T Cavalié
email: cavalie@obs.u-bordeaux1.fr

F Billebaud

T Fouchet

E Lellouch

J Brillet

M Dobrijevic

Observations of CO on Saturn and Uranus at millimeter wavelengths: New upper limit determinations

Keywords: Planets and satellites, individual, Saturn, Uranus -Radio lines, solar system

) and Hesman et al. (2007) observing technique.

. As their fluorescence model did not take scattering effects into account, we suggest that their value could be overestimated.

Introduction

Water and carbon dioxide have been detected in the stratospheres of the four giant planets and Titan (Feuchtgruber et al. 1997(Feuchtgruber et al. , 1999;;Coustenis et al. 1998;Burgdorf et al. 2006). The presence of these condensable species above the tropopause cold trap (except CO 2 , which does not condensate on Jupiter and Saturn) implies the presence of an external supply of oxygen. Different possible sources have been proposed: infall of interplanetary dust particles (IDPs), sputtering from the planetary rings and icy satellites, or massive impacts of kilometer-sized comets or asteroids [START_REF] Moses | The Collision of Comet Shoemaker-Levy 9 and Jupiter[END_REF].

In order to discriminate between these various sources, different techniques can be used. First, the spatial distribution (altitude and latitude) of oxygenated compounds can trace localized inputs. [START_REF] Lellouch | The Collision of Comet Shoemaker-Levy 9 and Jupiter[END_REF] showed that the water in Jupiter is restricted to pressure levels higher than the 0.5-mbar level and that CO 2 is enhanced in the southern hemisphere: both species must then originate from localized, recent events: the Shoemaker Levy 9 (SL9) impacts. A recent analysis of submillimeter spectra of Odin and SWAS (Submillimeter Wave Astronomy Satellite) space telescopes seems to confirm the cometary origin of the stratospheric water on Jupiter (Cavalié et al. 2008). A second possibility is to determine which chemical species, water or a species with a C-O bond (CO, CO 2 , ...), is the main oxygen carrier to the stratospheres of the giant planets. From observations in the millimetric range or at 5-µm, Bézard et al. (2002) and [START_REF] Lellouch | The Collision of Comet Shoemaker-Levy 9 and Jupiter[END_REF]Lellouch et al. ( , 2005) ) have measured the CO abundance in the stratospheres of Jupiter and Neptune, and demonstrated that the ratios of the CO to the H 2 O deposition rates are larger Send offprint requests to: T. Cavalié than 30 and can be as large as 200. These large CO/H 2 O ratios strongly favor large impacts as the main source of oxygen in the stratospheres of Jupiter and Neptune, as shown by SL9 (Lellouch 1996;Moses 1996;[START_REF] Moses | The Collision of Comet Shoemaker-Levy 9 and Jupiter[END_REF].

In between, the situation of Saturn and Uranus is still unclear, because the CO mixing ratio in the stratospheres of these planets is still unknown or uncertain, and also since the CO origin cannot definitely be attributed to an external flux. Indeed, a second, internal, source of CO exists in giant planets in the form of convective transport from the deep warm atmosphere where the equilibrium CO abundance is larger than in the shallow atmosphere. This internal source is observed on Jupiter and Neptune (Bézard et al. 2002;Lellouch et al. 2005). Its measurement could be very powerful to constrain the O/H ratio and the convective activity in the giant planet deep interiors (Bézard et al. 2002).

On Saturn, CO has been detected from 5-µm spectroscopy at a level of 1 ppb (if uniformly mixed) by Noll & Larson (1991) but the authors were not able to measure independently the stratospheric and tropospheric mixing ratios. Therefore it is not possible to reliably estimate the relative contributions of the external and internal flux. Later, Rosenqvist et al. (1992) could only derive an upper limit of 1 × 10 -7 from their observations at millimeter wavelengths. On Uranus, CO has been recently detected at 5-µm (Encrenaz et al. 2004) from fluorescent emission. These authors derived a stratospheric mixing ratio of 3 × 10 -8 but within an uncertainty of a factor of two. In the troposphere, they derived an upper limit of 2 × 10 -8 . The latter improved the previous upper limits from Marten et al. (1993) and Encrenaz et al. (1996), which were ∼ 3 × 10 -8 and 5 × 10 -7 respectively. The large error bars in Encrenaz et al. (2004) did not allow the authors to distinguish the external from the internal flux. For Uranus again, the CO/H 2 O deposition rate ratio could not be measured, and the origin of the external oxygen remains unconstrained. We present new observations to help constrain the abundance of CO in the atmospheres of Saturn and Uranus.

Our observations are described in section 2. The radiative transfer model is presented in section 3 and we show our results in section 4. Discussion and conclusions are given in section 5.

Observations

We observed Saturn and Uranus on 16-18 September 2006 at the Institut de RadioAstronomie Millimétrique (IRAM) 30-m telescope, at Pico Veleta, Spain. The angular sizes of the planets were 16.6 arcsec for Saturn and 3.7 arcsec for Uranus. The Sub-Earth Point latitudes were respectively -14.3 • and -5.25 • for Saturn and Uranus.

We used simultaneously two receivers (A and B) which share two backends at 1 and 3-mm wavelength in order to observe the CO J=1→0 and J=2→1 lines at 115.271 GHz and 230.538 GHz (respectively). Each receiver covers a bandwidth of 512 MHz at 115 GHz and 1024 MHz at 230 GHz. Both receiver spectral resolutions were set at 1 MHz. The receivers were operated in a single side band mode. The side band rejection was 0.001 at 115 GHz and 0.01 at 230 GHz. The wobbler switching technique was used and the wobbler frequency was set to 0.5 Hz. System temperatures ranged from 320 K to 430 K during the first night (from 380 K to 500 K during the second night) on the 1-mm band and from 250 K to 420 K during the first night (from 260 K to 400 K during the second night) on the 3-mm band. The zenithal atmospheric opacity value was about 0.15-0.20 at 115 GHz and 0.25-0.30 at 230 GHz. A large absorption feature appears at 231.280 GHz on Uranus and Saturn spectra (see Figs. 1 and2). This line is the terrestrial ozone line.

On Saturn and Uranus, the expected CO tropospheric/low stratospheric absorptions are broad; they cover several GHz in frequency. This is the reason why we adopted the observation technique that is described in Lellouch et al. (2005) to observe these lines. This technique consists of using both receivers A and B in parallel in order to cover the widest band available: by taking an overlap of 1/4 of the bands (128 MHz at 115 GHz and 256 MHz at 230 GHz), it is possible to cover 896 MHz and 1 792 MHz for one single tuning at the CO(1-0) and CO(2-1) frequencies (respectively). For Saturn and Uranus, we covered 1 280 MHz at 115 GHz and 4 096 MHz at 230 GHz (see Figs. 1 and2). All the tunings were observed for 1/2 h up to 1 h in a cyclic scheme.

At 115 GHz, 3 contiguous spectra were obtained, while we observed 5 spectral bands around 230 GHz. The intense terrestrial O 2 atmospheric absorption around 118.750 GHz prevents ground-based observations rightwards of 115.5 GHz. This is the reason why we only have 3 bands at 115 GHz, respectively centered around 115.271 GHz,114.887 GHz and 114.503 GHz. The 115 GHz spectrum was not probed over 115.5 GHz.

Once all the observations were carried out, the individual spectra were fit together. Indeed, as two contiguous spectra had generally not the same continuum level, we averaged the spectra on their overlap part and rescaled them one to another. So, all the spectra have been rescaled according to the continuum of the first spectrum of each band. The final spectrum frequency ranges are 114. GHz. Finally, the spectra have been reduced like classical spectra. After removing a polynomial baseline of maximum degree 3, we removed the ripples, which had the greatest amplitude with a FFT procedure. The ripple frequencies that have been removed were the ones present on every single spectrum (at 115 GHz and 230 GHz, for Saturn and Uranus). As we performed no absolute calibration, the lines were interpreted in terms of line to continuum ratios. All the lines were smoothed down to a 16 MHz spectral resolution in order to reduce the noise level. The 3-σ levels we get, depending on the continuum value, are stored in Table 1.

The best 3-σ level we get are at 230 GHz for the CO(2-1) line either on Saturn or Uranus, after converting the antenna temperature noise level intensity of each spectrum into a brightness temperature noise level intensity.

(a)

Radiative transfer

We modeled the observed millimetric radiations with a line-byline non-scattering radiative transfer model. We computed synthetic spectra in the wavelength ranges of the CO(1-0) and (2-1) lines. The program respects the approximate spherical geometry of the planets (except the rings on Saturn) so that planetary disk and limb contributions were taken into account. We assumed a uniform distribution of all opacity sources and, for the sake of simplicity, we adopted mean thermal profiles for both planets since our beam size is comparable to or larger than the planetary disks. Thus, the total opacity on a line of sight only depends on the total thickness of the atmosphere on this line of sight. Therefore, we have a symmetry, in terms of opacity sources and temperature distributions, around the central point of the disk. This is the reason why we only computed synthetic spectra on a given number of points located on a planetary disk radius. Assuming local thermal equilibrium, the specific intensity I ν W.m -2 .sr -1 .Hz -1 on a given line of sight is:

I ν = τ ν (∞) 0 B ν (T (z))e -τ ν dτ ν (1)
The parameter τ ν is the optical depth, B ν (T ) is the Planck function and T (z) is the atmospheric vertical thermal profile.

The thermal profile we used come, respectively, from Ollivier et al. (2000) for Saturn and from Lindal (1992) for the troposphere and Orton et al. (1987) for the stratosphere for Uranus (see Fig. 3). Continuum opacity is dominated by H 2 -He-CH 4 collision-induced absorption (Borysow et al. 1985;Borysow & Frommhold 1986;Borysow et al. 1988). Following Moreno (1998), the opacity due to the far wings of ammonia and phosphine lines is also included (only NH 3 in the case of Uranus). We used the Davis et al. (1996) ammonia andFletcher et al. (2007) phosphine mixing ratio vertical profiles for Saturn. The mixing ratio profile of ammonia on Uranus is in agreement with Hofstadter & Muhleman (1989) measurements (q NH 3 = 7.8 × 10 -7 below the condensation level). We tested the CO mixing ratio vertical profiles inferred by Noll & Larson (1991) for the atmosphere of Saturn. The first one is based on an internal source of CO (1ppb, constant with altitude), while the other is the signature of an external supply of CO (2.5 × 10 -8 in the stratosphere and 10 -10 in the troposphere). For Uranus, the CO lines were computed from the Encrenaz et al. ( 2004) models. The authors tentatively favored an external origin for CO (3×10 -8 above the 100 mbar level) and set an upper limit on the magnitude of a tropospheric source (vertical profile with 2 × 10 -8 constant with altitude). Our 3-σ mixing ratio values are discussed in section 4. All these profiles are shown on Figs. 4 and5.

Spectroscopic parameters for NH 3 , PH 3 and CO were taken from Pickett et al. (1998). The line widths are determined by the collisional line widths for H 2 and He broadening. The broadening γ and temperature dependence exponent n values we took for NH 3 , PH 3 and CO are summarized in Table 2. All lines, except the NH 3 ones, are assumed to be Voigt-shaped. Following Moreno (1998), we take a modified Van Vleck and Weisskopf line profile for ammonia. We performed our analysis in terms of line to continuum ratios. We compared our continuum model to previously published observations. On Saturn, the continuum disk brightness temperature of our model is 141.6 K at 115 GHz and 140.0 K at 230 GHz. These numbers are consistent with [START_REF] Klein | Saturn's microwave spectrum: Implications for the atmosphere and the rings[END_REF] (∼140±15 K at 115 GHz and ∼135±11 K at 230 GHz). The good match of our continuum model to observations is confirmed with the Ulich (1981) data. The author refers a disk brightness temperature of 150±7 K at 86 GHz and 164±12 K at 142 GHz. As The upper limits we have derived for the internal source (dotted lines) and for the external source (dashed lines) are also displayed. Giant planets have rapid rotation periods. This induces a smearing of the lines which is linked to the rotation period and radius of the planets (see Table 3). Indeed, the radiations emitted on a line of sight are Doppler shifted because of the projected velocity of the rotating gas. The smearing effect is maximum when the Sub-Earth Point latitude is close to 0 • . To take this into account, the planetary disk was partitioned into a regular square grid. On each point of the grid, a synthetic spectrum is calculated by interpolating the spectra which were computed on a planetary disk radius. The spectra of the grid are then shifted in frequency, according to their location, and averaged with appropriate weights by taking the antenna pattern (2-D gaussian function) into account [START_REF] Bergin | [END_REF][START_REF] Lellouch | The Collision of Comet Shoemaker-Levy 9 and Jupiter[END_REF]. Although the CO lines are formed in the lower stratosphere, the lines appear to be in absorption, because the temperature where the line is formed is lower than the tropospheric continuum temperature. According to our model, we do not expect any observable stratospheric emission peak in the line core at 115 GHz and 230 GHz, unlike lines at higher frequencies [START_REF] Cavalié | Herschel Open time Key Program Workshop[END_REF]). Because the lines are formed at high pressure levels, the CO lines are broad (Encrenaz et al. 1996). Thus, the smearing was found to have only a small effect on the final spectrum as compared to the line width.

The rings of Saturn are sources of flux emission/absorption at millimeter wavelength (Dunn et al. 2005;de Pater & Dickel 1991). With a Sub-Earth Point latitude of -14.3 • at the time of the observations, the rings hide ∼ 10% of the northern part of the disk and add their own flux. This is the reason why we pointed the telescope towards the planet between the southern limb and the rings. The contribution of the rings cannot be neglected when analysing submillimeter and infrared data. But Melnick et al. (1983) andde Pater &Dickel (1991) show that brightness temperature of the rings (∼20-30 K) is well below the temperature of the disk at 1-mm and 3-mm (∼140 K). So the flux added by the ring contribution is not too important at these wavelength. Therefore, we do not include the effect of the rings in our analysis. Moreover, as we do not analyse the data in terms of absolute intensities, we assume that the attenuation of the flux of the disk due the rings does not change the result. Nevertheless, we pointed the telescope towards the mid-point of a segment going from the inner ring to the southern limb in order to have the maximum flux coming from the planet. This shift of the pointing on the planetary disk was taken into account when modeling the emission of the planet.

Analysis

The CO(1-0) and (2-1) lines are the weakest CO lines below 1000 GHz. Indeed, the tropopause temperature on Saturn and Uranus are 82 K and 53 K respectively, so that the population density of the energy levels peak for higher frequency lines. However, atmospheric opacity is lower at 115 GHz and 230 GHz than for higher frequencies. The best compromise at such low frequencies is reached at 230 GHz. This is the reason why the upper limits come from the 230 GHz observation analysis. As a consequence of the broad width of the CO lines, we do not exactly reach the continuum on the 230 GHz spectrum of Saturn. Hence, the low frequency edges of the synthetic spectra have been rescaled to match the edge of the 3-σ model.

None of the CO lines we targeted were detected on Saturn or Uranus. From the 230 GHz observations of Saturn analysis, we derive a CO mixing ratio 3-σ upper limit of 3.9 × 10 -8 for a constant mixing ratio of CO with altitude. This is an improvement by a factor of 3 as compared to the Rosenqvist et al. (1992) measurement. However, this limit is well above the 1 ppb level of the internal origin model of Noll & Larson (1991). If the bulk of CO is restricted to the stratosphere, then we derive an upper limit for the CO mixing ratio of 6.3 ×10 -8 . Once again this value is above, but only slightly, the Noll & Larson (1991) value in the stratosphere (2.5 × 10 -8 ). All the spectra discussed above are displayed on Fig. 6.

The observations of Uranus at 230 GHz (see Fig. 7) led to an upper limit of 1.8 × 10 -8 for a mixing ratio constant with altitude (internal origin) and 2.7×10 -8 if all the CO is in the stratosphere. The first value is consistent with Encrenaz et al. (2004) upper limit for an internal source and the second one is consistent with the upper limit derived by Marten et al. (1993). However, the latter seems to be in contradiction with Encrenaz et al. (2004) detection level in the stratosphere, at face value. But, the authors mention that their fluorescence model could overestimate the CO abundance by at least a factor of 2.

Discussion

Observations of CO in the atmospheres of the giant planets are of key importance in order to determine the origin of oxygenated compounds that are present in their atmospheres. CO was first detected in the atmosphere of Saturn in 1985-1986. From their work, [START_REF] Noll | [END_REF] could not distinguish between an internal (q CO = 1 ppb constant with altitude) and an external (q CO = 2.5 × 10 -8 restricted to the stratosphere) supply of atmospheric CO. From their 1990-2230 cm -1 high-resolution spectrum, Noll & Larson (1991) detected thirteen CO lines but were unable to firmly conclude on an internal origin, even if one of the lines favored the latter hypothesis. From photochemical modeling, [START_REF] Moses | The Collision of Comet Shoemaker-Levy 9 and Jupiter[END_REF] could also not distinguish between the origin scenarios, although both an internal and an external origin was favored as the simplest method for explaining all the CO, CO 2 and H 2 O observations. In their photochemical model of the atmosphere of Saturn, Ollivier et al. (2000) proposed a cometary-like composition for the oxygen external flux. This scenario cannot provide enough CO in the stratosphere of Saturn, so Ollivier et al. (2000) tentatively inferred that CO was mixed upward from the deep interior levels. Contrary to Ollivier et al. (2000), [START_REF] Moses | The Collision of Comet Shoemaker-Levy 9 and Jupiter[END_REF] could not reproduce the CO 2 and H 2 O data from ISO (Infrared space Observatory) with an influx of water alone and required some external influx of carbonoxygen bonded species. The 3-σ upper limit of 3.9 × 10 -8 , restricted to the stratosphere, yielded by our work still can not discard the external origin model from Noll & Larson (1991). Besides, it can also not discard the external source model from [START_REF] Moses | The Collision of Comet Shoemaker-Levy 9 and Jupiter[END_REF], where the CO mixing ratio increases with altitude and becomes greater than our upper limit between 10 -1 and 10 -5 mbar. Because the atmospheric layers above 10 -1 -10 -2 mbar do not contribute to the radiation at these wavelengths, the line is not sensitive to an increase of the CO mixing ratio in this part of the stratosphere. New observations with better sensitivity and/or more observing time would be needed at this frequency (230 GHz) to reach the 2.5 × 10 -8 level.

The weakness of the internal heat source in Uranus may be the cause of the limited convection in the deep interior of the planet (Hunten 1978;Wallace 1980), implying the lack of disequilibrium species, like CO, in its troposphere. Moreover, the quasi-isothermal structure of Uranus in the 1-mbar to 1-bar pressure range does not favor spectroscopic detections of minor constituents. The first attempts to detect CO at millimetric wavelengths led to upper limits on the mid-atmospheric (10-mbar) CO mixing ratio (4 × 10 -8 , Rosenqvist et al. 1992; 3 × 10 -8 , Marten et al. 1993). From CSO-FTS observations, Encrenaz et al. (1996) inferred an upper limit for the tropospheric content of CO (5 × 10 -7 ). The fluorescence of CO was detected in the atmosphere of Uranus by Encrenaz et al. (2004). Their atmospheric model suggests that the CO mixing ratio is greater in the stratosphere (3 × 10 -8 if all CO restricted to the stratosphere) than in the troposphere (upper limit: 2 × 10 -8 ). However, we did not detect the CO (2-1) line and a noise level analysis led to a 1.8 × 10 -8 upper limit if the CO mixing ratio is constant with altitude (internal origin) and 2.7 × 10 -8 , if restricted to the stratosphere (external origin). Our upper limit for the internal origin is consistent with the one derived by Encrenaz et al. (2004). However, our external origin upper limit is slightly lower than Encrenaz et al. (2004) detection level. But, their model did not take scattering by cloud particles into account. Therefore, the authors conclude that their values could be overestimated by at least a factor of 2. Thus, we suggest that Encrenaz et al. (2004) could have overestimated their measured value of the CO mixing ratio in the stratosphere (external origin model) of Uranus and that scattering has a non negligible effect on the spectrum of Uranus at 4.7-µm. From our observations, there is no evidence of an external origin of CO. This molecule could be a good candidate for the production of the observed CO 2 (Burgdorf et al. 2006). Even if CO is not detected from our observations, the level of our upper limits does not imply that we should reconsider the possibility of a production of CO 2 from CO, since CO 2 was detected at the (4.0±0.5)×10 -11 level. A complete photochemical model of Uranus, which takes chemi-cal uncertainties into account (Hébrard et al. 2007;Hébrard et al. 2006;Dobrijevic et al. 2003) should be developed in order to test several hypothesis for the supply of H 2 O, CO 2 and CO in order to understand their chemistry and to model the observed abundances of these species.

Conclusion

Although this attempt to detect CO in the atmospheres of Saturn and Uranus has failed, new upper limits have been derived. Our upper limit on a CO internal source in the atmosphere of Saturn (3.9 × 10 -8 , constant with altitude) is better than the previously published one (Rosenqvist et al. 1992) by a factor of 3. This upper limit on the internal source is far from the 1 ppb detection level of Noll & Larson (1991) and the upper limit we derived on the external source (6.3 × 10 -8 , restricted to the stratosphere) is greater than the stratospheric abundance of CO in the external origin model of Noll & Larson (1991) (2.5 × 10 -8 ) by a factor of 2.5. New observations with better sensitivity and/or more observing time should permit a derivation of the magnitude of the CO external source on Saturn. Our non-detection of CO at the 1.8×10 -8 level in the atmosphere of Uranus is consistent with the latest published upper limit on an internal source of CO (2×10 -8 , Encrenaz et al. 2004). Moreover, the upper limit we derived on the magnitude of an external source (2.7 × 10 -8 , restricted to the stratosphere) is below the detection level of Encrenaz et al. (2004). However, our measurement is not incompatible with the result of Encrenaz et al. (2004) (3 × 10 -8 , restricted to the stratosphere). Their fluorescence model did not take scattering effects by cloud particles into account. As a consequence, the mixing ratio they derived could have been overestimated. According to the authors, the factor of uncertainty on their result could be at least a factor of 2. Moreover, our new upper limits add new constraints on a possible temporal variability of the abundance of CO in the atmospheres of these planets.

From our observations, it is still impossible to conclude on the strength of the external flux of oxygen in the atmosphere of Saturn and Uranus. This question could be answered by observing at higher frequency where the lines are stronger. Such observations would still need several hours of integration because of the sky opacity, which is higher at such frequencies. Such observations would be focused on the stratospheric narrow emission peak in order to derive the stratospheric CO abundance but could also be performed over a wide frequency range (Lellouch et al. 2005;Hesman et al. 2007). CO observations with the Herschel Space Observatory would also be very valuable with regard to the foreseen outstanding capability to detect faint lines in short integration times of this telescope. Such observations should focus on the lines which have the strongest predicted line strength (good signal to noise ratio in a relatively short integration time) and for which observations are independent from the ammonia and phosphine abundances [START_REF] Cavalié | Herschel Open time Key Program Workshop[END_REF]). Frequency [GHz] Fig. 6. Observed 12 CO (2-1) spectra on Saturn. The total bandwidth is 4 GHz and the spectral resolution is 16 MHz. The spectra resulting from the upper limits we derived are plotted as follows: our 3-σ upper limit of 6.3 × 10 -8 in the stratosphere (external source of CO) in solid line and our 3-σ upper limit of 3.9 × 10 -8 constant with altitude (internal source of CO) in long-dashed lines. The internal source of CO model (short-dashed lines) as well as the external source of CO (dotted lines) taken from Noll & Larson (1991) are also plotted for comparison with our upper limits. The absorption feature at 231.280 GHz is the terrestrial O 3 line. ) spectra on Uranus. The total bandwidth is 4 GHz and the spectral resolution is 16 MHz. The short-dashed lines correspond to the upper limit determined by Encrenaz et al. (2004) for a CO mixing ratio constant with altitude (q CO = 2 × 10 -8 ). The external source of CO model (dotted lines) is the Encrenaz et al. (2004) best fit model (q CO = 3 × 10 -8 restricted to the stratosphere). Our 3-σ upper limits are plotted in solid line (q CO = 2.7 × 10 -8 , restricted to the stratosphere) and in long-dashed lines (q CO = 1.8 × 10 -8 , constant with altitude). The absorption feature at 231.280 GHz is the terrestrial O 3 line.

Fig. 1 .

 1 Fig. 1. Spectra of Saturn at (a) 115 GHz and (b) 230 GHz in the antenna temperature scale. Receiver A has a corrupt channel. It can be seen in the middle of the first three lowest frequency bands at 230 GHz. This channel is removed before reducing the spectra. The terrestrial ozone absorption band can be seen at 231.280 GHz. This feature is not taken into account when computing the spectrum noise level.

Fig. 2 .

 2 Fig. 2. Spectra of Uranus at (c) 115 GHz and (d) 230 GHz. Same caption as Fig. 1.

Fig. 3 .

 3 Fig. 3. Thermal profiles of Saturn (solid line) and Uranus (dashed line).

Fig. 4 .

 4 Fig. 4. Ammonia, phosphine and carbon monoxide mixing ratio vertical profiles on Saturn, all in solid lines. The internal and external source of CO models come from Noll & Larson (1991) (N&L'91 on the plot).The upper limits we have derived for the internal source (dotted lines) and for the external source (dashed lines) are also displayed.

Fig. 5 .

 5 Fig. 5. Ammonia and carbon monoxide mixing ratio vertical profiles on Uranus, all in solid lines. Both CO internal (upper limit) and external (best fit model) origin profiles come from Encrenaz et al. (2004) (E'04 on the plot). Our upper limits are plotted in dotted lines for the internal source and in dashed lines for the external source.

Fig. 7 .

 7 Fig. 7.Observed 12 CO (2-1) spectra on Uranus. The total bandwidth is 4 GHz and the spectral resolution is 16 MHz. The short-dashed lines correspond to the upper limit determined byEncrenaz et al. (2004) for a CO mixing ratio constant with altitude (q CO = 2 × 10 -8 ). The external source of CO model (dotted lines) is theEncrenaz et al. (2004) best fit model (q CO = 3 × 10 -8 restricted to the stratosphere). Our 3-σ upper limits are plotted in solid line (q CO = 2.7 × 10 -8 , restricted to the stratosphere) and in long-dashed lines (q CO = 1.8 × 10 -8 , constant with altitude). The absorption feature at 231.280 GHz is the terrestrial O 3 line.

Table 1 .

 1 Continuum

		115 GHz	230 GHz
	T * a	continuum	3-σ	continuum	3-σ
	Saturn	30.1 K	213 mK	37.5 K	420 mK
	Uranus	1.65 K	9 mK	2.9 K	33 mK

This feature is not taken into account when computing the spectrum noise level. and 3-σ noise level on the antenna temperature scale on the CO(1-0) and CO(2-1) spectral bands.

Table 2 .

 2 Collisional line width γ cm -1 .atm -1 (at 300 K) and temperature dependence factor n for NH 3 , PH 3 and CO with H 2 and He and for Saturn and Uranus (a blank space means that no data is available).

	γ	H 2	He	Saturn Uranus
	NH 3	0.069		0.069	0.069
	PH 3 (1-0)	0.1078 0.0618 0.1011 0.1003
	PH 3 (2-1)	0.1064 0.0606 0.0997 0.0990
	CO (1-0)	0.0749 0.0479 0.0709 0.0696
	CO (2-1)	0.0715 0.0468 0.0679 0.0666
	n	H 2	He	Saturn Uranus
	NH 3	0.67		0.67	0.67
	PH 3	0.73	0.30	0.67	0.66
	CO (1-0) & (2-1)			0.60	0.60
	References: Berge & Gulkis (1976) and Brown & Peterson (1994) for
	NH 3 , Lévy et al. (1993) and Lévy et al. (1994) for PH 3 and Draegert &
	Dudley (1968), Varanasi (1988), Beaky et al. (1996) and Weisstein &
	Serabyn (1996) for CO.
			Saturn	Uranus
	Rotation period	10.2h	17.9 h
	Equatorial radius	60 268 km 25 559 km
	Polar radius	54 364 km 24 973 km
	H 2 molar fraction	0.8606	0.8320
	He molar fraction	0.135	0.152
	CH 4 molar fraction	4.4 10 -3	1.6 10 -4

Table 3 .

 3 Saturn

and Uranus physical characteristics

(Lindal 1992;[START_REF] Conrath | [END_REF]Conrath & Gautier 2000;de Graauw et al. 1997;Baines et al. 1995)

.

Acknowledgements. IRAM is supported by INSU/CNRS (France), MPG (Germany) and IGN (Spain).


Breno O Fragomeni

Daniela A L Lourenco

Yutaka Masuda

Andres Legarra

Ignacy Misztal

Erratum to: Incorporation of causative quantitative trait nucleotides in single-step GBLUP

Author details [START_REF] Fragomeni | Incorporation of causative quantitative trait nucleotides in single-step GBLUP[END_REF]

Open Access

Ge n e t i c s Selection Evolution *Correspondence: fragomen@uga.edu [START_REF] Fragomeni | Incorporation of causative quantitative trait nucleotides in single-step GBLUP[END_REF] Edgar L. Rhodes Center for Animal and Dairy Science, University of Georgia, Athens, GA, USA Full list of author information is available at the end of the article The online version of the original article can be found under doi:10.1186/s12711-017-0335-0.

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T Cavalié

F Billebaud

N Biver

M Dobrijevic

E Lellouch

J Brillet

A Lecacheux

Å Hjalmarson

Aa Sandqvist

U Frisk

M Olberg

The Odin Team

E A Bergin

Observation of water vapor in the stratosphere of Jupiter with the Odin Space Telescope

Keywords: Jupiter, atmosphere, water, spectroscopy, Odin space telescope

Observation of water vapor in the stratosphere of

Jupiter with the Odin Space Telescope.

Introduction

The Infrared Space Observatory has detected water vapor in the stratospheres of the giant planets (Feuchtgruber et al. 1997;[START_REF] Feuchtgruber | The Universe as Seen by ISO[END_REF][START_REF] Lellouch | The Universe as Seen by ISO[END_REF]. The large amount of water measured above the condensation level of vapor (cold trap at the tropopause) implies the presence of an external source of oxygenated compounds (Moses et al. 2000b; Lellouch et al. 2002). These compounds could be brought by interplanetary dust particles (IDP), sputtering from the rings and/or satellites and large cometary impacts. Observations of Jupiter carried out by ISO tend to prove that most of the stratospheric water is due to the Shoemaker-Levy 9 (SL9) comet impacts in July 1994 (Lellouch et al. 2002), whereas [START_REF] Bergin | [END_REF] obtained satisfactory fits to the Submillimeter Wavelength Astronomy Satellite (SWAS) data by considering IDP infall, with a constant flux of 2.0×10 6 cm -2 .s -1 .

The submillimeter satellite Odin was launched in 2001 and obtained a high resolution spectrum of Jupiter's water vapor line (1 10 -1 01 ) at 557 GHz on November 8, 2002. This spectrum is presented in this work as well as a reanalysis of SWAS observations. Spectral analysis combined with the use of our photochemical model (Ollivier et al. 2000, adapted to Jupiter) provides new clues which help understanding the origin of water vapor in the stratosphere * Tel: +33-5-5777-6123; fax: +33-5-5777-6110

Email address: Thibault.Cavalie@obs.u-bordeaux1.fr (T. Cavalié).

A description of the observations is given in Sect. 2. Our photochemical and radiative transfer models are described in Sect. 3. Our results are presented in Sect. 4 and the different sources of H 2 O are discussed in Sect. 5.

Observations

The space telescopes SWAS and Odin observed the water vapor 557 GHz line on Jupiter in 1999, 2001 (SWAS) and 2002 (Odin). The resulting brightness temperature spectra have a signal-to-noise ratios of ∼17, 10 and 16 respectively for the 1999, 2001 and 2002 observations. The spectral resolution is about 1 km.s -1 for the SWAS spectra and 0.6 km.s -1 for the Odin spectrum.

The SWAS spectra are corrected for the Double Side Band (DSB) response of the instrument. Nevertheless, the SWAS spectra show broad features at 100 km.s -1 and more, which cannot be reproduced in models. These wings, probably due to instrumental effects as mentionned in [START_REF] Bergin | [END_REF] and Lellouch et al. (2002), cause an uncertainty on the continuum level of the emission. More details on the SWAS 1999 and 2001 observations can be found in [START_REF] Bergin | [END_REF] and Lellouch et al. (2002).

The Odin observations were carried out with the Acousto-Optical Spectrometer (AOS) in a classical position switching mode (Olberg et al. 2003). The receivers are operated in a Single Side Band (SSB) mode. The spectral band is 1 GHz. As Jupiter has a strong continuum emission at this frequency, stationary waves are generated within the instrument, causing ripples on the spectrum (Fig. 1). The subtraction of the ripples is the source of an uncertainty of 10% The signal-to-noise ratio is 16.

on the line contrast and some uncertainty on the line wing shape.

As the beam size (3.3'×4.5' for SWAS and 2.1' for Odin) is larger than the planet size (∼35-40"), all the observed features correspond to the emission of the whole planet. The line width is dominated by the smearing effect because of limb equatorial velocity ∼12.6 km.s -1 of the planet [START_REF] Bergin | [END_REF]. As no absolute calibration has been done for the Odin observations, all results are discussed in terms of line-to-continuum ratios and the Odin/SWAS observed continuum have been rescaled to the brightness temperature scale of our model (T B =128.6 K).

Modeling

We describe, in this section, details of our data analysis procedure that can be summarized in the following way:

• A water vertical profile is simulated from a time-dependent 1D photochem-ical model. The main parameters that affect this profile are the altitude and the magnitude of water deposition (in the case of a sporadic cometary origin), the magnitude of the water influx (in case of a steady interplanetary dust particle flux) and the eddy diffusion coefficient in the stratosphere.

• A radiative-transfer model computes a synthetic spectrum for each water vapor profile.

• Comparison of observational data and synthetic spectra enables to constrain parameters of the photochemical model.

Photochemical modeling

We used a time-dependent photochemical model, derived from the model developed for Saturn by Ollivier et al. (2000) and which has been adapted to the case of the atmosphere of Jupiter. For each altitude and each chemical compound i, the code solves the continuity equation

dn i dt = P i -n i L i -div (φ i ) ( 1
)
where n is the concentration, P the chemical production, L the chemical loss and φ the vertical flux. This is a one-dimensional model since only the vertical transport is considered. ). The sporadic input of water due to the impacts was modeled via two parameters: the deposition pressure p 0 and the initial mixing ratio q 0 above the p 0 level (see Lellouch et al. 2002 for more details). For each computation, the value of q 0 was set to a constant value as a function of altitude (above the p 0 level).

Thus, we have two possibilities for the SL9 models. The first one consists of fixing the value of p 0 and adjusting the value of q 0 with the data. In the second case, we fix q 0 and adjust p 0 . Some constraints exist on both p 0 and q 0 . The most reliable constraint is probably the fact that the deposition level that was

Radiative transfer modeling

We modeled the observed submillimeter radiation with a line-by-line nonscattering radiative transfer model. We computed synthetic spectra of the H 2 O 557 GHz line. The program represents the approximate spherical geometry of the planet so that planetary disk and limb contributions are taken into account. We assumed an uniform distribution of all other opacity sources and we adopted a mean thermal profile (see Fig. 3) of the atmosphere of Jupiter (Fouchet et al. 2000a) since our beam size is larger than the observed planetary disk. Continuum opacity is dominated by H 2 -He-CH 4 collision-induced absorption (Borysow et al. 1985(Borysow et al. , 1986(Borysow et al. and 1988)). Following Moreno (1998), the opacity due to the far wings of ammonia and phosphine lines was also included. We used the Fouchet et al. (2000b) ammonia and phosphine mixing ratio vertical profiles (see Fig. 4). Spectroscopic parameters for NH The rapid rotation of Jupiter (9.9 h) induces the smearing of the disk-averaged line on the spectrum, because of the Doppler shifts due to the gas rotation velocity (12.6 km.s -1 at the eastern and western limbs). The way this effect is taken into account is described in [START_REF] Bergin | [END_REF].

We briefly come back to the use of disk-averaged vertical profiles of mixing ratio. The fact that the SL9 impacts were all located in the southern hemi- latitude of 44 • S. Longitudinal mixing proved to be efficient in the submillibar region. Indeed, HCN was observed at such pressure levels a few months after the comet impacts and the maps showed that it had already spread over several degrees in longitude (Bézard et al. 1997). So, the deposits quickly formed a longitudinal belt after the impacts. Thus we have to take into account the background amount of water present in the stratosphere of Jupiter, which is

due to the low IDP flux (Φ IDP H 2 O =4×10 4 cm -2 .s -1
), and the SL9 input located at 44 • S, which is modeled via the parameters p 0 and q 0 . By averaging those two kinds of vertical profiles over the surface of the planet, we obtain the kind of profile shown in Fig. 2 (see "hybrid model" in Lellouch et al. 2002),

where p 0 is determined by the SL9 input and where q 0 is multiplied by the ratio between the surface of the SL9 longitudinal belt and the total surface of the planet. Using this approach, the values of q 0 we derive from the observations are disk-averaged values. A disk-averaged water vertical profile is adapted since the beam size is greater than the planet size.

Results

The best-fit models have been determined with a χ 2 minimization process.

All profiles and column density values are disk-averaged. One must note that an uncertainty of 5 K on the thermal profile would add an uncertainty of 0.4 × 10 15 cm -2 on the water vapor column density, 0.3 × 10 -7 on q 0 (in the case of a SL9 origin) and 0.6×10 6 cm -2 .s -1 (in the case of an IDP origin).

SWAS data

The of our model (T B =128.6 K), then the fits of IDP models are far better (Fig. 7). The flux we derive is Φ IDP H 2 O =(3.7±0.5)×10 6 cm -2 .s -1 and the corresponding column density is (2.8±0.4)×10 15 cm -2 . The synthetic spectrum is within The SL9 model where q 0 is fixed to 6×10 -8 also gives a good fit for p 0 =(0.45± 0.09) mbar (see Fig. 8). Here, the error bar on the p 0 value is not due to the 1-σ level of the spectrum. Indeed, the synthetic spectra with either p 0 =0.37 mbar or p 0 =0.54 mbar are outside the 1-σ level of the spectrum. This error bar is due to the fact that the integration step of the photochemical model is 5 km. This results in 0.09 mbar steps in the 0.2-0.6 mbar region. Taking p 0 =0.45 mbar and q 0 =6×10 -8 , the column density of water is (2.6±0.6)×10 15 cm -2 at the time of the impacts.

Odin data

After removing the ripple pattern, the line shows some asymmetry in the line wings. This, as well as the noise level, is a limitation in the determination of the best-fit model. Testing the IDP fluxes leads us to retrieve of a lower flux than the flux retrieved from the SWAS data. Indeed, the χ 2 minimum is obtained for a flux value of Φ IDP H 2 O =(3.4±0.5)×10 6 cm -2 .s -1 (see Fig. 9). This result is compatible with the SWAS initial results (before shifting downward the Rayleigh-Jeans temperature continuum). If we try to fit the line with an averaged best-fit model to the SWAS/Odin data (Φ IDP H 2 O =3.6×10 6 cm -2 .s -1 ), then the line center is better reproduced (see Fig. 9). Nevertheless, such a modeling results in broader wings, but they still are within the 1-σ error bars.

As for the SL9 model, restraining the bulk of water above an initial pressure level of 0.2 mbar, results in narrower lines than the IDP model. The line center as well as the wings are well reproduced with the synthetic spectra.

When fixing p 0 to 0.2 mbar, the optimum water mixing ratio above this level is q 0 =2.0×10 -7 . The uncertainty is 0.5×10 -7 . When fixing q 0 to 6.0×10 -8 (Lellouch et al. 2002), p 0 is found to be (0.54±0.09) mbar (see Fig. 10). The latter model implies a column density of (3.2±0.6)×10 15 cm -2 at the time of the comet impacts.

Discussion

The best-fit model parameters for each observation, as derived from χ 2 minimization, are summarized in Table 2. From this set of parameters, we derived averaged values. For each model (IDP, SL9 with q 0 fixed and SL9 with p 0 fixed), the value obtained is affected a weight related to the signal-to-noise ratio of the observation. Doing this way, we obtained the averaged values used in Fig. 11.

First of all, when considering the SL9 source hypothesis and fixing the value of q 0 at 6×10 -8 , we derive a deposition pressure level p 0 in the range of 0.45-0.54 mbar. The column density we derived is consistent with the value of Lellouch et al. (2002). However, even if our model does not provide a more precise value of p 0 , the range of the values we derive is outside the ranges derived by Lellouch et al. (1997) and Moreno (1998) from CO observations at millimeter wavelengths at the time of the impacts, which are 0.04-0.2 mbar and 0.2±0.1 mbar (respectively). Therefore, we regard this possibility as unlikely with regard to both SWAS and Odin data.

So, the models we have to compare are the IDP model and the SL9 model with p 0 =0.2 mbar. We derived an external flux of water, originating from an IDP source, of Φ IDP H 2 O =(3.6±0.5)×10 6 cm -2 .s -1 . This value is greater than the one derived by [START_REF] Bergin | [END_REF] by a factor of less than 2. From their physical model, which only included vertical transport (no chemical or photochemical processes), the authors derived a deposition flux of 2.0×10 6 cm -2 .s -1 . To obtain a narrower line from their model and thus to obtain their best-fit model, they increased the mixing ratio over pressure slope (-d(log q)/d(log p)) of their physical profile from 0.8 to 1.3. Nevertheless, as noted by the authors, changing the slope could not simulate precisely the effects of photolysis, chemical reactions and the non-linearity of the interactions between these processes as well as vertical transport and condensation. Taking photolysis and chemical losses into account, they would probably have obtained a higher value for the flux consistent with our result.

With a SL9 model, we obtain q 0 =(1.9±0.5)×10 -7 when fixing p 0 =0.2 mbar.

Model Φ IDP H 2 O [cm -2 .s -1 ] p 0 [mbar] q 0 IDP (3.7±0.5)×10 6 - - SL9 (q 0 fixed) 4.0×10 4 (0.45±0.09) 6×10 -8 SL9 (p 0 fixed) 4.0×10 4 0.2 (1.8±0.5)×10 -7
Odin 2002

Model Φ IDP H 2 O [cm -2 .s -1 ] p 0 [mbar] q 0 IDP (3.4±0.5)×10 6 - - SL9 (q 0 fixed) 4.0×10 4 (0.54±0.09) 6×10 -8 SL9 (p 0 fixed) 4.0×10 4 0.2 (2.0±0.5)×10 -7
Table 2 Best-fit model parameters for each set of data and each model, from which the averaged best-fit value are derived (see text).

line wings than the IDP model (see Fig. 11). A χ 2 analysis clearly indicates that the SL9 model gives better fits to the data. However, all the IDP synthetic spectra are within the 1-σ error bars on all observations. So, this model cannot be ruled out at this stage.

If the observed water would come from the SL9 comet, then the non-steady state created by the deposition of the cometary material above the p 0 level in our model should evolve towards a steady state where the only observable source of water would be the low IDP flux (4×10 4 cm -2 .s -1 in our model).

From our computations, such a state is reached ∼400 years after the im-pacts. As a result, the downward diffusion of water as well as the photochemical/chemical losses effects would first desaturate the line. Thus, the line contrast should first increase with time (see Fig. 12). Our photochemical model predicts that the line center temperature of the line should increase by 0.76 K from 1999 to 2007. Taking the noise level of the SWAS 1999 observations into account, our model predicts that this effect could only be observed in 2007 by reaching a signal-to-noise ratio of 50 with the Odin telescope. Afterwards, the amount of water decreasing more and more with time at submillibar pressures, the line should become fainter and broader and should tend towards the line that would be due to the low IDP flux only (see Fig. 12). This change should be observable with Herschel-HIFI.

One must not forget that the shape of the water vertical profile computed with a photochemical model highly depends on the vertical eddy diffusion coefficient profile K(z). Due to strong uncertainties in the chemical scheme, profile is a conservative way of analysing the observed lines with regard to the implications noted above.

Conclusion

In this paper, we have shown that the high signal-to-noise ratio observations of water vapor in the stratosphere of Jupiter, carried out with SWAS and the Odin telescope between 1999 and 2002, favor a SL9 origin for water. Indeed, all observations are better fitted when the bulk of water is restricted to submillibar pressures. In our disk-averaged and simplified deposition model of the SL9 water, we derived a water mixing ratio of 1.9×10 -7 above an initial pressure deposition level of 0.2 mbar. In this model, a low IDP flux of 4×10 -100 -50 0 50 100 Brightness Temperature [K] Velocity [km.s -1 ] Fig. 10. Odin 2002 data modeled with SL9 models. When p 0 is fixed to 0.2 mbar, the derived q value is 2.0×10 -7 (solid line) whereas when fixing q 0 to 6.0×10 -8 , the derived p 0 pressure level is 0.54 mbar (long-dashed lines).

Fig. 1 .

 1 Fig. 1. Odin observations of Jupiter at the H 2 O (1 10 -1 01 ) line frequency on November 8, 2002. The observed antenna temperature is displayed as a function of velocity.

Fig. 2 .

 2 Fig. 2. Example of SL9 source vertical profiles of water at the time of the SL9 impacts (07/1994) in solid line and at the time of the Odin observations (11/2002) in dashed lines. The evolution of water abundance is computed by the photochemical model. The water vapor mixing ratio is displayed as a function of atmospheric pressure. Profiles correspond to a fixed value of p 0 =0.2 mbar, and an adjusted value of q 0 =1.9×10 -7 .

Fig. 3 .

 3 Fig. 3. Disk-averaged thermal profile of the atmosphere of Jupiter. The tropopause temperature is 109 K. The profile is isothermal (T =165.4 K) between 1 mbar and 10 -3 mbar. Reference: Fouchet et al. (2000a).

Fig. 4 .

 4 Fig. 4. Ammonia (solid line) and phosphine (long-dashed lines) mixing ratio vertical profiles as a function of pressure in Fouchet et al. (2000b).

Fig. 5 .Fig. 6 .

 56 Fig. 5. Best-fit model to the (a) SWAS 1999 and (b) 2001 data obtained with a SL9 model with the initial parameters p 0 =0.2 mbar and q 0 =1.8×10 -7 (solid lines). The IDP models (long-dashed lines) correspond to infall fluxes of Φ IDP H 2 O =3.4×10 6 cm -2 .s -1 for 1999 and 2001 respectively.

the 1 -

 1 σ error bars on the data over the [-80:+20] km.s -1 range. Finally, the IDP model cannot be ruled out at this stage, because of the uncertainty on the continuum level of each observation, even if a χ 2 analysis shows that the SL9 model gives a better match to the data than the IDP model. All the IDP models that are considered for the SWAS data in what follows are models with downward shifted continuum (to 125.4 K and 117.0 K, for 1999 and 2001 respectively).

  Lellouch et al. (2002) derived a column density of (2.0±0.5)×10 15 cm -2 at the time of the ISO observations. The column density we derived is (3.7±1.0)×10 15 cm -2 at the time of the impacts. This value is greater than the ISO value, but by taking photolysis, chemical reactions, vertical transport and condensation, this value decreases down to (3.1±0.8)×10 15 cm -2 at the time of ISO observations. This value is still above the Lellouch et al. (2002) value, but there is a small overlap on the ranges of values. Moreover, considering an uncertainty of 5 K on the thermal profile ends up in an additional uncertainty of 0.4 × 10 15 cm -2 on the column abundance. So, these values could well be consistent and an intermediate value of column density should be compatible with all inferred values. As the water vapor vertical profile of Lellouch et al. (2002) was computed from a vertical transport model, the ISO data should be re-analysed with a more complete photochemical model. This work still has to be done and its results could be directly comparable to ours. The SL9 model quoted above better reproduces the line contrast as well as the SWAS 1999 and 2001

4 cm - 2 .s - 1 .

 421 Fig. 7 of Moses et al. (2005), many different K(z) profiles have been inferred from past observations. At the submillibar pressure range, K(z) 5 × 10 4 cm -2 .s -1 within a factor of 2 (Moreno et al. 2003). According to the Moses et al. (2005) model C value used in this work, K(z) is equal to 7.8 × 10 4cm -2 .s -1 . At pressures between 0.1 mbar and 100 mbar (tropopause level),Gladstone et al. (1996) found values of K(z) higher by a factor of ∼3. So we have to consider the fact that another choice in the K(z) profile could change our results. In the lower stratosphere, our adopted K(z) profile gives

Fig. 7 .Fig. 8 .

 78 Fig. 7. Brightness temperature spectra as observed by SWAS (a) in 1999 and (b) in 2001. Both spectrum continuum have been rescaled so as to obtain a better fit of the line wings with an IDP model. (c): water mixing ratio vertical profile as a function of pressure resulting from the observed flux of Φ IDP H 2 O =3.7×10 6 cm -2 .s -1 .

Fig. 9 .

 9 Fig. 9. Odin 2002 data modeled with IDP models. The solid line corresponds to a flux of Φ IDP H 2 O =3.4×106 cm -2 .s -1 (χ 2 minimum value). The long-dashed lines correspond to the overall (SWAS and Odin data) best-fit model (Φ IDP H 2 O =3.6×10 6 cm -2 .s -1 ).

Fig. 11 .Fig. 12 .

 1112 Fig. 11. Overall best-fit models for the SWAS 1999 and 2001 and Odin 2002 observations. Solid lines: SL9 model with p 0 =0.2 mbar (fixed) and q 0 = 1.9×10 -7 ; longdashed lines: SL9 model with p 0 =0.45 mbar and q 0 =6×10 -8 (fixed); short-dashed lines: IDP model with a steady infall flux of water Φ IDP H 2 O =3.6×10 6 cm -2 .s -1 . The overall best-fit parameter have been obtained from Table 2 and by taking the signal-to-noise ratio of each observation into account. Doing this way, the SWAS 1999 observations have a lower impact on the results than the SWAS 2001 and Odin 2002 observations.

Table 1 Collisional

 1

	γ	H 2	He	Jupiter
	NH 3	0.069		0.069
	PH 3 (2-1) 0.1064 0.0606 0.1001
	H 2 O	0.0811 0.0228 0.0731
	n	H 2	He	Jupiter
	NH 3	0.67		0.67
	PH 3	0.73	0.30	0.67
	H 2 O	0.9	0.50	0.85

139

sphere is not a limitation to our hypothesis. All the impacts occured at the line width γ cm -1 .atm -1 (at 300 K) and temperature dependance factor n for NH 3 , PH 3 and H 2 O with H 2 and He and for Jupiter (a blank space means that no data are available). References: Berge & Gulkis (1976) and Brown & Peterson (1994) for NH 3 , Levy et al. (1993,1994) for PH 3 and Dutta et al. (1993) for H 2 O.

  4 cm -2 .s -1 was also taken into account. This suggests a localised input of water, in terms of altitude, which is contradictory with a steady state resulting from an IDP permanent flux. Nevertheless, all synthetic spectra obtained from anIDP flux of Φ IDP H 2 O =(3.6±0.5)×106 cm -2 .s -1 give fits that are within the 1-σ error bars of the observations, but the χ 2 value is greater than the one computed from the SL9 model. In view of these results, the ISO data of 1997 should be re-analysed using the model developped in this work. Further observations, reaching a higher signal-to-noise ratio are needed to state on the origin of water vapor in the stratosphere of Jupiter, even if the SL9 origin is favored by both SWAS and Odin observations. The analysis of the latest Odin observations (August 2007) is underway. Moreover, Herschel observations with the HIFI instrument (500 GHz -2000 GHz) should allow obtaining a signal-to-noise ratio with a comparable spectral resolution in reasonable times. Such a high signal-to-noise ratio would enable us to better resolve the line wing shape in order to discriminate between both origins. Moreover, a temporal variability of the line could be brought to light. Such a variability should not be expected with an IDP origin. Indeed, Moses et al. (2000a) suggested that the production of the IDP is dominated by short-period comets. Selsis et al. (2004) showed that 48 short-period (∼5-10 year periods) comets approach Jupiter's orbit at less than the Roche lobe radius of the planet. So, the IDP flux on Jupiter should be steady. Finally, using HIFI at the highest frequencies would result in a sufficient spatial resolution to carry out maps of

	Brightness Temperature [K]
	-100	-50			0		50		100
	(a)			Velocity [km.s -1 ]
	Brightness Temperature [K]
	-100	-50			0		50		100
	(b)			Velocity [km.s -1 ]
	0.001
	Jupiter at water vapor frequencies. A latitudinal inhomogeneous distribution 0.01
	of water, with an increase of its amount in the southern hemisphere would be 0.1
	a strong signature of a SL9 impact origin and could provide information on the horizontal diffusion at the submillibar level. 1 Pressure [mbar]
	10
	100
	1000
	1e-14	1e-13	1e-12	1e-11	1e-10	1e-09	1e-08	1e-07	1e-06
	(c)				Mixing Ratio


domain: shs sdv info

text:
Mylène Delosière

Agnès Thomas

Claudia Terlouw

Denis Durand
email: denis.durand@inra.fr

Plasma indicators of bovine health: impacts of diet supplementations and pre-slaughter stress

Keywords: Cows, plasma, health, diet, stress Comment citer ce

This data article reports the values of indicators of bovine health determined in the plasma of Normand cull-cows at different times of the about 100 days lasting finishing period and at slaughter. The data constitute a large dataset based on the quantification of metabolites and the evaluation of enzymes activities allowing the determination of antioxidant capacity, oxidative stress level, energy and lipid metabolisms, activity of the Hypothalamo-Pituitary-Adrenal axis and the hepatic status in cull-cows.

Specifications

analyze relationships between health indicators to increase our understanding of the metabolism of cows; iii) compared with other data from plasma or tissues of ruminants in order to understand effects of breeds, species, diet supplementation, or preslaughter-stress on health indicators. iv) be computationally aggregated with health indicators of other cows.

Data

The dataset reports the concentration of metabolites determined in plasma of cows in order to evaluate the health's parameters of animals (Table 1) during the finishing period before slaughter. The factors of variation were diet supplementations and pre-slaughter stress. Plasma sampling took place throughout the about 100 days finishing period: the day before experimentation, on days 50 and 100 and just after stunning, before bleeding. Details of the sampling procedures can be found in [START_REF] Bourguet | Characterising the emotional reactivity of cows to understand and predict their stress reactions to the slaughter procedure[END_REF] [START_REF] Bourguet | Characterising the emotional reactivity of cows to understand and predict their stress reactions to the slaughter procedure[END_REF] and [START_REF] Gobert | Prévention de la peroxydation des acides gras polyinsaturés par des antioxydants alimentaires chez le bovin en production[END_REF] [START_REF] Gobert | Prévention de la peroxydation des acides gras polyinsaturés par des antioxydants alimentaires chez le bovin en production[END_REF].

Indicators of oxidative stress level

In vivo

Gaz chromatography -mass spectrometry

Watters quattro micro [START_REF] Lesgards | Differential distribution of 4-hydroxynonenal adducts to sulfur and nitrogen residues in blood proteins as revealed using Raney nickel and gas chromatography-mass spectrometry[END_REF] adapted by [START_REF] Gobert | Pre-slaughter stress and lipoperoxidation: protective effect of vitamin E and plant extracts rich in polyphenols given to finishing cattle[END_REF] Indicators of energy and lipid metabolism

Dietary treatments.

All cows were given a straw (30%) and concentrate (70%)-based diet during 101 ± 3 days (Figure 1). Nine cows received no supplementation (Control group, C). For 64 of the cows the diet was supplemented with lipids (40g oil/kg diet DM, L group) provided by extruded oilseeds (Table 2). (LEP group). The PERP were prepared from rosemary (Rosemarinus officinalis), grape (Vinis vitifera), citrus (Citrus paradisi) and marigold (Calendula officinalis) by the Phytosynthèse company (Riom, France) (INRA patent #P170-B-23.495 FR). Cows were housed in 6 x 6 m pens with straw bedding combining 2 experimental and 2 non experimental animals according to a balanced design relative to feeding treatments. Pens were equipped with electronic feeding gates and individually offered the appropriate allowance of concentrates and straw per day for each cow. This periodically adjusted allowance supplied the required amount of nutrient to achieve a target growth rate of 1150 g/d. The finishing period of 101 d. in our study was slightly longer than the 70 d. French standards for cattle and was chosen to facilitate experimental organization and to achieve good production conditions in such cull-cows. Pre-slaughter treatments.

At the end of the finishing period, two slaughter conditions were used: limited stress (stress -) vs moderate stress (stress +), which are described in detail in Bourguet et al. (2010) [START_REF] Bourguet | Characterising the emotional reactivity of cows to understand and predict their stress reactions to the slaughter procedure[END_REF]. Of each of the L and LEP dietary groups, 7 L and 8 LEP cows were slaughtered under the moderate stress conditions, respectively, referred to as the LS and LEPS cows. All other cows were slaughtered under limited stress conditions. Briefly, each cow slaughtered under limited stress conditions was directly transported (6.9 ± 0. 8 communicating corridors containing many sharp bends (Figure 2). The labyrinth was partly indoors (138 m) and outdoors (122 m). Cows were unloaded at the entrance of the indoor part of the labyrinth and taken through it 3 times, by 2 purposely noisy experimenters (shouting and hitting metal structures with a stick) over a period of 28 min. They were subsequently individually transported (14.4 ± 0.8 min) to the experimental abattoir where they were immediately unloaded. Upon arrival at the abattoir, all cows were immediately stunned by a captive bolt and bled.

Figure 2: Schematic representation of the labyrinth used for the slaughter protocol with "moderate stress" conditions. The labyrinth was partly indoors and partly outdoors and included corridors with sharp bends and two cattle crushes in which the animal was repeatedly restrained. See [START_REF] Bourguet | Characterising the emotional reactivity of cows to understand and predict their stress reactions to the slaughter procedure[END_REF] for details.

Cows were slaughtered with one slaughter day per week over a period of 8 weeks. Slaughter took place between 07h30 and 11h10 am. Each slaughter day, two experimental cows from a same pen were slaughtered using each slaughter condition.

Sample collection.

Blood samples were collected in the morning, before food distribution day 0 (D0) and after the beginning of food distribution day 50 (D50) (not all cows) and day 100 (D100). The collection by venepuncture used tubes containing anticoagulants, such as trisodique ethylene diamine tetraacetique (EDTA K3), lithium heparin or trisodique citrate and were placed on ice until centrifugation. Blood samples were collected immediately after stunning (before bleeding) (Dslaugh) by venepuncture syringe and blood was distributed in tubes containing EDTA K3 or lithium heparin and placed on ice until centrifugation. The tubes were centrifuged at 4500 rpm (corresponding to 1950 g) during 15 min (+4°C), and plasma was transferred to Eppendorf tubes which were stored at -80°C. Some tubes were acidified using HClO 4 to achieve a deproteinization needed for β-hydroxybutyrate and lactate determinations. Some plasma tubes presented hemolysis and discarded explaining some missing values in the dataset.

Biochemical approaches were applied on the plasma in order to appreciate the global health of cows evaluating indicators of antioxidant capacity, oxidative stress level, energy and lipid metabolisms, Hypothalamo-Pituitary-Adrenal activity and hepatic status (detailed in Table 1).

  Comment citer ce document : Delosière, M., Thomas, A., Terlouw, C., Durand, D. (2018). Plasma indicators of bovine health: impacts of diet supplementations and pre-slaughter stress. Data in Brief, 21, 1-11. , DOI : 10.1016/j.dib.2018.10.009 3

  Comment citer ce document : Delosière, M., Thomas, A., Terlouw, C., Durand, D. (2018). Plasma indicators of bovine health: impacts of diet supplementations and pre-slaughter stress. Data in Brief, 21, 1-11. , DOI : 10.1016/j.dib.2018.10.009

5 Indicators 2 .

 52 Comment citer ce document : Delosière, M., Thomas, A., Terlouw, C., Durand, D. (2018). Plasma indicators of bovine health: impacts of diet supplementations and pre-slaughter stress. Data in Brief, 21, 1-11. , DOI : 10.1016/j.dib.2018.10.009 Comment citer ce document : Delosière, M., Thomas, A., Terlouw, C., Durand, D. (2018). Plasma indicators of bovine health: impacts of diet supplementations and pre-slaughter stress. Data in Brief, 21, 1-11. , DOI : 10.1016/j.dib.2018.10.009 6 Experimental Design, Materials and Methods Animals. Seventy-five, 48-60 months old, Normand cull cows (mean live weight 642 kg) were introduced into the experiment organized in three repetitions, conducted in winter 2007 (n=24), summer and autumn 2007 (n=25) and winter 2008 (n=26). Cows were identified by numbers between 251 and 331. Two cows presenting health problems during the finishing period (sarcosporidiose and abomasum flipping) were removed from the experimentation leaving 73 cows.

Figure 1 :

 1 Figure 1: Experimentation designed to study the effects of 3 diet supplementations (lipids, lipids and vitamin E, lipids, vitamin E and PERP) and 2 slaughter protocols relative to pre-slaughter stress conditions on plasma of Normand cull-cows over a 100 days finishing period and at slaughter.

  For the cows supplemented with lipids, 17 cows received a diet supplemented with vitamin E (155 IU/kg) (LE group) and 17 cows with vitamin E (155 UI/kg) and plant extracts rich in polyphenols (PERP; 7 g/kg diet DM) Comment citer ce document : Delosière, M., Thomas, A., Terlouw, C., Durand, D. (2018). Plasma indicators of bovine health: impacts of diet supplementations and pre-slaughter stress. Data in Brief, 21, 1-11. , DOI : 10.1016/j.dib.2018.10.009 7

  6 min) in a lorry (3 x 2 m) from the experimental farm towards the experimental abattoir of the INRA research center. During the journey of 2 km, each cow was accompanied by a non-experimental conspecific to avoid social isolation stress and were handled calmly. The objective of the moderate stress treatment was to combine in a standardized manner stress factors of psychological (novelty, social isolation, presence of active humans, noise) and physical origins. Each cow was individually transported (11.3 ± 0.3 min) towards an unfamiliar farm where a labyrinth was constructed with Comment citer ce document : Delosière, M., Thomas, A., Terlouw, C., Durand, D. (2018). Plasma indicators of bovine health: impacts of diet supplementations and pre-slaughter stress. Data in Brief, 21, 1-11. , DOI : 10.1016/j.dib.2018.10.009

  Comment citer ce document : Delosière, M., Thomas, A., Terlouw, C., Durand, D. (2018). Plasma indicators of bovine health: impacts of diet supplementations and pre-slaughter stress. Data in Brief, 21, 1-11. , DOI : 10.1016/j.dib.2018.10.009 9

Table Subject

 Subject

	area	Biochemistry, Biology
	More specific	Animal health, oxidative stress, lipid oxidation
	subject area
	Type of data	Tables, figures
	How data was	Spectrophotometry, spectrofluorometry, enzymatique… (precisions in table 1)
	acquired
	Data format	Analyzed and ready to use
	Experimental	Indicators were determined on total plasma or protein-depleted plasma
	factors
	Experimental	Registration of plasma indicators of bovine health during a finishing period
	features	studying diets and pre-slaughter stress treatments on animals
	Data source	INRA, Theix, St-Genès-Champanelle, France
	location
	Data accessibility Dataset is available in public repository: Portail Data INRA (data.inra.fr)
		Data identification number: 10.15454/UUASR4
		https://doi.org/10.15454/UUASR4
		During the reviewing process by Data In Brief, please find data in this private
		URL:
		https://data.inra.fr/privateurl.xhtml?token=a82e4fab-654b-4d7b-9848-
		4a2d11bfbe49
	Dataset citation

[1] Delosière M., Thomas, A., Terlouw, C., Durand, D., (2018) "Dataset : Plasma indicators of bovine health", https://doi.org/10.15454/UUASR4, Portail Data Inra, It is the first time a large dataset combines various health indicators in Normand cull cows. This dataset can be used to: i) assess of values of health indicators in cows; ii)

Table 1 :

 1 Plasma indicators of bovine health.

		Abbreviat
	Indicators	ions
		(Unit)

Biological meanings Meth ods Technologies Equipments Reference s Indicators of antioxidant capacity

	Kinetics of
	conjugated dienes (CD) generation 1 -Length of lag	Lp (min)	Resistance time of PUFA against peroxidation	Ex vivo	Spectrophotom etry	Uvikon XS	[5]
	phase
	Kinetics CD		Speed of the
	generation 2 -Maximal rate	Rmax (A 234 /min)	propagation reaction of PUFA	Ex vivo	Spectrophotom etry	Uvikon XS	[5]
	of peroxidation		peroxidation
	Kinetics CD
	generation
	3 -Maximum amount of CD accumulated after the	CDmax (A 234 )	Quantity of peroxidabilizable PUFA	Ex vivo	Spectrophotom etry	Uvikon XS	[5]
	propagation
	phase
			Antioxidant capacity
	Total Antioxidant Status	TAS (mmol TEAC/L)	determined comparatively to "trolox equivalent antioxidant	Ex vivo	Spectrophotom etry	Uvikon XS	[6] adapted by [7]
			capacity" (TEAC)
	Vitamin A	Vit A (µg/mL)	Lipophilic antioxidant	In vivo	High performance chromatography liquid	HPLC Kontron UV/Vis Sys1 -detector	[8]
					High
	Vitamin E	Vit E (µg/mL)	Lipophilic antioxidant	In vivo	performance liquid -detector chromatography	HPLC Kontron UV/Vis Sys1 -detector	[8]
					UV/Vis
			Antioxidant capacity
	Antioxidant Capacity of water-soluble-portion	ACW (nmol/mL )	of "water-soluble ascorbic acid portion" comparatively to	In vivo	Emission spectroscopy	Photochem Analytic jena	[9]
			capacity
			Antioxidant capacity
	Antioxidant Capacity of lipid-soluble portion	ACL (nmol/mL )	of "lipid-soluble portion" comparatively to	In vivo	Emission spectroscopy	Photochem Analytic jena	[9]
			Trolox capacity

Table 2 :

 2 Experimental composition of the diets. Cows were given a straw-concentrate-based diet (C, n=9), or the same diet enriched with lipids rich in PUFA from extruded linseed and rapeseeds (40 g oil/kg diet DM) provided alone (L, n=22) or with vitamin E (155 IU/kg diet DM) (LE, n=17) or with vitamin E plus plant extracts rich in polyphenols (7 g/kg diet DM) (LEP, n=25).

	Experimental diets

Acknowledgements

Authors acknowledge the French research agency (ANR) for the financial support (Lipivimus 2007(Lipivimus -2009)). We thank technicians from the Biomarqueurs and Caraïbe teams at UMR Herbivores (INRA) for samples collections and analyses. We thanks techniciens from Herbipôle (INRA) for animal care. We thank Dr Olivier Berdeaux from CSGA (INRA) for his support in free 4-HNE and 4-HHE determination.

Competing interests

Authors declared no competing interests.