, mainly due to efforts directed towards the development of sensitive and efficient radar equipment.

On the other hand, the international scientific community lost almost all communications. The exchange of information was minimum, with a great exception. A copy of the article published by Reber (who followed the Jansky's work) in the Astrophysical Journal arrived at the Leiden Observatory during the occupation of Holland. This publication catched the attention of Dr. Oort, who immediately envision the possibility of radio astronomical observations. Oort was one of the few astronomers who recognized the scientific potential of the radio data, previously regarded only as a curiosity in the field of engineering. He understood that observing emission lines of species that populate interstellar space would provide a tremendous tool for determining not only the quantity of the species but also of the energy necessary to excite it, and the information on the radial velocity, from which to derive the dynamics of the Galaxy.

Oort commissioned a young astronomer, Hendrick van de Hulst, the estimation of the probability of observing emission lines in radio wavelengths. Van de Hulst, after studying the issue, concluded that the element the most likely to be detected in interstellar space was neutral hydrogen. The frequency of the emis- sion, caused by the transition between the two hyperfine sublevels of the ground state of HI, is 1420 MHz (λ =21cm). Thelifetimeforeachexcitedatomisabout 11 million years so that the probability of spontaneous transition is very low, but this could be offset by the large number of atoms in the interstellar medium and the frequent collisional excitations and desexcitation between the atoms, so such emissions could be observable.

Once in 1945 the predictions by Hendrick Van de Hulst were published, it started a race between several countries to detect the 21 cm line of HI. The detection occurred almost simultaneously (with differences of a couple of weeks) in 1951 in three places: Harvard, Leiden and Sydney. Thereafter, radio astronomy acquired enormous importance which resulted in several projects of radio telescopes in the world and an explosive growth in the amount of information that was collected. The line of HI could be detected in almost any direction from the sky [START_REF] Verschuur | The Invisible Universe Revealed, The Story of Radio Astronomy,S p r i n g e r -V e r l a g43 von[END_REF].

We will not follow this historical development any further, except to note that the instrumental development has been toward higher sensitivity, shorter wavelength, and higher angular resolution.

The radio -far infrared window covers from λ ∼10 m to shortward of λ ∼ 0.1 mm.

Some basic concepts

Electromagnetic radiation in the radio is a wave phenomenon, but when the scale of the system greatly exceeds the wavelength of radiation, we can consider the radiation to travel in straight lines called rays.

As i n g l er a yc a r r i e se s s e n t i a l l yn oe n e r g y ,s oi ti sn e c e s s a r yt oc o n s i d e rt h e energy carried by sets of rays, which differ infinitesimally from the given ray. Consider a set of rays from the sky falling through an horizontal area σ at the surface of the earth. The infinitesimal power dP from a solid angle dΩ of the sky incident on a surface of area dσ is

dP = I ν cos θ dΩ dσ dν (2.1)
where dP =i n fi n i t e s i m a lpo w e r ,i nw a t t s , dσ =i n fi n i t e s i m a la r e ao fs u r f a c e ,c m 2 , dν = infinitesimal bandwidth, in Hz, θ =a n g l ebe t w e e nt h en o r m a lt od σ and the direction to dΩ, I ν =b r i g h t n e s so rs pe c i fi ci n t e n s i t y ,i nWm -2 Hz -1 sr -1 . Equation 2.1 is the definition of the brightness I ν .T h et o t a lfl u xo fas o u r c ei s obtained by integrating 2.1 over the total solid angle Ω s subtended by the source (See fig. 2.2).

S ν = Ωs I ν (θ, φ) cosθ dΩ (2.2)
where the flux density S ν is measured in units of W m -2 Hz -1 [START_REF] Kraus | Radio Astronomy[END_REF].

Antenna and brightness temperature

The radio telescopes measure the power coming from the sky per unit area per unit frequency, or as we already see the flux density.

When the antenna is pointed toward a cosmic source this provides the antenna temperature,T a , defined as the temperature of an equivalent resistor which would give the same power, as measured at the final output terminals, as the cosmic source.

When we point the antenna on a source and then away from it, the difference in the two outputs gives the contribution from the source alone. This provides the brightness temperature,T b , defined as the Rayleigh-Jeans temperature of an equivalent black body which would give the intrinsic temperature of the source.

Note that the antenna temperature is antenna-specific, and will vary from telescope to telescope. They are usually easy to measure, and are useful for understanding the system performance. In contrast the brightness temperatures is a property of the sources on the sky; provided the sources are resolved, different radio telescopes will measure the same value.

The concepts of antenna and brightness temperatures allow us to calculate the power received by single dishes without knowing the details of the Fourier coverage of the aperture. The consequences of the grading are taken into account through global factors, the most important being aperture efficiency and beam efficiency. The aperture ap is defined as

ap ≡ A e /A (2.4)
where A e is the effective collecting area, and A is the geometric area of the antenna.

The effective beam efficiency,B ef f ,c o r r e s p o n d st ot h ea m o u n to fs i g n a lt h a t is contained in the main lobe and is defined by

B ef f ≡ T a /T MB (2.5)
where T a is the antenna temperature corrected for atmospheric extinction, it is: T a =T a e τν and T MB is the main beam temperature.T h em a i nb e a mt e m p e r a t u r e relate the forward efficiency (the the total amount of the signal that comes from the front of the telescope) with the beam efficiency as follows:

T MB = T a F ef f B ef f (2.6)
Under the assumption that both the source and the beam are Gaussian shaped, the brightness temperature, T b ,i sr e l a t e dt ot h em a i nb e a mt e m p e rature, T MB (see Eq. 2.9), by:

T b = T mb θ 2 s + θ 2 MB θ 2 s (2.7)
The aperture efficiency is usually measured via a celestial source of known flux density, whose antenna temperature is obtained by comparison with a noise standard, such as the difference between hot and cold loads. The beam efficiency may bee calculated from the beam pattern [START_REF] Wilson | [END_REF].

The beam

The sky brightness is in general a function of the angle. This may be expressed explicitly by I(θ, φ). Thus, the spectral power or power per unit bandwidth, becomes

dP = I(θ, φ) cosθ dσ dΩ (2.8)
where it is considered the case where dσ is the flat horizontal surface of a receiving antenna with the power pattern of the antenna directed toward the zenith (θ=0).

In this case the equivalent area is called the effective aperture of the antenna. The power pattern P n (θ, φ)isameasureoftheresponseoftheantennatoradiationas afunctionoftheanglesθ and φ.I ti sad i m e n s i o n l e s sq u a n t i t ya n di sn o r m a l i z e d at its most sensitive direction (ideally, this will be along the physical axis of the antenna).

P n (θ, φ)= 1 P max P n (θ, φ) (2.9)
The pattern solid angle Ω A or more commonly called beam area or solid angle of an antenna is

Ω A = 4π P n (θ, φ) dΩ (2.10)
this is measured in steradians (sr). The integration is extended over the full sphere 4π, such that Ω A is the solid angle of an ideal antenna having P n =1 for all of Ω A and P n =0e v e r y w h e r ee l s e . S u c ha na n t e n n ad o e sn o te x i s t ;f o r most antennas the normalized power pattern has considerably larger values for a certain range of both θ and φ than for the remainder of the sphere. This range is called the main beam or main lobe of the antenna; the remainder are the side lobes (See fig 2 .3). In analogy to 2.10,i tc a nb ed e fi n ea st h emain beam solid angle Ω MB by Ω MB = main lobe

P n (θ, φ) dΩ (2.11)
For a Gaussian shap ed main b eam, the solid angle is given by

Ω=1.133 θ 2 b (2.12)
where θ b is the full width to half maximum (FWHM), also know as full width to half power (FWHP). The full width half maximum of the main beam defines the resolution of the telescope and therefore it is important to have an accurate measurement of the telescope's main beam.

If the beam is convolved with a gaussian source of halfwidth θ s ,t h er e s po n s e pattern will also be gaussian, with halfwidth θ r given by

θ 2 r = θ 2 s + θ 2 b (2.13)
(in practice, this relation can be trusted only for θ s >θ b /2) [START_REF] Wilson | [END_REF].

Since the response of the telescope is a convolution of the telescope's beam with a source's signal, a bright and isolated point source should be used to measure the shape of the main beam. Using such an object allows one to be confident area and efficiency of the antenna and the sensitivity of the radio receiver used to amplify and detect the signals.

When the wavefronts reaches the antenna the input signal has to be amplified, because cosmic radio sources are weak, and modulated at a certain frequency that matches with the electronics of the telescope system.

The equipment behind the antenna feed is usually formed by a receiver and a pre-amplifier, followed by a mixer to obtain an intermediate frequency (IF). In millimeter astronomy, because of the absence of good amplifiers at the radio frequency, the signal from the sky is fed directly to the mixer. The output of the mixing of the radio frequency (RF) signal with a local oscillator signal passes through one or several IF amplifier stages, and is then fed to a backend detector, whose output goes to a computer.

The noise of the radiometer is characterized by a receiver temperature, T R , the temperature of an equivalent resistor with the same noise power, as given by the Nyquist formula, P = kT R ∆ν. Normally, the receiver temperature consists of contributions from the mixer plus the IF stages, the latter being multiplied by the conversion loss, L, of the mixer, T R =T M +L T IE .

The system temperature is the temperature of an equivalent resistor which would give the same noise power as the entire system, that is, the receiver, and the contributions received by the antenna from the sky and the ground:

T sys = T R + T sky + T ground (2.15)
where

T sky = F ef f T amb (1 -e -τν ) (2.16)
and

T ground =(1-F ef f ) T amb (2.17)
where T amb is the ambient temperature (assumed to be the same for the air and the ground, in this example), and F e ff is the forward efficiency of the antenna. The r.m.s sensitivity, ∆T a ,o ft h es y s t e mi sg i v e nb y ∆T a = T sys (∆ντ) -0.5 (2.18) where ∆ν=b a n d w i d t hp e ro b s e r v i n gc h a n n e l ,a n dτ is the integration time on source. The signal to noise ratio S/N is then define as

S/N = T a T sys √ ∆ντ (2.19)
2.3 Observing below the atmosphere

Atmopheric transmission

The Earth's atmosphere below 80km is a mixture of gases with the principle constituents nitrogen (N 2 :7 8 . 0 9 %b yv o l u m e ) ,o x y g e n( O 2 :2 0 . 9 5 % )a n da r g o n (Ar : 0.93%). This mixture is almost constant in the lower atmosphere, but there are several minor constituents whose relative percentage may vary both with altitude and time.

The most important of these is water vapor (H 2 O). Its concentration varies erratically with the local weather conditions and with altitude. Carbon dioxide (CO 2 )withana v eragepercen tageof0.03%sho wsbothseasonalv ariationsanda secular trend. In recent years it has come to prominence in connection with the Observing below the atmosphere greenhouse effect.

The atmosphere absorbs electromagnetic radiation particularly at ultraviolet, infrared, X-ray, and gamma-ray wavelengths, while it is mostly transparent in the visible, near-infrared and radio-submm wavebands. Thus, ground-based astronomy is confined to those atmospheric windows.Then the ground-based astronomy observations is constrained to these windows (See Table 2.1 for a list of telescopes in the radio-FIR range).

The visible window is relatively narrow in terms of frequency or wavelength from 0.3 to about 1 µm. The absorption at the short-wavelength side of the optical window is mainly caused by oxygen and ozone (O 3 ). (See Fig. 2.5).

As the wavelength of the radiation approaches X-rays, atomic absorption dominates, and in a very hard X-ray and gamma-ray region, nuclear absorption is the most important.

Ta b l e 2 . 1 : The largest radio-mm-submm telescopes and interferometers to observe in radio to submm wavelengths. ground-based radio observations and refracts the signal and there are changes in the path length. Usually these effects change slowly with time, but there can also be rapid changes such as scintillation and anomalous refraction. Thus, the propagation properties of the atmosphere and detailed features of its radiation must be taken into account, if the astronomical measurements are to be interpreted properly.

Single-dish

Calibration procedure

In the millimeter and submillimeter wavelength range, the atmosphere has a large influence because the atmospheric absorption at these wavelengths can be significant, and varies with elevation. Therefore, we must make accurate corrections to obtain useful calibrated data. In addition, in the mm range, most large telescopes are close to the limits caused by their surface accuracy, so that the power received in the error beam may be comparable to that received in the main beam. Thus, one must use a appropriate value of beam efficiency.

In order to solve this atmospheric problem it is useful to take the sky itself as a calibrator source, as its signal also varies with elevation in the same way as the attenuation, so such a calibration yields the source temperature outside the atmosphere.

This calibration method, often called the chopper wheel method [Penzias and Burrus, 1973], defines the calibration signal to be the difference between an absorber at ambient temperature and the sky.

The procedure consist of: (1) the measurement of the receiver output when an ambient (room temperature) load is placed before the feed horn, and (2) the measurement of the receiver output, when the feed horn is directed toward cold sky at certain elevation. For (1), the output receiver while measuring an ambient load, T amb ,i sV amb :

V amb = G(T amb + T R ) (2.20)
For step (2), the load is removed; then the resp onse to empty sky noise, T sky and receiver cabin (or ground), T ground from eqs. 2.16 and 2.17,i s

V sky = G[F ef f T sky +(1-F ef f ) T ground + T R ] (2.21)
where F ef f is the forward efficiency, this is basically the fraction of power in the forward beam of the feed. G is the varying gain factor to be calibrated out, T R is the receiver temperature. If we take the difference of V amb and V sky ,w eh a v e

∆V cal ≡ V amb -V sky = GF ef f T amb e -τν (2.22)
where τ ν is the atmospheric absorption at the frequency of interest. The response to the signal received from the radio source is

∆V sig = GT A e -τν (2.23)
Hence, solving for T

A , T A = ∆V sig ∆V cal F ef f T amb (2.24)
where T

A is the antenna temperature of the source outside the earths atmosphere. We define

T * A ≡ T A F ef f = ∆V sig ∆V cal T amb (2.25)
The quantity T * A is commonly referred to as the corrected antenna temperature. It can be thought of as a "forward-beam brightness temperature". It is the brightness temperature of an equivalent source which fills the entire 2π sterradians of the forward beam pattern. Note that although the chopper wheel method corrects for atmospheric attenuation, it does not correct antenna temperature for all telescope losses. In particular it does not correct for the very significant telescope losses due to surface irregularities [START_REF] Downes | Evolution of Galaxies: Astronomical Observations[END_REF].

For an extended source filling a large part of the sky T

A is the T MB .F o r small sources, one must still correct for the telescope beam efficiency, which is commonly referred to as B ef f .T h e n

T MB = F ef f B ef f T * A (2.26)

Spectral line Observations

Observing spectral lines, we can probe the physical, chemical and dynamical conditions of the interstellar matter (ISM) in the Milky Way and in external galaxies.

The intensities of the lines can tell us about gas temperature and density, the chemical composition and therefore about the abundances and the evolutionary state, the ionization and the magnetic fields.

Besides, frequencies and line widths can be used to constrain dynamical models (clouds collapse, protostellar outflows) and to derive redshifts (age and distance).

Spectral lines form throught the interaction between a quantum system (molecules or atoms) and a single photon. Theses systems can change their states only in discrete amounts of energy ∆E. The transition between these states leads to emission or absorption of light at a single frequency ν =∆ E / h ,t h es oc a l l e d rest frequency.

Spectrally, this transition is seen as a line. Molecules, which consist of multiple atoms, can vibrate or rotate around an axis. When a vibrational or rotational state of molecules discretely changes, obeying quantum dynamics, they emit or absorb line radiation. Electron energy transitions as well as vibrational and rotational transitions make molecular lines. These lines are exited, near HII regions radiatively via UV fields, in cold molecular clouds via CMB but predominantly via collisions with H 2 .I nc a s eo f al e v e li n v e r s i o no n eg e t sam a s e r .

Depending on the density and temperature, spectral line emission can be optically thin or thick. In the case of optical depth τ 1, one can look through a cloud and determine column densities and internal dynamics. For τ 1, one can see only the surface of an object. Using radiative transfer codes one can calculate the cloud temperature and density.

Energy uncertainty causes a small broadening, the "natural line width". Thermal motion of emitters leads to Doppler shifted line frequencies. Considering statistical ensembles one can derive a gaussian shape for the broadened line, but only the line of sight, i.e. the radial component adds to this effect. Larger Doppler shifts can occur due to dynamical processes in molecular clouds and stars as ejecta phenomena (outflows).

To observe these spectral lines we need an heterodyne frontend, usually a downconverter from observing frequencies to a low (0-4 GHz) intermediate frequency (IF) band and a spectrometers to analyze the signal.

Spectrometers measure the frontend signal in many frequency bins across the available bandwidth. There are several techniques: (1) filter banks: a series of analog filters; complex electronics or (2) auto-correlators: special purpose computers; correlation function of time series signals; low number of bits.

As we already see the source signal is partially absorbed by the Earth atmosphere and the atmosphere itself radiates and thus contributes to the signal, the telescope beam picks up ground spillover, the receiver etc. add a signal too. To remove the atmospheric and instrumental emissions one observes the target and then a position on sky out-source.The difference of the two measurements contains only the source signal (still weakened by the differential atmospheric absorption).

On-Off measurements can be taken by moving the telescope between two positions. If the source is small, then one can use wobbler switching which is faster.

In IRAM 30-m telescope there are four different procedures available for line observations: position switching, frequency switching, wobbler switching and raster map in position switching.

Position Switching In this observing mode the antenna moves between the source position and a reference position (also called off position) where there is "blank sky" making a given angle. This procedure will provide the simplest means to eliminate the standing waves in the baselines. Note that the half of the observation time is spent off source.

Frequency Switching In this observing mode the local oscillator is shifted by several MHz at a rate of several Hz. The result is two spectra one shifted respect to the other. One substract the two spectra and make an average to improve the signal to noise ratio. This mode eliminates bandpass shape from components Observing below the atmosphere after the mixer. The advantage of this mode is that the integration time is larger than in the other modes because this frequency switching can be done much quicker than moving the telescope, and the antenna is always on-source.

Wobbler Switching In this observing mode the secondary reflector switch between the source and a reference position located at a wobbler throw distance from the source. The subscans are taken alternatively with telescope pointing at minus and plus the wobbler throw distance from the source (maximum 240 arcseconds), so that the reference positions are alternatively at each side of the source.

Raster mapping This procedure is equivalent to the position switching, but spectra for several positions are obtained during a single scan. Raster mapping allows to use the same reference subscan for several positions. The IRAM 30-meter telescope is a single-dish antenna that was built between 1980 and 1984. It is located on Pico Veleta in Sierra Nevada, Spain, it is one of the two radio astronomy facilities operated by IRAM. In table 2.2 are indicated the principal parameters of this telescope. The telescope is an alt-azimuth mounted Cassegrain-Nasmyth system. In order to minimize thermal deformations, the telescope is covered on the entire outside by thermal insulation, only the panels of the panels of the main reflector are exposed to the atmospheric conditions. In addition, an active control system keeps the temperature homogeneity and gradients of the reflector backup structure well within one degree. The main reflector is constructed on the principle of homology [von Hoerner, 1967]. It deforms under the influence of gravity such that at any elevation angle ε the best-fit surface of the deformed reflector is a parabola with deviations within the tolerance specifications of the reflector precision (generally <λ/20 of shortest wavelength of observation). Figure 2.7 shows the gain correspondant to homology deformation for four wavelengths as a function of elevation. The position of the focus and the focal length of the best-fit parabola may change with elevation, however, these effects are to a large extent compensated by corresponding shifts and tilts of the subreflector (built in computer controlled corrections). The panels are optimally adjusted for an elevation ε • =43 • where gravitational deformations are eliminated by a corresponding pre-derfomation of the surface.

Single-dish IRAM-30m telescope

The 30-meter telescope is equipped with a series of single pixel receivers operating at 3, 2, 1 and 0.8 millimeters. The atmospheric conditions for increasing amount of precipitable water vapor (pwv) in air are the worst, especially for higher frequencies. The shorter wavelengths (1 and 0.8 mm) are best suited in winter, when the air is colder and drier.

It is also equipped with two cameras working at 1 millimeter: HERA with 9 pixels for the mapping of molecular gas in extended nebulae.

Receivers and backends

The IRAM 30m telescope is equipped with a series of heterodyne receivers and bolometers operating at four different wavelength bands: 3 mm (100 GHz), 2 mm (150 GHz), 1.3 mm (230 GHz), and since 2010, the 0.9 mm (330 GHz) is available.

The heterodyne receivers that can be used at the 30-m telescope are:

EMIR (Eight MIxer Receiver) It is working since April 2009. It provides a minimum instantaneous bandwidth of 4 GHz in each of the two polarizations. The four EMIR bands cover the four main atmospheric windows in the millimeter range they are denominated as E090, E150, E230, and E330 according to their center frequencies in GHz. The E150 band has SSB mixers with a single sideband available at a time, and the E090, E230, and E330 bands can be operated in 2SB mode where both sidebands are available for connection to backends. In figure 2.8 shows an overview of EMIR receiver and possible combination bands and figure 2.9 shows the atmospheric transmission for 2mm and 4mm precipitable water vapors with the EMIR bands marked together with the frequencies of a few important molecular transitions.

HERA (HEterodyne Receiver Array) consists of two arrays of 3×3p i x e l s with 24" spacing. This receiver has a bandwidth of 1 GHz. HERA covers a band between 215 and 272 GHz (See fig 2.9)T h eb e a mw i d t hv a r i e s between 12" and 9". The main purpose of this first multibeam spectral line receiver on the 30m telescope is to allow, together with the related backends, for rapid raster and on-the-fly imaging of spectral lines in the 1.3mm atmospheric window.

The receivers deliver 18 IF bands, every one being one GHz wide. The autocorrelator VESPA (VErsatile SPectrometer Array), is very flexible and offers a high spectral resolution but narrow band width and thus it is not capable of processes the totality of the bandwidth delivered by the receivers. For this reason a special spectrometer has been designed: WILMA (Wideband Line Multiple Autocorrelator), which has no flexibility at all, but which is capable of crunching all the bandwidth at once (16 GHz ot 18 × 930 MHz), but with a low spectral resolution.

Software

The software used to reduce and/or analyze the data acquired with the IRAM 30M telescope and Plateau de Bure Interferometer PDBI is called GILDAS ( Grenoble Image and Line Data Analysis System) software. It is a collection of softwares oriented toward millimeter and submillimeter radioastronomical applications including both single-dish or interferometric observations. The utilities offered by GILDAS are the following: CLASS (Continuum and Line Analysis Single-dish Software) is a software package for reducing and analyze spectroscopic data obtained on a single-dish telescope.

GREG (GREnoble Graphics) is a graphic utility package whose goal is to help people prepare plots of their data or of the results of theoretical computations.

ASTRO (A Software To pRepare Observations) is a software developed to help astronomers to prepare their observations with the IRAM telescopes, for example, computing the position of sources as seen at a given time from a given location.

PAKO is a user interface used during the observations to command the telescope.

MIRA (Multichannel Imaging and Calibration Software for Receiver Arrays) is used also during the observations are carried out to reduced and view the data immediately after it has been taken.

Observing above the atmosphere

The availability of astronomical satellites starting with IRAS, and then ISO, ODIN, MSX, CHANDRA and SPITZER have increased the number of spatial observations. Somewhat more specialized are the radio telescopes dedicated to the study of the 3 K microwave background. These include the satellite missions COBE and WMAP (See table 2.3). Taken together, these have changed our concepts of astronomy. This trend continue with the launch of the Herschel Satellite Observatory and the start of scientific measurements with the Stratospheric Observatory for Infrared Astronomy, SOFIA. The Herschel Space Observatory observe at sub-mm and far infrared wavelengths with bolometer and heterodyne instruments. The antenna is a 3.5 m paraboloid made of silicon carbide, more than four times larger than any previous infrared space telescope and almost one and a half times larger than that of the Hubble Space Telescope. The structure behind the antenna is a sun shield; the structure below the main reflector is the cryostat containing the receiver systems.

Herschel is equipped with two bolometer cameras, SPIRE (covering 250-520 µm) and PACS (covering 75-170 µm, and a very high-resolution spectrometer, HIFI (covering 157-212 and 240-625 µm). The bolometers also have spectroscopic capability. The detectors in these instruments are cooled to temperatures close to absolute zero by a sophisticated cryogenic system. It is estimated that its lifetime will end in February/March 2013 until the cryostat runs out of helium.

Herschel is being operated as an observatory. About two-thirds of its observing time is available to the worldwide scientific community. The rest of the observing time is being allotted to the instrument consortia. Herschel is operating autonomously, sending acquired data to Earth over a three-hour period every Ta b l e 2 . 3 : List of some satellites that work(ed) from near-infrared to radio wavelengths.

Satellite

Cover Chapter 3

The source: Cep E-mm

The Cep E molecular cloud

The Cepheus E molecular cloud is one of the six molecular clouds inside the Cepheus OB3 association. This association has 40 early-type members at 725 pc [Blaauw et al., 1959] divided in two subgroups [Blaauw, 1964], Cep OB3a and Cep OB3b with ages of 7.5 and 5.5 Myr respectively [Jordi et al., 1996]. The Cepheus OB3 association has an extension of 20 × 60 pc and a average radial velocity of -10 km s -1 [Sargent, 1977] (See Figure 3.1).

The Cepheus E molecular cloud was first identified by [Sargent, 1977]b yC O (1-0) maps over the region of Cepheus OB3. Sargent's CO observations revealed the presence of several molecular clouds in the Cepheus OB3 association. She labeled them as Cep A, B, C, D, E, F (See Figure 3.2).

The Cep E molecular cloud has a size of ∼ 1.5 pc , T ex ∼ 17 K and a density n(H 2 )a b o u t2 . 1 ×10 3 cm -3 estimated from 18 CO(J=1 -0) maps following the assumption of LTE [Yu et al., 1996].

The Cepheus OB3 association was mapped at the positions of the strongest CO emission detected by [Sargent, 1977]inthe1 11 -1 10 transition of formaldehyde at 6cm wavelength, showing that the Cep E molecular cloud is the second most massive (M H 2 of 2100 M ⊙ )a n dd e n s eo ft h eC e pO B 3m o l e c u l a rc l o u d s [START_REF] Few | [END_REF]. Thermal NH 3 and OH (1665/1667 MHz) emission [Wouterlo ot et al., 1988] have been also detected toward Cepheus E molecular cloud. [Dobashi et al., 2005]. The lowest contour of the extinction is at A V =1m a g ,a n dt h ei n c r e m e n ti s 0.7 mag. For details see [START_REF] Kun | Star Forming Regions in Cepheus[END_REF].

Finally, during a survey carried out by [Wouterloot and Walmsley, 1986]u sing 100-m Effelsberg telescope to detect H 2 Om a s e r -e m i s s i o nf r o mt h eC e p h e u s molecular clouds, they found eleven new sources, among them Cep E-mm source, or IRAS 23011+6126. Several authors have studied this source, after the first identification by [Lefloch et al., 1996]a sab o n afi d eC l a s s0 ,a c c o r d i n gt ot h ec r i t e r i ae s t a b l i s h e d by André, Ward-Thompson, & Barsony 1993 (See section 1.2.4). [Crimier et al., 2010], [Moro-Martín et al., 2001] studied the properties of the molecular gas and dust envelope. [Fro ebrich et al., 2003b]s t u d i e dt h en a t u r eo ft h ec e n t r a lp r o t ostellar source, which was questioned at about the same time by [Noriega-Crespo et al., 2004].

The Cepheus E-mm source

All these studies allow to constrain the source Cep E parameters as reported in Table 3

The dusty envelope

The protostellar envelope emits in the millimeter continuum, due to the dust emission of all the material surrounding the protostar.

There was some debate on Cep E belonging to Class 0 or Class I classes. Based on millimetric observations [Moro-Martín et al., 2001]a ffi r m e dt h a tt h eS E Do f profiles for a IM-protostars sample, finding that Cep E shows a density profile consistent with the inside-out collapse theory [START_REF] Shu | [END_REF], where the envelope density follows a single index power law with α=1.9 at wavelengths larger than 60 µm. This is;

n(r)=n 0 × r 0 r α (3.1)
where the power law index, α,t h ed e n s i t y ,n 0 ,a tr 0 are free parameters of their model. The envelope is assumed to start at a radius r in and extends to r out .

In the Table 3.3 are listed the input parameters as well as the best fit parameters and some relevant physical quantities obtained from the radiative transfer analysis of Cep E protostar. Crimier et al. found that the maximum difference between T dust and T gas is 5K so gas and dust are coupled. This model yields a mass envelope of 35 M ⊙ and a best fit value for the luminosity of 100 L ⊙ .

Others envelope mass values have been obtained by others authors assuming power law-density ∼13 M ⊙ and bolometric luminosity of ∼30 L ⊙ [Moro-Martín et al., 2001]; [Noriega-Crespo et al., 2004]ac o n s t a n td e n s i t yo v e ras p h e r i c a lv o lume; 10 M ⊙ [START_REF] Ladd | Herbig-Haro Flows and the Birth of Stars,V o l .1 8 2o fIAU Symposium[END_REF]a n di nd u s tt e m pe r a t u r ed e pe n d e n tm od e l s ; 18 M ⊙ [Lefloch et al., 1996]. SED models based on a grey body approximation [Fro ebrich et al., 2003b]p r e d i c t∼78 L ⊙ and a envelope mass of 7 M ⊙ taking a β=1.0±0.3, where β is the submm slope of the SED. The summary of these studies are listed in Table 3 [Lefloch et al., 1996] The Cepheus E-mm source Ta b l e 3 . 3 : Summary of the dust radiative transfer analysis for Cep E protostar. Adapted from [Crimier et al., 2010].

Fixed input parameters Stellar temperature, T 5000 K Dust temperature at r in , T in 300 K Dust opacity (OH5) at 100 µm, κ 100 86.5 cm 35 800 Radius at T dust =100K,r 100 K , (AU) 223 H 2 density at r 100 K , n 0 ,( c m -3 )2 .0 × 10 8 Envelope mass, M env ,( M ⊙ )3 5 (T dust -T gas ) max ,( K ) 5

Single or binary protostar

Cep E is proposed to be a double source separated by 1.4" (∼1000 AU) based on continuum images at 22 GHz [Moro-Martín et al., 2001]. (See Fig. 3.5 at the right).

Moreover, line and continuum observations (See section 3.3), showed two separate outflows structures emanating from the mm central source suggesting the presence of a second driving source (e.g. [Eisloffel et al., 1996], [Ladd and Hodapp, 1997]) See Figures 3.5 and 3.10. However, [Velusamy et al., 2011]d on o t find any evidence of a double source in the IR bands using the Spitzer imaging with a subarc resolution (∼ 0.6" -0.8") enough to resolve any double source separated by 1.4". These observations imply that Cep E-mm is a binary with a separation smaller than ∼1". This issue is still under discussion. depletion factor and the gas temperature [Roberts et al., 2003]. However, it is not possible to establish an evolutionary trend among Class 0 sources based on simple parameters such as the average CO depletion and the average N 2 H + deuterium fractionation. This is due to the complexity of these regions and because the observations do not reach the sufficient angular resolution to trace the inner regions of the envelopes. Alonso-Albi et al. derived also the radial abundance profiles across the sample envelopes of C 18 O, N 2 H + and N 2 D + .T h e i rr e s u l t so nC e pEa r es h o w e di nt h e Figure 3.6.

Using a spherical-envelope steady-state chemical models [Caselli et al., 2002] and assuming a CO binding energy of ∼1100 K [Alonso-Albi et al., 2010]f a i l e d to reproduce their observations towards Cep E . The qualitative behavior is reproduced for the C 18 Oa n dN 2 H + emission, but the model reproduces the line integrated intensities poorly. The line intensity predictions are a factor of 2-4 higher than the observations. Increasing the CO binding energy and thus reducing the efficiency of cosmic-ray-induced desorption does not improve the fit (see model 2 in 3.6). In addition, this also causes an increase in deuterium fractionation, which is not observed.

The authors propose that this discrepancy is due to the assumption of spherical geometry in the model and which do not take into account important phenomena like bipolar outflows and UV radiation from the star.

Cepheus E bipolar outflow

Cep E drives a bipolar outflow. The first detection was obtained by Fukui in 1989, as part of an IRAS data base survey. Years later, in 1994 the first image of Cep E was taken by Hodapp during the near-infrared K'-band imaging survey of all regions associated with CO molecular outflows from the list of Fukui 1989 (Fig. 3.7). Hodapp described the bright and compact object as a bipolar nebula, however it was not possible to know if it was a reflexion nebula or a line emission region.

The next year, Eislöffel and colleagues, carried out a near-infrared imaging survey for molecular hydrogen emission in CO outflows. They found that Cep E 1) line contours of Cep E that can explain the wiggle structure in the southern lobe and the regular sideways positional offsets of the bows of the northern lobe. Taken from [Eisloffel et al., 1996].

lines. [Lefloch et al., 1996] estimated an age for the outflow of about 3 ×10 3 yr. The 12 CO(2-1) line presents high velocity wings extended up to -125 km s -1 relative to the ambient gas (v lsr ≈-11 km s -1 )i nt h eb l u e -s h i f t e dl o b ea n dt o 80 km s -1 in the red-shifted lobe (Fig. 3.9).

In addition, the Cep E outflow shows the presence of bullets of high-velocity material, in both lobes, situated symmetrically at almost the same distance from the central source, with average velocities of 100 km s -1 and 70 km s -1 relative to the ambient gas.

The outflow lobes are relatively compact, with a typical size of ≈ 30" × 10" at half power. The estimated kinematical ages for the blue and red lobe are 1.1×10 3 and 1.8×10 3 yr respectively [Lefloch et al., 1996].

The SiO emission traces the high-density material of the bullets detected by 12 CO lines, and not the ambient gas. Comparison with the H 2 observations [Eisloffel et al., 1996], and the continuum 1.25 mm dust emission [Lefloch et al., 1996]s u g g e s t st h a to u t fl o wi n t e r a c t ss t r o n g l yw i t ht h ec i r c u m s t e l l a rm a t e r i a l , sweeping up the gas and dust of the parental envelope.

In addition, water emission has been detected in the high-velocity gas close to [Lefloch et al., 1996].

the protostar by [Lefloch et al., 2011]. They observed a masing line at 183 GHz, in order to elucidate the origin of water emission in protostellar systems. This water emission can be observed only at the IRAM 30m telescope, under excellent and stable weather conditions in winter. Although this emission is unresolved in the main-beam of the IRAM30m telescope, the authors used PdB observations at 2" of resolution to help establish the origin of the emission detected, and the physical conditions in the emitting gas. Assuming an ortho-to-para ratio of 3, it is possible to derive a water column density of N(H 2 O)≈ (1.2-8.0) ×10 17 cm -2 .

The jet gas kinetic temperature and column density values were determined by [Lefloch et al., 2011]obtainingaT K ≈ 120 K and a N(CO)≈ 2.5 ×10 17 cm -2 , adopting the standard value for the CO abundance in the gas, we have N(H 2 )≈ 2.5 ×10 21 cm -2 . Based on the ratio SiO J=8→7 to J=3→2 the authors estimate HH 377 S2 3 0 1 1 + 6 1 2 6r e g i o n . T h ed u s t emission trace the quadrupolar outflow observed by [Ladd and Hodapp, 1997]( S e et e x t ) . T a k e nf r o m [ Chini et al., 2001]. [Eisloffel et al., 1996]. The southern lobe is coincident with HH 377 (See section 3.4). The low-level emission seen at 850 and 1300 µm to the south-east is coincident with the bipolar outflow estructure detected in 12 CO J= 2-1 by [Ladd and Hodapp, 1997].

3.11). The elongation of the central source at the higher level contours at 450 µmi nt h en o r t h -s o u t hd i r e c t i o ni sc o i n c i d e n tw i t ht h eb i p o l a ro u t fl o wd e t e c t e d in molecular hydrogen by

In summary, the observations indicate the presence of two bipolar outflow structures, a north-south bipolar outflow traced by H 2 1-0S ( 1 ) ,as o u t h -e a s t bipolar outflow traced by 12 CO J= 2-1. A group of knots suggest the presence of a single lobe traced also by H 2 1 -0 S(1) eastward. However, it is not clear so far if this is a binary or multiple source or a time and angle variable driving source that generate these outflow structures. This issue is still open.

HH 377

[Noriega-Crespo and Garnavich, 1997]d e t e c t e dt h eo p t i c a lc o u n t e r p a r to ft h e southern bowshock observed at 2µmb y [ Eisloffel et al., 1996]. This optical knot was observed in Hα and [SII] λλ6717,6731 emission and has been named HH 377 by [START_REF] Devine | Herbig-Haro Flows and the Birth of Stars,V o l .1 8 2o fIAU Symposium[END_REF].

The structure of the H 2 outflow and the optical emission of only one compact knot is illustrated in Figure 3.12.I nt h i sfi g u r et h e r ea r eH 2 1-0 S(1) contour maps of Cepheus E, overlaid on gray-scale and contour representations of optical images at left with Hα and at the right with [SII]. These maps clearly show the coincidence between the optical knot and the brightest region of the southern lobe of the IR flow. Cep E Herbig Haro object is one cases like HH 1-2 [Davis et al., 1994] where the J-type (inferred from the optical spectra) and C-type (inferred from the IR spectra) seems to coexist spatially [START_REF] Ayala | [END_REF]].

HH377 has characteristics typical of low-excitation (∼ 1% ionization) Herbig-Haro objects and an anomalous high electron density, n e ∼ 4100 cm -3 ,suggesting an unusually high preshock density of ∼ 10 5 cm -3 [START_REF] Ayala | [END_REF]. The proper motion of this optical knot is 107±14 km s -1 [Noriega-Crespo and Garnavich, 2001]w i t har a d i a lv e l oc i t yo f-7 0±10 km s -1 [START_REF] Ayala | [END_REF].

Hydrodynamic models of molecular jets have been produced to simulate observed molecular outflows from protostellar sources e.g., [Suttner et al., 1997]. In their work they carried out 1D, 2D and 3D simulations of molecular jets to model observed complex structure of the jets. The simulates emission maps are in good correspondence to a number of observations of deeply-embedded outflows. They compare the 3D simulate images of the 1-0S(1) H 2 emission to a sample of observed outflows; Cep E, HH21, HH212, HH288 and L1634. They remarked that the H 2 lines peak at the bow shock, 1-0S( 1) being somewhat more extended the 2-1S (1). At the tip the emission originates mainly from shocked jet gas, in the wings ambient H 2 survives and radiates. A bow-like arc upstream of the main bowshock is formed and similar features had been observed in protostellar jets e.g. HH288, L1448, DR 21 [Davis and Smith, 1996]. Rotating the simulated images this bowshock looks like and elongated knot similar to HH212 [START_REF] Zinnecker | Disks and Outflows Around Young Stars,V o l .4 6 5o fLecture Notes in Physics[END_REF]HH211 [McCaughrean et al., 1994]. So authors achieve to reproduce features which seem to be multiple bow shocks are generated naturally by continuous jet. [Suttner et al., 1997]f o u n dt h a tC e pEj e tr e q u i r eaj e td e n s i t yo f1 0 5 cm -3 , they remark that even if the density in optical jets is hardly above 10 3 cm -3 this value is due to: (1) the hydrogen molecule is observed to be close to LTE, (2) the overall energetics require high densities (3) dense clumps and cores have been deduced from observations, (4) jets would take excessively long time to escape from cores if their density were relatively low. This value is consistent with the [START_REF] Ayala | [END_REF]a n dl a t e rw i t ht h e [ Smith et al., 2003]r e s u l t s . T h i sc o u l db e as i g no fav e r ye a r l ys t a g eo fj e td e v e l o p m e n ta n ds u p p o r tt h es c e n ei nw h i c h Cep E is a very young outflow breaking through its parental molecular core.

In Table 3.4 we present a summary of Cep E outflow properties, how they were deduced and their references.

Shock models of Cep E

Many authors have tried to interpret the wealth of observations toward Cep E making use of shock models (e.g., [Eisloffel et al., 1996], [START_REF] Ayala | [END_REF], [Smith et al., 2003], [Noriega-Crespo et al., 2004]). The comparison between the measured fluxes and the predictions from two different shock models, C-type and J-type, can help to constrain the properties of the outflows (shock velocity, brich et al., 2003a]. However in the case of Cep E, simple planar J-shock or planar C-shock models do not reproduce the observations. Single J-type shock waves predict high excitation H 2 spectra and C-type shock waves predict constant excitation across a wide range of upper energy levels [Smith et al., 2003].

Better results come from multi-shock fits, as shown in left panels in the Fig. 3.13 for the southern lobe. [Smith et al., 2003]f o u n dt w os t e a d y -s t a t eC -t y p e planar shocks that fit the full set of data. However it is not clear how two shocks Steady-state C-type shocks with a paraboidal bow shock geometry seem account for the molecular excitations of H 2 10 S(1) rovibrational and CO rotational transitions from the with a bow speed of >38km s -1 [Smith et al., 2003].

Steady-state C-type shocks are strongly favored in the case of Cep E (e.g., [Eisloffel et al., 1996], [Ladd and Hodapp, 1997], [START_REF] Ayala | [END_REF], [Smith et al., 2003]).

AfurtherdiagnosticisprovidedbytheintegratedintensityoftheCOemission. The C-bow shock model predicts a CO(3-2)/CO(18-17) ratio of 6, close to the measured ratio of 4 for the north lobe and 2 for the southern lobe [Smith et al., 2003].

High-resolution H 2 1-0 S(1) imaging show that lobes consist of numerous emission knots of size (1-3)×10 15 cm [Ladd and Hodapp, 1997]. To reproduce the integrate emission H 2 1-0 S(1) of each lobe the model required around 20 bow shocks, close to the paraboidal shape [Smith et al., 2003].

Chapter 4

The spectral survey of Cep E

Introduction

Spectral surveys are a powerful tool to obtain a complete census of the molecular complexity in star forming regions, and to understand the interplay between the molecular composition and the star forming process.

The dynamic of the collapse is regulated by the molecular composition, as well as the gas temperature and the coupling between matter and the magnetic fields. Conversely, the gravitational collapse influences the gas physical conditions, as the infalling material releases its energy into photons that heat the collapsing envelope, and, consequently change the chemical composition and gas cooling.

In fact, the thermal balance of the gas depends on the molecular composition, through the gas cooling function. Conversely, the temperature and density controls the molecular composition through the chemical network which depend on the gas temperature, density and their abundance.

In addition, the ambipolar diffusion (the coupling of matter to the magnetic field) is controlled by the ionization fraction.

Besides, the emission lines can be a powerful diagnostic tool, because different lines from the same molecule probe different excitation conditions and permit to trace regions of different physical conditions, from multifrequency analysis and modeling. Likewise, lines are the only way to study the dynamics.

The few unbiased spectral surveys carried out in the millimeter and submil-Introduction limeter bands accessible from the ground were obtained mostly toward high-mass stars (e.g. Orion KL [Tercero et al., 2010], [Tercero et al., 2011]and [Tercero et al., 2012], AFGL 2591, IRAS 20126+4104 [Plume et al., 2007], IRAS 23385+6053 [Thompson and Macdonald, 2003], NGC 6334 I, I(N)) [Thorwirth et al., 2007]and compact HII regions as W3 IRS5, IRS4 [Kim et al., 2006], G34.3+0.15 [Kim et al., 2000], G5.89-0.39 [Thompson and MacDonald, 1999]. This is mainly because of their strong molecular lines and their rich chemistry [Herbst and van Dishoeck, 2009]. Only two solar-type sources have been studied in the mm/submillimeter window: IRAS 16293-2422 [Caux et al., 2011], L1527 [Sakai et al., 2008].

Finally, the launch of the Herschel Space Observatory has permitted the spectral surveys in the 500 -2000 GHz in sample of low to high mass star forming regions [Ceccarelli and CHESS Consortium, 2010]. It is important to note that the intermediate-mass regime has not been covered yet in the mm/submm range. Hence their chemistry, the influence of the central object and its radiation field are not well known. Before this thesis, there was no unbiased line surveys of a molecular outflow. At the same time when Cep E was been observed, spectral line surveys were carried out towards L1157-B1 (IRAM and Herschel); and another one at 3 mm with Nobeyama Radio Observatory (NRO) 45m telescope [Sugimura et al., 2011].

Line surveys towards outflows and jets give access to unique information to the feedback of outflows on the surrounding molecular gas, in particular the impact on the gas chemical composition in the region.

In the last years, observational studies were carried out towards young protostellar outflows, though mostly fragmentary studies e.g. [Bachiller, 1996], [Bachiller et al., 2001], [Jiménez-Serra et al., 2004]. The most studied source, until now, has been L1157-B1, whose investigation of the chemical composition was restricted to bright molecular line emission. The exceptional line richness of L1157-B1, associated with a peculiar gas composition, distinct from that of the protostellar envelope, led R. Bachiller to propose the existence of a class of young protostellar outflows, "chemically active" (See section 1.4.4). These biased surveys together with the new generation of models with both, collimated and wide-angle wind components (e.g. [START_REF] Ferreira | Herbig-Haro Flows and the Birth of Stars,V o l .1 8 2o fIAU Symposium[END_REF]][ Shang et al., 2006], [Machida et al., 2008]), are an attempt to achieve a consistent picture of outflows morphology, kinematics and chemical abundance enhancements.

The motivation to carry out an unbiased survey of CepE with the IRAM 30m telescope are: 1) To constrain the actual comp osition of a typical intermediatemass star forming region (envelope and molecular outflow); 2) to understand how the gas chemical composition is altered by the outflow process; 3) comparison with shock models for the outflow emission. ), emission is associated with the large scale, extended outflow entrained material; at moderate (between 30 and 45 km s -1 )v elocit yemissionisassociated with the cavity at the base of the north jet; the high (≥ 45 km s -1 )v e l o c i t yi s associated with the north jet. Towards the south outflow lobe, (between -65 and -80 km s -1 )i tc a nb ed i s t i n g u i s hac o n c e n t r a t i o no fC O( 2 -1 )a tt h ea p e xo ft h e south jet and at extremely high velocity (≤-110 km s -1 )i sd e t e c t e de m i s s i o n associated with the south jet.

Kinematics and morphology of the outflow

Observations

The unbiased line spectral survey towards Cep E was carried out with the IRAM 30-m radiotelescope at Pico Veleta, Spain, in three different runs: March and July 2010 and February 2011.

Two positions were observed: toward the driving protostellar at source offset position (0, 0), at 23h03m13.0s +61 • 42'21" (J2000) and towards the apex of the south outflow lobe at offset position (-12", -18"). These two positions were chosen in order to a) obtain a census of the molecular composition of the protostellar envelope, including a possible hot corino, b) obtain a census of the molecular composition of a typical bowshock region (HH377), at (-12,-18), c) distinguish between protostellar and shock origin for the molecular species detected.

The observed positions are marked by yellow crosses in the CO (2-1) emission map of Figure 4 weather conditions, with a precipitable water vapor (pwv)o f∼1m m ,as y s t e m temperature in the range of 250 to 600 K and an opacity between 0.1 to 0.4 during the observations. Overall, the observations required a total of about 80 hours of observation time.

The beam of the telescope varies from 7 arcsec for the highest frequencies (350 GHz) to 30 arcsec for the lowest frequencies (80 GHz) (see Fig. 4.3). The achieved rms after averaging all the observations for each band is reported in Table 4.1.

Data calibration 4.3.1.1 Atmospheric calibration

In the millimeter and submillimeter wavelength range, the atmosphere has a large influence and can change rapidly, affecting the signal received by the telescope. Such perturbations will be recorded as fluctuations over the continuum along all the spectra, as it is showed in Fig. 4.4.

In the survey the atmospheric calibration was obtained performing three measurements in the "cal" procedure, including a measure on the sky. The sky opacity and calibration was obtained using the program ATM (Atmospheric Transmission at Microwaves) installed at the IRAM 30m telescope (Cernicharo 1985). The observing mode Wobbler switch was chosen precisely to minimize the drifts in atmospheric calibration by subtracting the emission of the reference position every sec.

Thus, the first step is to remove the effect of the changing baselines, caused by variations in the receiver gain with antenna temperature. Usually, only linear baselines were removed from the spectra. In some spectra, we found it difficult to fit the raw spectra with only the first order of polynomial function, and in these cases we extract the baselines with a higher polynomial order or in some cases we also used sinusoidal functions. Then we can start to work with these clean spectra. Later, in order to convert Ta* (flux units of the spectra) into an intensity as close as possible to the true brightness temperature of a source that has a size of the same order as the main-beam, we correct for the telescope beam efficiency.

For the IRAM 30 m telescope, the forward efficiency (this is basically the coupling of the telescope with the sky. See section ??)i sF ef f ∼ 0.9 down to 3 mm wavelength (See table 4.1), as provided by IRAM's website.O nt h eo t h e r hand, B ef f varies with the wavelength according to the Ruze formula1 :

η a = η ill × exp((4πσ s /λ) 2 ) (4.1)
For our data analysis we made a (simplified) linear interp olation b etween values determined by the IRAM staff at given frequencies. These interpolation fuctions are listed in table 4.2.

Spurious features

Spikes There is another problem that can appear in the spectra, the presence of spurious lines in form of spikes. These "spikes" come from interferences in the electronic chain of nearby receivers; they are tracked and catalogued by the IRAM's staff. However, there are other spikes that can appear and these can be recognized because they appear as an extremely bright emission or absorption features in just one or two channels. These spikes were flagged from the spectra. Contamination from image band The rejection of the image band is limited to 10 -15 dB. This is not sufficient anymore at the level of sensitivity reached by our observations. We indeed observed "ghost" lines from bright lines in the image band (See figure 4.6). This ghost lines were flagged in the individual spectra.

Errors

The detection limit for emission lines in the survey was established at 3 rms noise level, in order to assure a reliable line identification.

In practice, the integrated flux over a velocity range is limited at 3 σ,w h e r e the σ and the rms are related as follows:

T mb dv ≤ 3σ =3 * rms √ ∆vδv (4.2)
where, rms is the noise over one channel, δvt h es pe c t r a lr e s o l u t i o n( 2 M H z ) ,a n d ∆v the width of the velocity range over which we are integrating.

This expression was calculated as the quadratic sum of the errors of each channel along the velocity interval in which one makes the integral (Fig. 4.7). Then, the error of the integrated flux over a velocity interval is:

σ 2 Int.flux = σ 2 channel 1 δv 2 + σ 2 channel 2 δv 2 + σ 2 channel 3 δv 2 + ... σ 2 Int.flux =Σ N i=1 σ 2 channel δv 2 (4.3)
but the σ channel is the rms and if we assume that it is the same over each tunning, .9: Three kinematical components were found: the ambient gas at ∼ -10.9 km s -1 ,t h ee n t r a i n e d gas between -90 → -12 and -9 → +30 and the jet, detected at -125 km s -1 (south jet) and at around +50 km s -1 (north jet).

Most of the detected lines have upper level energies lower than 50K, normally are associated with cold gas. Conversely, only 14% of the lines have upper level energies larger than 100K, suggesting that Cep E spectrum is unlikely dominated by hot gas emission. Some line profiles, notably those of CO and HCO + ,s h o wt h ee x i s t e n c eo f three velocity components (Fig. 4.9), likely associated with three different physical components: the ambient gas at ∼ -10.9 km s -1 , the entrained gas at (-90 → -12) and (-9 → +30) and the jet, mostly detected at -125 km s -1 (south jet) in the (-12,-18) position and at +50 km s -1 (north jet) in the (0,0) position. All the components are detected in both positions, but not in all molecular tracers (Table 5.1).

In Figure 4.9 we can see that some line spectra like CO show extremely high velocity (EHV) features, that was already seen in [Lefloch et al., 1996]. This means that molecular spectra present a secondary component less intense than the main emission peak from the ambient gas but enough intense and placed at high velocities to be identified as emission from a driving jet. e.g., [Bachiller et al., 2001], [Jørgensen et al., 2007], [Santiago-García et al., 2009], [Hirano et al., 2010].

In an preliminary analysis, we detect a total of 559 lines whose 428 identified. Tables 4.11 and 4.12 list all the detected lines and the integrated emission be- 4.26 present the spectra with the labels of each line identification listed in Tables 4.11 and 4.12 for (0,0) and (-12,-18) offset respectively.

In chapter 5 we are going to present the results of the detailed study of the jet component.

Line identification

The systemic velocity of Cep E is -10.9 km s -1 [Sargent, 1977]. Line identification was carried out using the package of CLASS called WEEDS [Maret et al., 2011]. Two databases available online were used in order to identify lines: (1) Cologne Database for Molecular Spectroscopy 1 (CDMS; see Mller et al. 2001Mller et al. , 2005) ) and ( 2) Molecular Spectroscopy database of Jet Propulsion Laboratory2 (JPL; see Pickett et al. 1998).

A close match in frequency between an observed and cataloged line is a necessary, but not sufficient, condition for identification. One must take into account all the transitions available in the observed spectral band(s), the spontaneous emission Einstein coefficient and verify the consistency. In total we have 132 unidentified lines out of 559 lines, namely 23% in our preliminary analysis. Table 4.5 lists the species detected in the survey. Furthermore, the chemical richness is more important towards the central source than towards the south lobe.

Rotational diagrams

The abundance and excitation conditions of the molecular species detected were derived in first approximation under the hypothesis of LTE and optically thin emission, carrying out a simple rotational diagram analysis. (See Appendix A).

The method consists in plotting the column density per statistical weight of a number of molecular energy levels as a function of their energy above the ground state, following the standard formula e.g. [Goldsmith and Langer, 1999]:

ln N u g u = - 1 T rot E u k + ln N -ln Z, (4.10)
where N u is the column density of the upper energy level with statistical weight g u ,Ni st h et o t a lc o l u m nd e n s i t y ,T rot is the rotational temperature (∼kinetic temperature, if LTE applies) and Z is the partition function. From a straight line least-squares fit to the observed values of ln(Nu/gu) versus E u ,t h ey -i n t e r c e p t gives the column density, and the negative reciprocal of the slope gives the rotational temperature.

As a result of the LTE analysis we derived the T rot and column density values for the different species. The line emission was integrated between -20 and 10 km s -1 for each line. The quoted values refer therefore to the quiescent gas, namely the envelope toward the central position and the cloud toward the south lobe position. Note that we did not include a beam dilution correction as the source size is large (about 50 arcsec, ∼ 35 000 AU [Crimier et al., 2010]at730pc of distance) than the largest telescope beam. The values of T rot and the column densities are listed in Table 4.4.

Overview of the chemical structure of the envelope

The detection of lines from complex molecules like the acetaldehyde (CH 3 CHO), the methylformate (CH 3 OCHO) and probably the dimethyl ether (CH 3 OCH 3 ) reveals that Cep E has a "hot core" like region. See Table 4.8.

These chemical characteristics are thought to arise from thermal evaporation of ice mantles from the grains close to the protostar, possibly followed by rapid high-temperature gas-phase reactions. In this scenario, the fully hydrogenated species are first generation molecules produced by surface chemistry on the grains, whereas complex organic molecules like HC 3 N, CH 3 OCH 3 and HCOOCH 3 are second generation produced by gas-phase chemistry between evaporated species e.g., [Ceccarelli, 2007]; [Herbst and van Dishoeck, 2009].

We detected emission from ions at both positions, towards the source (HCO + , DCO + ,H 13 CO + , HC

Molecular deuteration

Cep E presents a strong deuteration with 11 different detected species in all the survey (See Table 4.7).

Toward the central p osition, we detected CH 2 DOH, DCO + ,D 2 CO, HDCO, HDO, DNC, CCD, all species already detected toward IRAS 16293-2422 protostar [Caux et al., 2011]justtomentionthalowmassprotostarwhereasimilarspectral survey has been carried out in the past. In addition, we detected DCN, NH 2 D, N 2 D + and HDCCO. The brightest deuterated species line is from DCO + ,w i t h intensity ∼ 0.33 K.

Toward the position (-12,-18), we detected in addition species like CH 2 DOH, NH 2 D, HDCO already detected toward LL1157-B1 (Codella et al. 2012), DCO + and HDCCO. The brightest line are from DCO + , HDCO and HDCCO with an intensity about 0.030 K. Figure 4.17 shows some of the line profiles of the different deuterated species observed towards Cep E. They are, therefore, emitted by the cloud hosting Cep E.

To derive the rotational temperatures and the column densities of these species we carried out a rotational diagram analysis for the deuterated and non deuterated molecules in order to compare their abundances. The results are shown in Table 4.9 and in Table 4.10 are listed the abundance ratios.

The CH 2 DOH/CH 3 OH ratio is found to be 0.6 +0.5 -0.3 for the ambient gas at central position. This ratio is similar to what found in IRAS 16293-2422 with a value of 0.37 +0.38 -0.19 [Parise et al., 2006]andhigherthantheCH 2 DOH/CH 3 OH ratio found towards L1157-B1 (0.2-2) ×10 -2 [Sugimura et al., 2011]; [Codella et al., 2012]. The DNC/HNC ratio in Cep E is 0.1 larger than in IRAS 16293-2422 (0.03±0.015 [?]). In the case of N 2 D + /N 2 H + ratio, the value found for Cep E is 0.08 in agreement with the value of 0.03 derived by [Alonso-Albi et al., 2010]. It is also similar to the values found in other Class 0 sources, ∼0.1 for L1448mms and ∼0.2 for IRAS 03282 [Emprechtinger et al., 2009]. For HDCO/H 2 CO, we find a ratio of 0.1

Species

Position (0,0) Position (-12,-18) No. Ref. 39,44,192,204,221,222,281,339,341,347,366,374,375,378,385,486 CH 3 CHO 14 7 64,65,72,73,75,84,85,126,127,135,136,159,415 8,13,181,188, 189,190,233,295,307, 318,319,335, 22,86,111,178,203 322,349,361,445,455 466,522,528,542,552 34 SO 3 1 32, 168, 179 SO 2 22 7 5,100,101,143,144,163,216,256,260,287,297,316,317,331,338,382,419,422,425,450,488,510,544 OCS 10 1 10,79,112,197,244 291,336,372,408, ,98,99,102,162,164,264 276,277,278,279,280,294 379,390,406,469,520, ,55,92,110,165,193,234,263 300,334,365,386,421,426 ,150,215,249,253,259,299, 395,430,446,476,505,532 H 2 CO 10 9 169,191,195,196,225,418,427,428,232, Ta b l e 4 . 7 : Summary of deuterated molecules detected. Between brackets the total number of transitions in the case of multiplets are indicated. In the third column are the reference numbers from Tables 4.11 and 4.12.

Species

Position (0,0) Position (-12,-8) No. Ref.

CH 2 DOH 16 1 39,44,192,204,221,222,281,339, 341,347,366,374,375,378,385, ,150,215,249,253,259,299, 395,430,446,476,505,532 Ta b l e 4 . 8 : Summary of complex organic molecules detected. Between brackets the total number of transitions in the case of multiplets are indicated. In the third column are the reference numbers from Tables 4.11 and 4.12.

Species

Position (0,0) Position (- 39,44,192,204,221,222,281,339,341,347,366,374,375,378,385,486 CH 3 CHO 14 7 64,65,72,73,75,84,85,126,127,135,136,159,415 191,195,196,225,418,427,428,232 34.1 +20.5 -12.9

Ta b l e 4 . 1 0 : Ratios between some of the deuterated molecules and its corresponding non-deuterated molecule.

In the first column there are the values obtained in our survey towards (0,0) position. The values listed here correspond to the ambient gas component (envelope) towards offset (0, 0). See the text. The other columns show the same ratios towards Class 0 low-mass protostars, pre-stellar cores, L1157-B1 and IM protostars in order to compare our results. 4 .11). Features appearing as absorption spikes were removed because they are artifacts. Signals are in antenna temperature (K). 2 0 0 -1 1 1 0 8 2 0 9 3 , 5 5 96 , 4 0 , 0 1 4 0 , 0 0 3 ≤ 0,060 0,153 ± 0,025 ≤ 0,055 5S O 2 8 1 7 -8 0 8 8 3 6 8 8 , 0 9 33 6 , 7 0 , 0 1 3 0 , 0 0 3 ≤ 0,060 0,268 ± 0,024 ≤ 0,055 6U L 8 3 7 5 9 , 0 2 9 0 , 0 1 0 0 , 0 0 3 ≤ 0,060 0,192 ± 0,024 ≤ 0,055 7C H 3 O H 5 -1 -4 0 0 8 4 5 2 1 , 1 7 24 0 , 4 0 , 1 6 3 0 , 0 0 1 0 , 1 9 1 ± 0,015 2,419 ± 0,012 0,038 ± 0,014 8c -C 3 H 2 3 2 -3 1 3 8 4 7 2 7 , 6 8 81 6 , 1 0 , 0 0 5 0 , 0 0 1 0 , 1

Line catalogue

- -- -- -- UL 247650,048 - - -- -- -- c-C3H2
5 2 3 0-4 3 2 0 249054,368 41 0,026 0,005 ≤ 0,058 0,199 ± 0,024 ≤ 0,053 NO 3 1 3 4 -2-1 2 3 250436,848 19,2 0,025 0,004 0,137 ± 0,024 0,351 ± 0,020 ≤ 0,044 NO 3 1 3 3 -2-1 2 2 250440,659 19,2 0,025 0,004 0,137 ± 0,024 0,351 ± 0,020 ≤ 0,044 NO 3 1 3 2-2-1 2 1 250448,530 19,2 0,014 0,004 0,412 ± 0,024 0,372 ± 0,020 ≤ 0,044 CH3OH 11 0 11 +0-10 1 10 +0 250506,980 153,1 0,049 0,004 ≤ 0,048 0,399 ± 0,020 0,358 ± 0,022 CH3COCH3 1410 5 0-13 9 4 0 250634,375 87,7 0,016 0,004 ≤ 0,048 0,069 ± 0,020 ≤ 0,044 NO 3-1 3 4-2 1 2 3 250796,176 19,3 0,020 0,004 0,141 ± 0,024 0,177 ± 0,020 0,075 ± 0,022 NO 3-1 3 3-2 1 2 2 250815,594 19,3 0,030 0,004 ≤ 0,048 0,146 ± 0,020 0,170 ± 0,022 NO 3-1 3 2-2 1 2 1 250816,954 19,3 0,030 0,004 ≤ 0,048 0,146 ± 0,020 0,170 ± 0,022 SO2 13 113-12 012 251199,675 82,2 0,048 0,004 0,301 ± 0,024 0,810 ± 0,020 0,356 ± 0,022 SO2 8 3 5 -8 2 6 251210,586 55,2 0,023 0,004 0,266 ± 0,024 0,697 ± 0,020 ≤ 0,044 c-C3H2 7 0 7 0-6 1 6 0 251314,337 50,7 0,064 0,004 ≤ 0,048 0,261 ± 0,020 0,167 ± 0,022 c-C3H2 6 2 5 0-5 1 4 0 251527,302 47,5 0,030 0,004 ≤ 0,048 0,305 ± 0,020 ≤ 0,044 CH3OH 6 3 3 -0-6 2 4 +0 251738,520 98,5 0,040 0,004 ≤ 0,048 0,175 ± 0,020 ≤ 0,044 CH3OH* 5 3 -0-5 2 + 0 251811,956 84,6 0,075 0,004 1,634 ± 0,024 3,020 ± 0,020 0,924 ± 0,022 SO 6 5-5 4 251825,770 50,7 0,207 0,004 1,476 ± 0,024 3,105 ± 0,020 1,491 ± 0,022 CH3OH 8 3 6 +0 -8 2 7 -0 251984,702 133,5 0,035 0,004 ≤ 0,048 0,171 ± 0,020 0,532 ± 0,022 CH3OH 9 3 7 +0-9 2 8 -0 252090,380 154,4 0,030 0,004 0,074 ± 0,024 0,214 ± 0,020 ≤ 0,044 CH3OH 10 3 8 +0-10 2 9 -0 252252,850 177,6 0,027 0,004 0,072 ± 0,024 0,111 ± 0,020 ≤ 0,044 CH3OH -30 TadV 1C C -1 3 S 6 7 -5 6 7 8 1 1 5 6 , 0 7 81 5 , 3 --------2C C S 6 7 -5 6 8 1 5 0 5 , 1 7 01 5 , 4 0 , 0 1 0 0 , 0 0 3 0 , 0 5 0 ± 0,025 ≤ 0,054 0,141 ± 0,022 3H C 3 N 9 -8 8 1 8 8 1 , 4 6 81 9 , 6 0 , 0 3 5 0 , 0 0 3 ≤ 0,049 0,205 ± 0,027 0,499 ± 0,022 4c -C 3 H 2 2 0 0 -1 1 1 0 8 2 0 9 3 , 5 5 96 , 4 0 , 0 0 8 0 , 0 0 2 ≤ 0,037 ≤ 0,040 0,050 ± 0,016 5S O 2 8 1 7 -8 0 8 8 3 6 8 8 , 0 9 33 6 , 7 0 , 0 0 8 0 , 0 0 2 0 , 0 5 0 ± 0,018 0,064 ± 0,020 0,178 ± 0,016 6U L 8 3 7 5 9 , 0 2 9 --------7C H 3 O H 5 -1 -4 0 0 8 4 5 2 1 , 1 7 24 0 , 4 0 , 1 1 2 0 , 0 0 1 ≤ 0,025 0,337 ± 0,014 Chapter 5

The jet of Cep E

The jet molecular tracers

Over all the unbiased spectral survey some species present line profiles with a broad wing component that starts at the ambient cloud velocity (V LSR =-1 0 . 9 km s -1 )a n de x t e n d st o w a r d st h er e di nt h ec a s eo ft h eo ff s e tp o s i t i o n( 0 ,0 )a n d towards the blue in south position (-12, -18). This wing component decreases gradually in intensity until it merges at high velocities in the spectrum as a separate secondary peak of intensity.

For example, Figures 5.1 and 5.2 show the CO 1 -0 to 4 -3 and SiO 2 -1t o8-7s p e c t r at o w a r d st h ec e n t r a lp o s i t i o na n dt h es o u t ho u t fl o wl o b e . The observed profiles have emission up to high velocities (∼ 100 km s -1 from the rest velocity). As we already mention in chapter 4,f r o mt h el i n ep r o fi l e s it can be distinguish clearly two kinematical components, in addition to the ambient velocity one: i) wing emission, observed especially in the CO lines, prominent at velocities lower than ∼ 40 km s -1 from the rest velocity; ii) high velocity emission, particularly bright in the SiO lines compared to the ambient and wing emission, at velocities larger than ∼ 40 km s -1 from the rest velocity. This emission is relatively weak and broad between ∼-80 and ∼-20 km s -1 and much more intense and evident in the velocity range 30 -80 km s -1 . Note that the contribution of the CO wing emission at velocities ≥ 30 km s -1 decreases with increasing J: at J=4 there is almost no trace of CO wing

The jet molecular tracers emission. Analogously, for the southern lobe (offset -12", -18"), the gas forms ab r o a dw i n go fr e l a t i v e l ys l o wg a st h a ts t a r ta tt h ea m b i e n tc l o u dv e l o c i t y and decrease smoothly towards the negative velocities up to find a secondary intensity peak centered about -125 km s -1 . Similar (wing and high velocity) components have been observed in several sources and are usually considered as due to the entrained gas and the high velocity jet respectively [Bachiller et al., 2001]; [Jørgensen et al., 2007]; [Santiago-García et al., 2009]; [Hirano et al., 2010].

For the analysis, we separate the wing emission in moderate velocity (MV) and high velocity (HV) components towards the central position and in extremely high velocity (EHV) component for the south position. Adopted velocity ranges are listed in Table 5. 1 and illustrated in Fig 4.9. Note that in the table, the MV and HV components are label 'north' because the gas emission from these components traces the north ejection of material. Similarly, the EHV component is labed 'south'.

At the central position we could detected CO, SiO, HCO + , HCN and SO emission coming from the North-MV and North-HV components (see Figs 5.1 to 5.5). On the contrary, only CO, HCO + and HCN have been detected in the South-EHV component. Table 5.1 summarizes the situation.

At the central position, the SiO present the brightest emission from the North MV and HV components with peak intensities about 0.35 K from J transitions between 2 and 8 (See Figure 5.2). Our observations show that in both positions the HCO + emission has a line profile similar to the CO line profile, with intensities in the North-HV component around 0.15 K. On the contrary, at the offset position (-12, -18) SiO does not shows clearly the presence of a jet, but only a prominent wing component. As discussed by other authors (e.g. [Tafalla et al., 2010]), even a single velocity component may mask more than one physical component and only high resolution observations can help to fully disentangle the different physical components in the line of sight. To this end, the CO high spatial resolution channel map, obtained with the IRAM Plateau de Bure Interferometer (Lefloch et al. in preparation) and shown in Fig. 4.1,i so fg r e a th e l pi nu n d e r s t a n d i n g the origin and nature of the emission observed with the single-dish reported in the present work. (See section 4.2 to see a description of the maps.) Table 5.2 lists the detected lines for each species and their integrated emission in the North-MV, North-HV and South-EHV velocity interval components -12,-18) position. In both pannels the spectra of lowest transition were observed with IRAM-30m telescope and transition CO (4 -3) were observed with CSO for the (0,0) offset and with JCMT for the offset (-12,-18). Intensities are in units of antenna temperature (K). A vertical offset is added to each emission line to help the comparison. The vertical dotted lines show the velocity range where appear the EHV component associated with the jet. (See the text).

Derivation of the column density

Rotational diagrams

We start with a first level of analysis, by means of a rotational diagram analysis, performed using the CASSIS package (Centre d'Analyse Scientifique de Spectres Infrarouges et Submillimtriques; http://cassis.cesr.fr). For this analysis, the line emission was integrated between 30 and 80 km s -1 (North MV+HV component) for each line. The analysis gives as result the column densities ranging from 6.6×10 12 for the HCN to 1.2±0.1×10 16 for the CO. The rotational temperatures cover a range between 10 to 30 K (See table 5.3). The analysis was carried out taking into account a beam dilution correction for a source size of ∼3"×3".

In figures 5.7 and 5.8 are shown the rotational diagrams for the molecules that present emission from the jet. and SO spectra at (0,0) offset with detected jet emission at high velocities. Vertical dash line marks the ambient cloud velocity (∼ -10.9 km s -1 ), and vertical dotted lines mark the velocities at 30, 50 and 80 km s -1 .I nt h er a n g e3 0t o4 5k ms -1 we have the moderate velocity component of the north jet (North-MV), and between 50 to 80 km s -1 we have the high velocity component (North-HV). (right)T h es a m em o l e c u l a rt r a n s i t i o n sa t( -1 2 , -1 8 )o ff s e t . U n l i k et h es p e c t r ao nt h el e f t ,o n l yC Oa n dH C O + spectra show clearly detected jet emission at South-EHV. Vertical dotted lines mark the velocity limits between -100 and -150 km s -1 .

Ta b l e 5 . 3 : Molecular column densities and rotational temperatures derived by interpreting the emission between 30 to 80 km s -1 (MV and HV components together). The abundances are expressed with respect to H 2 and CO-normalized, i.e, X CO 4 =N/N CO ×10 4 .T h eb e a md i l u t i o nc o r r e c t i o nw a sp e r f o r m e dc o n s i d e r i n gas o u r c es i z e of 3" ×3".

Molecule

N (cm -2 )T ex (K) X CO 4 CO 1.2±0.1×10

North-MV component

To find the best fit solution, we adopted the following strategy for the North-MV component: 1) we first tried to best fit the observed SiO lines 1 with a single physical component; 2) the fit was not good enough and, therefore, we best fitted the SiO data with two physical components, where we fitted simultaneously the SiO and SO lines, to limit the degeneracy in the parameter space; 3) we then fit the other species by using the same two physical components derived in step 2.

As shown in Fig. 5.9,o n es i n g l ec o m po n e n td oe sn o tfi ta l lt h eo b s e r v e dS i O lines. In practice, while the J=5 to 8 lines are well fitted by a gas with a density of 2 -4 × 10 6 cm -3 and temperature 100-200 K, the 2-1 line requires the presence of colder and less dense gas. This second component is well constrained when considering together the SiO and SO emission. With this further constraint, the gas of the second component has a density (3 × 10 5 cm -3 )a n dat e m p e r a t u r e ∼ 20 K. The warm component covers about 4" while the cold one is slightly more extended, 5". These two components fit fairly well also the CO and HCN emission. Conversely, the HCO + emission cannot be fitted by neither of the two components, as it requires a less dense gas, of less than ∼ 5×10 4 cm -3 .L o o k i n ga t the HCO + spectra (Fig. 5.6), it is likely that this is because the HCO + North-MV line is largely contaminated by the entrained gas component. In the following, unfortunately, we will not consider any further, therefore, the HCO + case.

The warm and dense component has a size of 4" in good agreement with the dimensions of the cavity shown by the map of Fig. 4.1. We conclude that this physical component, the one marked II in the Tab. 5.5,describesthephysicaland chemical conditions in the cavity created at the base of the north jet. On the contrary, the component I probably describes the physical and chemical conditions of the entrained gas at large scale.

From a chemical point of view, we remark the following. First, both the SiO/CO and SO/CO abundance ratios are similar, within a factor 3, in the component II, associated with the cavity, and in the component I, associated with the entrained gas. In both components, SO is about five times more abundant than SiO. On the contrary, the HCN/CO ratio is very different in the two components, with the warm and dense component about about a factor 30 lower.

North-HV component

We followed a procedure similar to that used for the North-MV component. Also in this case, two physical components are required to fit the SiO, SO, CO and HCN lines (Tab. 5.4). In fact, the two physical components are approximatively Ta b l e 5 . 6 : List of the abundance with respect to CO of the species present in L1157-B1, CepE, L1448, I04166. The abundances are expressed with respect to H 2 and CO-normalized, i.e, X CO 4 =N/N CO ×10 4 . a Taken from [Tafalla et al., 2010]a n dc a l c u l a t e df r o m [Bachiller and Perez Gutierrez, 1997]. b, c Taken from [Tafalla et al., 2010]a tE H Vr e g i m e .

Cep E L1157-B1

Discussion

Comparison with other outflow sources

In this section, we compare the molecular composition obtained with the previous analysis with the composition observed towards low-mass protostars. Some examples of outflows line spectral surveys driven by a low-mass protostar have been carried out selecting some specific frequency line transitions e.g. [Bachiller, 1996], [Bachiller and Perez Gutierrez, 1997], [Bachiller et al., 2001] towards L1157-BI, [Jiménez-Serra et al., 2004]t o w a r d sL 1 4 4 8a n d [ Tafalla et al., 2010] towards both, L1448 and I04166. Mainly in CO, SiO, SO, H 2 OandCH 3 OH molecular lines.

As we already mentioned, the L1157 outflow is the prototype of the so-called chemically active outflows driven by a Class 0 protostar. [Bachiller et al., 2001] observed this object with IRAM-30m telescope and report 18 molecular transition detected towards the blue lobe. In later works [Codella et al., 2010]d e t e c t e da total of 27 emission lines in their survey towards B1 as part of the Herschel key program CHESS (Chemical HErschel Surveys of Star forming regions). [Tafalla et al., 2010] observed towards two YSO's, L1448 and I04166 in 33 frequencies corresponding to CO, 13 CO, SiO, SO, CH 3 OH, HC 3 N, CS, HCO + , HCN, H 2 CO, SO 2 ,H 2 S, SiS and HNC also with IRAM-30m (See Table 5.6).

In Table 5.6 we present the abundance values of Cep E in the HV regime For exemple, the chemical richness revealed in this source can b e compared with other low-and high mass sources, helping to understand similarities and differences between the different mass ranges.

Despite the difficulties in estimating the abundances, that often suffer large uncertainties due to the poorly known physical conditions, geometry, and sizes, this is of paramount importance for understanding the interplay between gas phase and grain surface chemistry in protostellar envelopes, and eventually the chemical evolution from hot cores/corinos to protoplanetary disks.

To derive the physical prop erties of a molecular cloud from line observations, one has to know or assume the molecular excitation conditions. The Sobolev approximation neglecting nonlocal radiative interactions is often used for a simple estimate. It is based on the existence of large velocity gradients throughout the cloud. Ossenkopf [1997].

It is assumed that the line widths are due to the existance of large velocity gradients throughout the cloud, so that the variation of the velocity between points relatively close is bigger than the local velocity dispersion. In other words, a line is Doppler shifted across its own width on a short length. Whenever one can assume that all other physical quantities remain constant at this scale, the radiative transfer integral can be solved locally for each point.

The statistical equilibrium equations describe the variation of the population of level i with energy E i due to both collisional and radiative transitions.

dn i dt = - j =i R ij n i + j =i R ji n j - j =i C ij n i + j =i C ji n j (B.1)
where, R ij is the probability of radiative transition from i to j level. Both, R ij R ji can bewritten in terms of Einstein coefficients. The C ij are the collisional transition probability. C ij = nγ ij and C ji = nγ ji where γ ij and γ ji are the excitation and de-excitation collisional coefficients; and n is the number density of the collisors, these can be H 2 , H or He. The collisional coefficients depend on the temperature of the gas, the density of collisors, geometric cross section of the collision and the velocity of the particles.

R ij = A ij +4π JB ij (E i >E j ) (B.2) R ji =4π JB ji (E i >E j ) (B.
The calculation of the collisional coefficients, is a very complicated task, which requires the most advanced knowledge of quantum molecular physics coupled with the most advanced numerical techniques and; as far possible, very difficult laboratory experiments.For this reason, at present, we know the collisional coefficients of a relatively small number of molecules: CO, H 2 ,a n da p p r o x i m a t e de s t i m a t e s for H 2 O, H 2 CO, CN, HCN, etc. with a degree of accuracy that can be poor.

For a stationary situation, dn i /dt=0, ∀i. Thus, there is a equation system that allow, in principle, to compute the population of the levels. However, in eq.B.2 to eq.B.4 appear the J in each point under consideration and that depends in turn of the population levels. Thus, the solutions of this equation system can be complicated. This can be avoid using the Sobolev approximation. Then, the equation system is split into many independent equations. It is assuming that the velocity field is such that there is no radiative interaction between distant points, so that a photon emitted by a molecule that is not absorbed locally by another molecule in the near zone may escape of the cloud. The points in the same velocity zones have similar physical conditions and therefore the same excitation conditions.

In the radiative interaction limit, the value of the mean intensity at any point in the cloud when, one integers over the line profile depends only on the local value of source function and of the probability that a photon emitted in that point can escape from the cloud without interact. Sobolev's approach therefore implies a localized radiative transfer, ie, a single point interact radiatively with their surroundings so that J does not depend on what happens outside of that

  and a new paradigm (right panel) emerged. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 1.2 Stages of the low-mass star formation process. Taken from Hogerheijde 1998 after Shu et al 1987. . . . . . . . . . . . . . . . . . . .

  E-mm dust temperature (left panel)a n dH 2 density (right panel)p r o fi l e sf r o mt h eb e s t -fi tm o d e l . F o rd e t a i l ss e e[ Crimier et al., 2010]. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 70 3.5 (left) HiRes (Spitzer) maps of the inner 10" region around the Cep Ep r o t o s t a r . T h ec r o s s e sr e p r e s e n tt h ep o s i t i o n so ft h em md o uble source from [Moro-Martín

  . . . . . . . . 73 3.6 (top) Model 1: In the first panel shows the standard chemical model for Cep E assuming a binding energy of 1100 K. The C 18 O abundance is shown in black, N 2 H + in red, and N 2 D + in blue. In the center; comparison between the predicted C 18 O1→0l i n ei ntensities (continuous line) and the observed values (filled squares). At the right: the same for N 2 H + 1 →0. (middle top) Model 2. (middel bottom) Model 3 (bottom) Model 4. Taken from[Alonso-Albi et al., 2010]. . . . . . . . . . . . . . . . . . . . . . . . . . . . 75 3.7 K' band image of the Cep E region in a ∼ 8 × 3 field. In the center of the image we see the double lobe morphology of the Cep Eo u t fl o w . T a k e nf r o m[ Hodapp, 1994]. . . . . . . . . . . . . . . . 76 3.8 Contour plot of Cep E in H 2 1-0 S(1) emission line. The solid line trace the precession model plotted over the 1-0 S(1) line contours of Cep E that can explain the wiggle structure in the southern lobe and the regular sideways positional offsets of the bows of the northern lobe. Taken from [Eisloffel et al., 1996]. . . . . . . . . . . 77 xii List of Figures 3.912 CO(2-1) emission at three different positions respect to the reference position of the IRAS 23011+6126 source. The top panel shows line emission comes from north lobe, a first component is centered at the ambient velocity (≈-11 km s -1

  . . . . . . . . . . . . . . . . . . . . . . . . . . . LIST OF FIGURES xiii 3.12 Panel (a) shows a H 2 1-0 S(1) contour map overlaid with grayscale representations of the Hα image. In panel (b) again the H 2 1-0 S(1) contour map overlaid with the [S II] λλ6717, 6731 image. Panels (c) and (d) are closed-ups to the bowshock region. They show the Hα and [S II] 6717, 6731 images (respectively) as dark contour maps, over the H 2 light contours. Taken from [Ayala et al., 2000]. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 81 3.13 The upper panels displays the H 2 data. The lower panels displays the CO rotational fluxes. Taken from [Smith et al., 2003]. (left) Planar shock models for Cep E south. (a) Ah o tC -s h o c k component with 32 km s -1 provides a fit to the vibrationally excited columns but does not contribute to the CO rotational lines. (b) Ac o o lC -s h o c kw i t hs p e e d1 1k m s -1 provides a fit to the rotationally excited H 2 and CO but does not contribute to the vibrational lines. (right) C-type bow shock model for Cep E South 1 with speed 120 km s -1 . In all the cases the density is 10 5 cm -3 .T h eH 2 models values are displayed for the ground (solid line),fi r s t(dashed line) ans second (dot-dashed line) vibrational levels. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 84 3.14 The upper panels displays the H 2 data. The lower panels displays the CO rotational fluxes. Taken from [Smith et al., 2003]. (left) C-type bow shock model for Cep E North (Nd).T h em o d e l bow speed is 120 km s -1 .( right) J-type bow shock model for Cep E North (Nd) with shape parameter s=1.5 and a speed of 60 km s -1 ................................ 85 4.1 CO (2-1) velocity maps taken with PdB toward Cep E region. (top)T h ec r o s sm a r kt h ep o s i t i o no ft h ec e n t r a ls o u r c e . T h eg a s emission comes from two main components, the cavity between 30 and 45 km s -1 and the jet between 50 and 80 km s -1 .( bottom)Cep Es o u t hj e t ,t h ee m i s s i o nc o m e sm a i n l yf r o mt h eg a sbe t w e e n-1 0 0 and -130 km s -1 .. . . . . . . . . . . . . . . . . . . . . . . . . . . 90 xiv List of Figures 4.2 SiO (2-1) velocity maps taken with PdB toward Cep E region.

  spectra in blue...................... 4.7 Cep E CO (2-1) spectra. This image illustrates the addition of errors of each channel along the velocity interval from v 1 to v 2 .. . 4.8 Upper level energy distribution in the Cep E unbiased spectral survey. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

  km s -1 (north jet). . . . . . . . . . . . . . . . . . . . . . . . . . . . LIST OF FIGURES xv 4.10 Overview of all the detected lines in the 3 mm band from 80-116 GHz. At the bottom the (0,0) position and at the top the (-12, -18) position. The temperature axis is zoomed to appreciate the line richness in both position. Signals are in antenna temperature (K). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4.11 Overview of all the detected lines in the 2 mm band from 164-170 GHz. At the bottom the (0,0) position and at the top the (-12, -18) position. The temperature axis is zoomed to appreciate the line richness in both position. Signals are in antenna temperature (K). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4.12 Overview of all the detected lines in the 1.3 mm band from 207-260 GHz. At the bottom the (0,0) position and at the top the (-12, -18) position. The temperature axis is zoomed to appreciate the line richness in both position. Signals are in antenna temperature (K). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4.13 Overview of all the detected lines in the 0.9 mm band from 260-350 GHz. At the top the (0,0) position and at the bottom the (-12, -18) position. The temperature axis is zoomed to appreciate the line richness in both position. Signals are in antenna temperature (K). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4.14 Rotational diagrams of CO, SiO and SO. The squares are the data and the red line are the fit. In some cases the error bar is smaller than the symbol.

4.25 1 1

 11 .3 mm band from 207-260 GHz. Signals are in antenna temperature (K). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4.26 0.9 mm band from 260-350 GHz. Signals are in antenna temperature (K). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4.27 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . Spectra of the CO emission lines observed toward both positions.

  -1 .. . . . . . . . . . . . . . . . . . . . . . . . . . .184 xviii List of Figures 5.7 Rotational diagrams of CO and SiO molecules. The squares are the data and the red line are the fit. In some cases the error bar is smaller than the symbol. The resulting column density and rotational temperature are listed in table 5.3.. . . . . . . . . . .185 5.8 Rotational diagrams of SO, HCO + and HCN molecules. The squares are the data and the red line are the fit. In some cases the error bar is smaller than the symbol. The resulting column density and rotational temperature are listed in table 5.3.. . . . . 186 5.9 Observed (red points) and best fit model predictions (green line) of the SiO lines in the North-MV (left) and North-HV (right) components (Tab. 5.1). The blue lines show the warm and dense components (marked as Component II in Tab. 5.4)w h i l et h em a g e n t a lines at the left bottom corners show the less dense and colder components (marked as Component I in Tab. 5.4.......... 190 5.10 Observed (red points) and best fit model predictions (green line) of the SiO lines in the South-EHV component observed towards the south lobe (Tab. 5.1). The magenta line shows the warm and dense component (marked as Component II in Tab. 5.4)whilethe blue line shows the less dense and colder component (marked as Component I in Tab. 5.4.. . . . . . . . . . . . . . . . . . . . . .191 5.11 Comparison between the normalized CO abundances found by [Tafalla et al., 2010]( red diamonds)t o w a r d sL 1 4 4 8a n do u rv a l u e s from the LVG analysis (stars) for the HV I (top)andMVIregime (bottom). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 193 5.12 Comparison between the normalized SiO abundances found by [Tafalla et al., 2010]( red diamonds)t o w a r d sL 1 4 4 8a n do u rv a l u e s from the LVG analysis (stars) for the HV I (top)andMVIregime (bottom). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 194

1. 2 Figure 1 . 1 :

 211 Figure 1.1: Before the 80's, the collapse was treated in the spherical symmetry approximation (left panel). With the discovery of the outflow phenomena and circumstellar disks, it was clear that the spherical approximation was not valid and a new paradigm (right panel)e m e r g e d .

Figure 1 . 2 :

 12 Figure 1.2: Stages of the low-mass star formation process. Taken from Hogerheijde 1998 after Shu et al 1987.

Figure 1

 1 Figure 1.3: Fragmentation is a multi-scale pro cess from spiral arms of a galaxy down to protostars whose characteristic lengths spanning from 10 kpc to 100 AU. Taken from[START_REF] Schulz | From Dust To Stars Studies of the Formation and Early Evolution of Stars,P r a w i svii[END_REF]

  Al a r g en u m b e ro fp r e -s t e l l a rc o n d e n s a t i o n sh a v eb e e no b s e r v e d ,b o t h in molecular line tracers of dense gas such as NH 3 ,C S ,N 2 H + and HCO + (e.g.

Figure 1 . 6 :

 16 Figure 1.6: (top) Interferometer map of the molecular jet in the HH211 outflow. The CO J=2-1 line (white lines)[Gueth and Guilloteau, 1999]i so v e r l a i do naf a l s ec o l o ri m a g eo fH 2[McCaughrean et al., 1994]. The red lines delineate the 230 GHz continuum emission from dust concentration around the protostar. (bottom) In color scale we see the near infrared H 2 emission, in contours blue and red there are the blue-and red-shifted emission of SiO, the green contours mark the dust emission at 870 micron surrounding the protostar.[Lee et al., 2007]

Figure 1 . 8 :

 18 Figure 1.8: Overabundance factors with respect to dense core values for a number of molecules towards the outflows from L1448-mm and IRAS 04166+2706. Data fromTafalla et al. (2010). Taken from[START_REF] Tafalla | IAU Symposium,V o l .2 8 0o fIAU Symposium[END_REF]

Figure 1 . 9 :

 19 Figure1.9: left:C O2 →1e m i s s i o nf r o mt h eL 1 1 5 7o u t fl o w . right:M a p so fi n t e g r a t e di n t e n s i t yf o rd i ff e r e n t molecular species towards the blue (southern) lobe of the L1157 outflow. Note the different extension and peak position of each species. Taken from[Bachiller and Perez Gutierrez, 1997].

Figure 1 . 10 :

 110 Figure 1.10: HH 1/2. These are among the brightest HH objects in the sky, they form bow shocks in a highly collimated bipolar outflow with dimension of 2.5 arsec (0.34 pc), symmetrically opposite from a young star which is ejecting material along its polar axis.

Figure 1 . 14 :

 114 Figure 1.14: The structure of J shocks.

Figure 1 . 15 :

 115 Figure 1.15: The structure of C shocks.

Figure 2 . 1 :

 21 Figure 2.1:In 1931 Karl Jansky built the first radiotelescope. This 14.6 meter rotatable, directional antenna system was designed, to investigate interference problems on long distance communication circuits. Using this antenna, Jansky was able to make the first positive identification of radio waves of extra-terrestrial origin.

Figure 2 . 4 :

 24 Figure 2.4: Typical radio telescope, a large parabolic antenna collects incoming radio waves and focuses them onto a smaller antenna called the feed horn. The signal is then carried to the radio receiver.

Figure 2 . 6 :

 26 Figure 2.6: IRAM30-m antenna at Pico Veleta in Spain.

Figure 2 . 7 :

 27 Figure 2.7: Gain elevatin dependence.

Figure 2 . 8 :

 28 Figure 2.8: Overview of EMIR bands. Possible band combinations: E090, E150, E230, E330, E090&E150, E090&E230 and E150&E330.

Figure 2 . 9 :

 29 Figure 2.9: Atmospheric transmission for 2mm and 4mm of precipitable water vapors at IRAM 30m. The EMIR and HERA bands are marked together with the frequencies of a few important molecular transitions.

Figure 2 . 10 :

 210 Figure 2.10: Herschel space observatory.

Figure 3 . 1 :

 31 Figure 3.1: The Cepheus OB3 active star formation region overlaid on the map of visual extinction, obtained from 2MASS [Lombardi and Alves, 2001]. Large circles denote the clouds associated with young stars. The meaning of the different symbols are as follows: Filled triangles are T Tauri stars, filled squares are Herbig Ae/Be stars, filled circles are weak-line T Tauri stars, open squares are photometric candidate and possible PMS members, X are Hα emission stars, crosses are T Tauri candidates. For details see[START_REF] Kun | Star Forming Regions in Cepheus[END_REF].

Figure 3 . 2 :

 32 Figure 3.2: Cep OB3b region. In the image the molecular clouds Cep A to Cep F, the dark cloud Lynds 1211, as well as the most prominent associated young stars are labeled. In contours the distribution of the visual extinction[Dobashi et al., 2005]. The lowest contour of the extinction is at A V =1m a g ,a n dt h ei n c r e m e n ti s 0.7 mag. For details see[START_REF] Kun | Star Forming Regions in Cepheus[END_REF].

  In the interior of the Cep E molecular cloud is embedded the Cep E mm source, (hereafter called just Cep E). Cep E coordinates are α(2000) 23h 03m 13s, δ (2000) 61 • 42' 21". (See Figure 3.2.)

Figure 3 . 6 :

 36 Figure 3.6: (top) Model 1: In the first panel shows the standard chemical model for Cep E assuming a binding energy of 1100 K. The C 18 Oa b u n d a n c ei ss h o w ni nb l a c k ,N 2 H + in red, and N 2 D + in blue. In the center; comparison between the predicted C 18 O1→0l i n ei n t e n s i t i e s( c o n t i n u o u sl i n e )a n dt h eo b s e r v e dv a l u e s( fi l l e d squares). At the right: the same for N 2 H + 1 →0. (middle top) Model 2. (middel bottom) Model 3 (bottom) Model 4. Taken from[Alonso-Albi et al., 2010].

Figure 3 . 8 :

 38 Figure 3.8: Contour plot of Cep E in H 2 1-0 S(1) emission line. The solid line trace the precession model plotted over the 1-0 S(1) line contours of Cep E that can explain the wiggle structure in the southern lobe and the regular sideways positional offsets of the bows of the northern lobe. Taken from[Eisloffel et al., 1996].

Figure 3 . 9 :

 39 Figure 3.9: 12 CO(2-1) emission at three different positions respect to the reference position of the IRAS 23011+6126 source. The top panel shows line emission comes from north lobe, a first component is centered at the ambient velocity (≈-11 km s -1 )a n di te x t e n d st or e dv e l oc i t i e su pt or e a c has e c o n dc o m po n e n tt h a tpe a k at about 50 km s -1 .T h emiddle panel shows emission comes from the central source, the spectrum display a main component centered at ambient velocity and weak emission comes from the blue-shifted velocities. The bottom panel exhibit the spectrum comes from the south lobe. It shows again a bright component at ambient velocity and a second component that peak around -125 km s -1 .T h e s es e c o n dc o m p o n e n t sa th i g hv e l o c i t i e s are interpreted as evidence of the presence of jets.[Lefloch et al., 1996].

Figure 3 . 11 :

 311 Figure 3.11: The figure shows the 450, 850 and 1300 µmi m a g e so ft h eI R AS2 3 0 1 1 + 6 1 2 6r e g i o n. T h ed u s t emission trace the quadrupolar outflow observed by[Ladd and Hodapp, 1997]( S e et e x t ) . T a k e nf r o m[ Chini et al., 2001].

Figure 3 . 12 :

 312 Figure 3.12: Panel (a) shows a H 2 1-0 S(1) contour map overlaid with gray-scale representations of the Hα image. In panel (b) again the H 2 1-0 S(1) contour map overlaid with the [S II] λλ6717, 6731 image. Panels (c) and (d) are closed-ups to the bowshock region. They show the Hα and [S II] 6717, 6731 images (respectively)as dark contour maps, over the H 2 light contours. Taken from[START_REF] Ayala | [END_REF].

Figures 4 .

 4 Figures 4.1 and 4.2 show velocity maps of CO and SiO, giving the morphology of Cep E outflow. Towards the central position, at low velocity (≤ 30 km s -1 ), emission is associated with the large scale, extended outflow entrained material; at moderate (between 30 and 45 km s -1 )v elocit yemissionisassociated with the cavity at the base of the north jet; the high (≥ 45 km s -1 )v e l o c i t yi s associated with the north jet. Towards the south outflow lobe, (between -65 and -80 km s -1 )i tc a nb ed i s t i n g u i s hac o n c e n t r a t i o no fC O( 2 -1 )a tt h ea p e xo ft h e south jet and at extremely high velocity (≤-110 km s -1 )i sd e t e c t e de m i s s i o n associated with the south jet.

Figure 4

 4 Figure 4.4: (left) Aspectrawithfluctuationsinthecon tin uumbaseline,itisdifferen ttozeroalongthespectra.(right) The same spectra of the left but corrected, it has been removed a baseline in order to obtain a continuum close to zero.

Figure 4 . 5 :

 45 Figure 4.5: Examples of spikes present in the spectra. Note that the intensity of these spikes are 4 times brighter than the emission from the lines.

Figure 4 . 6 :

 46 Figure 4.6: Example of a ghost line: CS (5-4) spectra (black line) with a contamination in the wing by CO. For comparison we overplot the CS (6-5) spectra in blue.

Figure 4

 4 Figure4.9: Three kinematical components were found: the ambient gas at ∼ -10.9 km s -1 ,t h ee n t r a i n e d gas between -90 → -12 and -9 → +30 and the jet, detected at -125 km s -1 (south jet) and at around +50 km s -1 (north jet).

Figures 4. 19

 19 Figures 4.19to 4.26 present the spectra with the labels of each line identification listed in Tables 4.11 and 4.12 for (0,0) and(-12,-18) offset respectively.

  Figures 4.19to 4.26 present the spectra with the labels of each line identification listed in Tables 4.11 and 4.12 for (0,0) and(-12,-18) offset respectively.

Figure 4 . 16 :

 416 Figure 4.16: Some brightest ion spectra at both positions. At the top we show the N 2 H + lines at the middle the HCO + lines and at the bottom the DCO + spectra. Vertical dotted line indicates the systemic velocity.

Figure 4 . 17 :

 417 Figure 4.17: Assembly of the line profiles of deuterated molecules observed towards Cep E (Offset 0,0). The dash line indicates the systemic velocity.

Figure 4 . 19 :

 419 Figure 4.19: 3m mb a n dt a k e nt o w a r dC e pEr e g i o n . E a c hl i n ei sl a b e l e dw i t ht h em o l e c u l a rs p e c i e( S e eT a b l e4.11). Features appearing as absorption spikes were removed because they are artifacts. Signals are in antenna temperature (K).

Figure 4 . 20 :

 420 Figure 4.20: 2m mb a n df r o m1 6 4 -1 7 0G H z . E a c hl i n ei sl a b e l e dw i t ht h em o l e c u l a rs p e c i e( S e eT a b l e4.11). Features appearing as absorption spikes were removed because they are artifacts. Signals are in antenna temperature (K).

Figure 4 .

 4 Figure 4.21: 1.3 mm band from 207-260 GHz. Each line is labeled with the molecular specie (See Table4.11). Features appearing as absorption spikes were removed because they are artifacts. Signals are in antenna temperature (K).

Figure 4 .

 4 Figure 4.22: 0.9 mm band from 260-305 GHz. Signals are in antenna temperature (K).

Figure 4 .

 4 Figure4.23: 3m mb a n df r o m8 0 -1 1 6G H z . S i g n a l sa r ei na n t e n n at e m p e r a t u r e( K ) .

Figure 4 . 24 :

 424 Figure4.24: 2m mb a n df r o m1 6 4 -1 7 0G H z . S i g n a l sa r ei na n t e n n at e m p e r a t u r e( K ) .

  , HCO + , HCN, SO South-EHV -150 to -100 -12,-18 south jet CO, HCO + , HCN Other molecules like SO presents important emission also around 0.2 K for the North-HV peak emission at central position (See figure 5.5).

Figure 5 . 1 :

 51 Figure5.1: Spectra of the CO emission lines observed toward both positions. At the left (0,0) position, and at the right the(-12,-18) position. In both pannels the spectra of lowest transition were observed with IRAM-30m telescope and transition CO (4 -3) were observed with CSO for the (0,0) offset and with JCMT for the offset(-12,-18). Intensities are in units of antenna temperature (K). A vertical offset is added to each emission line to help the comparison. The vertical dotted lines show the velocity range where appear the EHV component associated with the jet. (See the text).

Figure 5 . 3 :

 53 Figure 5.3: HCO + jet emission from both; the source and bowshock positions. A vertical offset is added to each emission line to help the comparison. The vertical dotted lines show the velocity range components.

Figure 5 . 4 :

 54 Figure 5.4: HCN jet emission from both; the source and bowshock positions. A vertical offset is added to each emission line to help the comparison. The vertical dotted lines show the velocity range components.

Figure 5 . 5 :

 55 Figure 5.5: SO jet emission from both; the source and bowshock positions. A vertical offset is added to each emission line to help the comparison. The vertical dotted lines show the velocity range components.

Figure 5 . 6 :

 56 Figure 5.6: (left)C O ,S i O ,H C N ,H C O +and SO spectra at (0,0) offset with detected jet emission at high velocities. Vertical dash line marks the ambient cloud velocity (∼ -10.9 km s -1 ), and vertical dotted lines mark the velocities at 30, 50 and 80 km s -1 .I nt h er a n g e3 0t o4 5k ms -1 we have the moderate velocity component of the north jet (North-MV), and between 50 to 80 km s -1 we have the high velocity component (North-HV). (right)T h es a m em o l e c u l a rt r a n s i t i o n sa t( -1 2 , -1 8 )o ff s e t . U n l i k et h es p e c t r ao nt h el e f t ,o n l yC Oa n dH C O + spectra show clearly detected jet emission at South-EHV. Vertical dotted lines mark the velocity limits between -100 and -150 km s -1 .

  ν)i st h el i n ep r o fi l e .

  In 1931 Karl Jansky built the first radiotelescope. This 14.6 meter rotatable, directional antenna system was designed, to investigate interference problems on long distance communication circuits. Using this antenna, Jansky was able to make the first positive identification of radio waves of extra-terrestrial origin. . . . .2.2 Schematic diagram whichshows how the specific intensity is defined. 2.3 Schema of an antenna pattern. . . . . . . . . . . . . . . . . . . . . Top Atmospheric transmission of the Earth's atmosphere for electromagnetic radiation. Bottom Radiation transmitted by the atmosphere between 0.2 and 70 µm. The pannel shows the absorption bands in the Earth's atmosphere and above the individual absorption spectrum for major greenhouse gases plus Rayleigh scattering. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 2.6 IRAM30-m antenna at Pico Veleta in Spain. . . . . . . . . . . . . 2.7 Gain elevatin dependence. . . . . . . . . . . . . . . . . . . . . . .

	x	List of Figures
		2.5 2.8 Overview of EMIR bands. Possible band combinations: E090,
		E150, E230, E330, E090&E150, E090&E230 and E150&E330. . .

.................... 1.14 The structure of J shocks. . . . . . . . . . . . . . . . . . . . . . .

1.15 The structure of C shocks. . . . . . . . . . . . . . . . . . . . . . . 2.1 2.4 Typical radio telescope, a large parabolic antenna collects incoming radio waves and focuses them onto a smaller antenna called the feed horn. The signal is then carried to the radio receiver. . . . . . . . 2.9 Atmospheric transmission for 2mm and 4mm of precipitable water vapors at IRAM 30m. The EMIR and HERA bands are marked together with the frequencies of a few important molecular transitions. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 2.10 Herschel space observatory. . . . . . . . . . . . . . . . . . . . . . .

  The resulting column density and rotational temperature are listed in table 4.4.. . . . . . . . . . . . . . . . .106 4.15 Rotational diagrams of HCO + , HCN, CH 3 OH and H 2 CO. The squares are the data and the red line are the fit. In some cases the error bar is smaller than the symbol. The resulting column density and rotational temperature are listed in table 4.4.I nt h e CH 3 OH rotational diagram are grey lines that indicate that there are transitions at different Eup for a same frequency. . . . . . . . 4.16 Some brightest ion spectra at both positions. At the top we show the N 2 H + lines at the middle the HCO + lines and at the bottom the DCO + spectra. Vertical dotted line indicates the systemic velocity. Assembly of the line profiles of deuterated molecules observed towards Cep E (Offset 0,0). The dash line indicates the systemic velocity. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4.18 Rotational diagrams of CH 2 DOH, DCN, N 2 D + , HDCO and D 2 CO .Thesquares are the data and the red line are the fit. The resulting column density and rotational temperature are listed in table 4.9. 4.19 3 mm band taken toward Cep E region. Each line is labeled with the molecular specie (See Table 4.11). Features appearing as absorption spikes were removed because they are artifacts. Signals are in antenna temperature (K). . . . . . . . . . . . . . . . . . . .

	xvi	List of Figures
	4.17

4.20 2 mm band from 164-170 GHz. Each line is labeled with the molecular specie (See Table 4.11). Features appearing as absorption spikes were removed because they are artifacts. Signals are in antenna temperature (K). . . . . . . . . . . . . . . . . . . . . . . . . 4.21 1.3 mm band from 207-260 GHz. Each line is labeled with the molecular specie (See Table 4.11). Features appearing as absorption spikes were removed because they are artifacts. Signals are in antenna temperature (K). . . . . . . . . . . . . . . . . . . . . . . 4.22 0.9 mm band from 260-305 GHz. Signals are in antenna temperature (K). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4.23 3 mm band from 80-116 GHz. Signals are in antenna temperature (K). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4.24 2 mm band from 164-170 GHz. Signals are in antenna temperature (K). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

  Summary of the molecules detected. The first and second columns number indicate the number of lines detected toward the central and the south position respectively. Between brackets the total number of transitions in the case of multiplets are indicated. In the third column are the reference numbers from Tables4.11 and 4.12...................................113 4.6 Summary of ions detected. Between brackets the total number of transitions in the case of multiplets are indicated. In the third column are the reference numbers from Tables 4.11 and 4.12.... 115 EHV component). . . . . . . . . . . . . . . . . . . . . 182 5.3 Molecular column densities and rotational temperatures derived by interpreting the emission between 30 to 80 km s -1 (MV and HV components together). The abundances are expressed with respect to H 2 and CO-normalized, i.e, X CO 4 =N/N CO ×10 4 .T h eb e a md ilution correction was performed considering a source size of 3" ×3".184 5.4 Temperature, density and sizes of the velocity components of Tab.

	xx	List of Tables	LIST OF TABLES	xxi
	5.2 Integrated flux intensities between 30 to 45 km s -1 (MV compo-nent) and 50 to 80 km s -1 (HV component) towards the central position, and between -100 to -150 km s -1 towards the south out-4.5 118 flow lobe (
		4.10 Ratios between some of the deuterated molecules and its corre-
		sponding non-deuterated molecule. In the first column there are
		the values obtained in our survey towards (0,0) position. The
		values listed here correspond to the ambient gas component (enve-
		lope) towards offset (0, 0). See the text. The other columns show
	4.3 Component velocity limits. The jet component present two com-the same ratios towards Class 0 low-mass protostars, pre-stellar
	ponents, MV (moderate velocity) and the HV (high velocity) gas. cores, L1157-B1 and IM protostars in order to compare our results. 118
	4.4 Molecular column densities and rotational temperatures derived 4.11 Overview of the Cep E spectral survey, 0,0 offset . . . . . . . . . . 152
	towards Cep E for ambient velocity component (See Figs 4.14 4.12 Overview of the Cep E spectral survey, -12, -18 offset . . . . . . . 164

2.1 The largest radio-mm-submm telescopes and interferometers to observe in radio to submm wavelengths. . . . . . . . . . . . . . . . . 2.2 IRAM 30m telescope parameters . . . . . . . . . . . . . . . . . . 2.3 List of some satellites that work(ed) from near-infrared to radio wavelengths. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 3.1 Cep E protostar parameters. . . . . . . . . . . . . . . . . . . . . . 3.2 Cep E envelope parameters calculated from different models. . . . 3.3 Summary of the dust radiative transfer analysis for Cep E protostar. Adapted from [Crimier et al., 2010]. . . . . . . . . . . . . . . 3.4 Summary of the Cep E outflow properties. . . . . . . . . . . . . . 4.1 Observational parameters. . . . . . . . . . . . . . . . . . . . . . . 4.2 Frequency ranges and interpolation functions used for derive the Beff for each line frequency in the Cep E survey. The Beff values were taken from IRAM"s website. . . . . . . . . . . . . . . . . . . and 4.15). The abundances are expressed respect to H 2 and COnormalized, i.e, X CO 4 =N/N CO ×10 4 .................. 4.7 Summary of deuterated molecules detected. Between brackets the total number of transitions in the case of multiplets are indicated. In the third column are the reference numbers from Tables 4.11 and 4.12.. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 116 4.8 Summary of complex organic molecules detected. Between brackets the total number of transitions in the case of multiplets are indicated. In the third column are the reference numbers from Tables 4.11 and 4.12.......................... 116 4.9 Column densities and T rot derived using a rotation diagram analysis. These values are used to compute the deuterated fraction (see Table 4.10). The abundances are expressed respect to H 2 and CO-normalized, i.e, X CO 4 =N/N CO ×10 4 .. . . . . . . . . . . . . .5.1 Velocity components. . . . . . . . . . . . . . . . . . . . . . . . . . 178 5.1.................................... 187 5.5 Column densities and abundances. Note: the abundances are computed assuming a CO abundance with respect to the H 2 of 10 -4 . 188 5.6 List of the abundance with respect to CO of the species present in L1157-B1, CepE, L1448, I04166. The abundances are expressed with respect to H 2 and CO-normalized, i.e, X CO 4 =N/N CO ×10 4 ..192 Introduction Intermediate-mass (IM) protostars (2 ≤ M * ≤8M ⊙ )a r et h el i n kbe t w e e n low and the high mass stars as they cover also an intermediate range of luminosities, densities and temperatures

  .1.

	Ta b l e 3 . 1 : Cep E protostar parameters.
	Parameter	Value	Reference
	Distance	730 pc	[Blaauw et al., 1959]
	Age	3×10 4 yr [Fro ebrich et al., 2003b]
	L bol	100 L ⊙	[Lefloch et al., 2011]
	Envelope mass	35 M ⊙	[Crimier et al., 2010]
	Envelope size	∼50"	[Crimier et al., 2010]

  .2.

	Ta b l e 3 . 2 : Cep E envelope parameters calculated from different models.
	Model	Menv	L bol	Reference
	grey body approximation		78 L ⊙	[Fro ebrich et al., 2003b]
	power-law density and temperature distributions		30 L ⊙	[Moro-Martín et al., 2001]; [Noriega-Crespo et al., 2004]
	dust temperature dependent models		100 L ⊙	[Lefloch et al., 1996]
	power-law density	35 M ⊙		[Crimier et al., 2010]
	grey body approximation	7 M ⊙		[Fro ebrich et al., 2003b]
	power-law density	13 M ⊙		[Moro-Martín et al., 2001]; [Noriega-Crespo et al., 2004]
	constant density over a spherical volume	10 M ⊙		[Ladd and Howe, 1997]
	dust temperature dependent models	18 M ⊙

  .3. Note. The numbers between curly brackets indicate the frequency at which the Feff was measured.

			Ta b l e 4 . 1 : Observational parameters.
	Band Frequency Resolution rms	HPBW	Feff	Calibration
		(GHz)	(km s -1 )( m K )( a r c s e c )	%	e r r o r ( α)
	E090	83-117	7.2-5.1	2-5	30-21	(86)95	10%
	E150	164-170	3.7-3.5	5	14.8-14.3 (145)93	15%
	E230	201-267	3-2.2	5-9	12.1-9.1 (210)94	20%
	E330	260-305 330-348	2.3-1.2 1.8-1.7	6-10 12-23	9.4-8 7.4-7	(260)88 (340)81	30% 30%

  Ta b l e 4 . 2 : Frequency ranges and interp olation functions used for derive the Beff for each line frequency in the Cep E survey. The Beff values were taken from IRAM"s website.

	Frecuency range	Beff	Interp olation function
	(GHz)	%
	86 -145	80 -74	-0.12 x + 91.20
	145 -210	74 -63	-0.17 x + 98.54
	210 -260	63 -53	-0.20 x + 105.00
	260 -340	53 -35	-0.23 x + 111.50
	4.3.1.

  18 O + , HCS + ,N 2 D + and N 2 H + )a n dt o w a r d st h es o u t hl o be (HCO + ,DCO + ,H 13 CO + , HCS + and N 2 H + ). The spectra are shown in Fig. 4.16. The list of species are in Table 4.6.

  Summary of ions detected. Between brackets the total number of transitions in the case of multiplets are indicated. In the third column are the reference numbers from Tables 4.11 and 4.12.

			HCS +	2	1	12, 451
			N 2 D +	2	0	248, 480
			N 2 H +	2	2	63, 407
						454
	H 2	13 CO	6	1	174,202,391,424,435,537
	H 2 CCO	3	1		95, 274, 410
	Ta b l e 4 . 6 : Species	Position (0,0) Position (-12,-18)	No. Ref.
		DCO +	1(6)	1(6)		180
		HCO +	2	2		43, 373
		H 13 CO +	3	1		27, 355, 554
		HC 18 O +	3	0		11, 337, 527

  Column densities and Trot derived using a rotation diagram analysis. These values are used to compute the deuterated fraction (see Table4.10). The abundances are expressed respect to H 2 and COnormalized, i.e, X CO 4 =N/N CO ×10 4 .

	Ta b l e 4 . 9 : Molecule	N (cm -2 )T rot (K)	X CO 4
	H 2 CH 3 CCH 13 CO H 2 CCO	6 9 3 CH 2 DOH 4.5 +1 -1.1 ×10 13 1 6 1 CH 3 OH 8.0 +5 -3 ×10 13 DNC 2.1×10 11 HNC 2.1×10 12 7.8×10 10 N 2 D+ 1.0×10 12 N 2 H+ D 2 CO 1.3±0.2×10 12 HDCO 1.6 +0.4 -0.3 ×10 12 H 2 CO 1.5 +0.8 -0.5 ×10 13	,454 174,202,391,424,435,537 14,15,96,97,217, 218,271,272,388 55 +29 102.3± 29.7 -14 33 +34 -11 181.8 +126.0 -76.4 13 0.5 +0.03 -0.02 15 4.8 +0.3 0.2 17 0.2±0.01 95, 274, 410 21 2.3±0.1 53 +12 -8 3.0 +0.7 -0.6 34 +7 -5 3.6 +1.1 -0.8 84 +448 -39

  Ta b l e 5 . 5 : Column densities and abundances. Note: the abundances are computed assuming a CO abundance with respect to the H 2 of 10 -4 . Il ≤ 2 × 10 12 ≤ 1 × 10 -4 ≤ 1 × 10 -8 Il ≤ 5 × 10 11 ≤ 2 × 10 -5 ≤ 2 × 10 -9

	Species Comp.	N x (cm -2 )	N x /N CO	N x /N a H2
			North-MV component
	CO	I5 ×10 16 II 1×10 17	1.0 1.0	1×10 -4 1×10 -4
	SiO	I3 ×10 13 Il 2×10 13	6×10 -4 2×10 -4	6×10 -8 2×10 -8
	SO	I1 ×10 14 Il 1×10 14	2×10 -3 1×10 -3	2×10 -7 1×10 -7
	HCN	I5 ×10 13 Il 3×10 12	1×10 -3 3×10 -5	1×10 -7 3×10 -9
			North-HV component
	CO	I1 . 5 ×10 16 II 2×10 16	1.0 1.0	1×10 -4 1×10 -4
	SiO	I8 ×10 12 Il 1×10 13	5×10 -4 5×10 -4	5×10 -8 5×10 -8
	SO	I2 ×10 13 Il 4.5×10 13	3×10 -3 2×10 -3	1×10 -7 2×10 -7
	HCN	I6 ×10 12 Il 2×10 12	4×10 -4 1×10 -4	4×10 -8 1×10 -8
			South-EHV component
	CO	I4 ×10 16 II 2×10 16	1.0 1.0	1×10 -4 1×10 -4
	HCO +	I1 ×10 13	3×10 -4	3×10 -8
	HCN	I6 ×10 12	2×10 -4	2×10 -8
	SiO SO HCO +	16 -0.3 ×10 13 1.1 +0.4 3.1 +0.5 -0.4 ×10 13 9.6×10 12	17 +2 -1 16 +3 -2 27 +2 -1 10±48 1×10 4 9.2 +4 -3 25.8 +6 -5
	HCN		6.6×10 12	10±35 . 5

  a L1448 b I04166 c

	SiO	5	10	11	12.5
	SO	22.5	21	19	23
	CH 3 OH		191	28	20
	HCO +	31 . 5	0 . 8 8	0 . 9
	HCN	1	26.4	2.1	0.88
	H 2 CO		14.5	0.40	<1.4

Mixed with the gas, molecular clouds contain interstellar dust grains. These particles absorb light with wavelengths smaller than their diameters (∼ 0.1 µm) and reradiate this energy in infrared wavelengths. Regions where the dust prevents the passage of visible light from background stars are traditionally known as dark clouds, because they trace regions of strong obscuration in optical photographs.

Theory of the outflow phenomenon

www.iram.es/IRAMES/telescope/telescopeSummary/telescope summary.html

http://www.astro.uni-koeln.de/cgi-bin/cdmssearch

http://spec.jpl.nasa.gov/ftp/pub/catalog/catform.html

We used the SiO rather than the CO lines because the SiO lines are sensitive to higher densities than the CO lines can probe. Consequently, a fit with the CO lines would miss the high density gas probed by the high J SiO lines.

Population Diagram

Line identification

Ta b l e 4 . 3 : Component velocity limits. The jet component present two components, MV (moderate velocity) and the HV (high velocity) gas.

Component

Velocity limits (km s -1 ) Offset 0,0 Offset -12, -18 Ambient gas -10.9km s -1 -10.9km s -1 Entrained gas -9 → 30 -12 → -100 Jet gas: MV 30 → 45 Jet gas: HV 50 → 80 -100 → -150 tween the velocity ranges of the three different kinematical components in both positions.

The jet molecular tracers respectively. Note that the table lists also the J=12-11 and 13-12 transitions observed towards the south jet by [Gómez-Ruiz et al., 2012]. On the contrary, the table does not report the SO transition at 261.843 GHz because of the presence of one methanol line in the North-MV interval.

Finally, the profiles of selected lines in the 201-267 GHz interval, zoomed around the North-MV, North-HV and South-EHV velocity intervals, are shown in Fig. 5.6. Towards the central position, CO and HCO + have very similar spectra profile with (little) contribution from the wings, while SiO and SO profiles are rather dominated by the North-MV and North-HV components. HCN, on the contrary, presents a very important contribution from the wing component. In the case of SiO, the lines in the North-MV and North-HV components are even brighter than those from the ambient velocity. Clearly, SiO is the best tracer of the North-MV and North-HV components. Surprising enough, however, SiO is very weak in the South-EHV component towards the south jet.

Comp.

T n size (K) (cm

Non-LTE LVG analysis

Once we had a general idea about the column densities and rotational temperatures, we carried out a detail analysis using the non-LTE Large Velocity-Gradient code described in [START_REF] Ceccarelli | [END_REF]. The code was adapted to take into account the emission from a jet for the analysis of the North-HV and South-EHV components, whereas it uses a standard circular gaussian source for the analysis of the North-MV component. Specifically, the filling factor in the case of an infinite jet with a transversal size j a observed with a telescope beam HPBW is j a /HPBW. The collisional coefficients are taken from Yang et al. (2010) for CO, Turner et al. (1992) for SiO, Flower (1999) for HCO + ,D u m o u c h e le ta l . ( 2 0 1 0 ) for HCN and Green (1994) for SO.

We obtained a grid of models, where we varied the gas density and temperature, the species column density and the sizes of the source and jet, covering a large parameter space. We then best fitted the observations against the model predictions.

The results of the fits in the North-MV, North-HV and South-EHV components are reported in Tabs. 5.4 and 5.5.I nt h en e x ts e c t i o n s ,w ew i l ld i s c u s si n detail how they have been obtained ad the results.

Conclusions and perspectives

My thesis I was founded on the study of the young IM Cep E protostar. This is mainly because this kind of objects are crucial for understanding the overall star forming process as they constitute the link between low-and high-mass range (e.g., [Fuente et al., 2012]; [Crimier et al., 2010]).

Recent studies have started to shed light on the first phases of intermediate mass star formation. [Crimier et al., 2010] Cep E-mm is the only Class 0 IM protostar so far studied systematically over the entire mm bands. This type of study provides an excellent template upon which to base interpretations of other IM young protostars. The wealth of data obtained on this source allowed us to derive the physical conditions and chemical abundances of different species, both in the envelope and in the outflowing material.

Summary

The unbiased spectral survey results can be summarized as follows.

We have presented an unbiased sp ectral survey of the 3, 1, and 0.9mm bands accessible from the ground towards the IM Class 0 source Cep E. We used the IRAM-30 m telescope, performing about 80 h of observations at two offset Discussion positions: (0",0") and (-12", -18").

More than 500 lines with S/N ≥ 3a n du p p e re n e r g yl e v e l sl o w e rt h a n2 5 0 K have been identified by comparing the observed spectra with the JPL and CDMS catalogs. They correspond to 60 chemically distinct species.

From a chemical p oint of view, the Cep E-mm source present in general a chemical richness both in the envelope and in the outflowing material. In the envelope, we could observe strong deuteration as well as richness in complex molecules, typical of to hot core chemistry.

The analysis of the line profiles of this large set of identified species allow us to recognize that emission predominantly originates from three kinematical components: the ambient gas at ∼ -10.9 km s -1 ,theen trainedgasat(-90→ -12) and (-9 → +30) and the jet, mostly detected at -125 km s -1 (south jet) in the (-12,-18) position and at +50 km s -1 (north jet) in the (0,0) position (see chapter 4).

The line profiles that exhibit emission from the jet component are CO, SiO, SO, HCN and HCO + toward the position (0,0) and CO, HCN and HCO + toward the position (-12,-18) (see chapter 5).

To derive reliable estimates of the corresp onding chemical abundances from the jet component, we carried out a non-LTE radiative transfer modeling. The results of this analysis is summarized in Tables 5.4 and 5.5.

Perspectives

The work presented in this thesis, is one of the first steps toward a more complete description of the formation process of intermediate-mass stars. Evidently, we are far from fully understand the star formation process of this type of IM objects. The contribution of the present work gives some clues that should continue to be explored.

Appendix A Population Diagram

The population diagram is a technique to estimate the cloud physical conditions from the molecular line emission; like the column densities of the observed species and the temperature in the region covered by the telescope beam.

It is based on the relation between the observed emission line intensities and the column densities of such emitting molecules. An optically thin transition produces an antenna temperature that is proportional to the column density in the upper level of the transition being observed. If besides, all transitions are thermalized, it is possible to convert that particular column density in the total column density of the specie.

The term population or rotation diagram refer more precisely to a plot of the column density per statistical weight of a number of molecular energy levels, as afunctionoftheirenergyabo v ethegroundstate. Seefigure2.9.T h et r a d i t i o n a l use of this type of diagram has been to assume that the emission is optically thin and that LTE applies.

In the case of local thermodynamic equilibrium (LTE) in the gas, this will just be a Boltzmann distribution, so a plot of the natural logarithm of N u /g u versus E u /k will yield a straight line with a slope of 1/T, where g u is the statistical weight of the level u and E u is the energy above the ground state. The temperature inferred is often called the rotational temperature (T rot ). The temperature obtained would be expected to be equal to the kinetic temperature (T kin )i fa l l levels were thermalized [Goldsmith and Langer, 1999].

In a spectral survey, it may have sufficient number of transitions to check the 200 Population Diagram self consistency of the assumptions of optical thinness and thermalization. An either, an analysis of the excitation is possible.

However, this method presents many uncertainties associated with the assumption of complete thermalization and thin line emission.

The column density in the upper state transition of an observed molecule can be expressed in function of the integer intensity of the line by

where ν ul is the transition frequency between the upper and lower level, and A ul is the Einstein coefficient of spontaneous emission; the transition probability between the two levels, T bg (2,7 K) is the cosmic background radiation, T ex is the excitation temperature of the transition, τ is the optical depth of the transition, T B is the observed line bright temperature and J ν (T ) is the source function, defined as

Assuming that the emission is optically thin (τ 1) and that the emission comes from a region of the spectra where apply the Rayleigh-Jeand approximation (hν kT) and that the background radiation is negligible (T bg T ex ), the equation A.1 can be reduced as

where, it is written for short W= T B dv and all the constants replaced by γ u .

If the emission is not optically thin, an optical depth correction factor can be defined as

(A.4) so that gives the follow relations

In the case of arbitrary excitation, a different temperature may characterize the population of each level relative to that of the ground state or relative to that of any other level. The excitation temperature T ex is defined by the relative populations or column densities of any two levels i and j of statistical weights g i and g j and energies E i and E j relative to an arbitrary common reference, through

For a molecule in lo cal thermal equilibrium (LTE), all excitation temp eratures are the same, then the population function of the upper lever is given by Boltzmann N u = N Z g u e -Eu/kT (A.7)

where N is the total column density of the species in question, and Z is the partition fuction, given by Z =

All levels N i (A.8)

We can rewrite this equation to obtain

Obviously, the eq.A.9 is a equation of a line with slope -1/T and a y-intercept of ln Nln Z. y = ax+ b (A.10)

When the optical depth is not small then, using the correction factor in the eq.A.9 gives .11) here it still assuming populations of all levels are in LTE at temperature T.

205

surroundings, is a local function.