Saturn upper atmospheric structure from Cassini EUV/FUV occultations

Saturn upper atmospheric structure from Cassini EUV/FUV occultations D. E. Shemansky 1 and X. Liu 1 Planetary and Space Science Div., Space Environment Technologies, Pasadena, CA, USA dshemansky@spacenvironment.net ABSTRACT The Cassini UVIS experiment provides data on Saturn atmospheric physical properties in three observational programs: 1) Solar and stellar occultations provide transmission spectra in the EUV/FUV range that allow extraction of vertical profiles of H 2 and hydrocarbon abundances from the top of the atmosphere to about 300 km above the 1 bar pressure level. 2) Dayglow spectral images have been obtained which together with the occultation results constrain model calculations to provide properties of the activated atmosphere. The model calculations of the vertical profile of the state of the dayside atmosphere are further constrained by the Cassini radio science measurements of ionospheric structure. 3)Images of the magnetosphere show the escape profile of atomic hydrogen from the top of the Saturn atmosphere, inferred to be the product of the primary heat deposition process responsible for the high thermospheric temperatures first discovered in the Voyager Program. This section discusses the impact of the Cassini UVIS observations in relation to other results in this and earlier relevant research, and the implications for the understanding of the state of the atmosphere. 1. The Cassini UVIS experiment The Cassini UVIS (Ultraviolet Imaging Spectrograph) experiment is fully described by Esposito et al. (2004). The observations utilized in the work referenced here are obtained using the EUV and FUV spectrograph units identified as Channels (Esposito et al. 2004). The EUV and FUV Channels in a single exposure produce 64 spectral vectors of maximum length

2 1024 pixels (pxs). The sky projection of spatial pixel size is (pxs X pxw) = (0.25(spectral) X 1.0 (spatial)) mr, spectral pixel size is pxs = 0.6096524 Å (EUV) and pxs = 0.779589 Å (FUV). The spectral range of the UVIS is 563. 1182 Å (EUV) and 1115. - 1912. Å (FUV). The airglow ports of these instruments are used in stellar occultation observations in addition to spatially extended airglow measurements. Unlike the Voyager UVS experiment, the UVIS contains a solar occultation port only in the EUV Channel. The unfortunate significance of this exclusion for the exploration of the Saturn atmosphere is that UVIS solar occultations cannot provide information on the hydrocarbon structure. This is a serious limitation given the restricted number of high quality stellar occultations available in the Cassini Program. Descriptions of the three observational programs, occultations, dayglow exposures, and images of the escape of atomic hydrogen from the top of the atmosphere, obtained using the UVIS are described below. Analysis of the data obtained to date is incomplete. A limited number of interpreted reductions are presented. The three observational programs described here each provide critical independent supporting physical information allowing the development of a consistent description of the state of the upper atmosphere. The primary limitation at this time is the lack of adequate measured latitudinal and longitudinal distributions. 2. Determination of vertical atmospheric structure from occultation measurements Occultations in the EUV/FUV photon spectrum provide measurements of atmospheric structure and composition from the exobase to 300 km altitude above the 1 bar pressure level at Saturn. The Cassini UVIS experiment provides the most accurate platform to date for extracting atmospheric structure in outer planet research primarily through higher spectral resolution, signal rates, and dynamic range. To date 15 stellar occultations have been obtained over a range of latitudes of 43 50. To date 7 solar occultations have been obtained over a range of latitudes of 66 60. 2.1. Stellar occultations The quality of the stellar occulations depends on the magnitude of the star, the vertical rate of passage, and the latitude variation of the impact parameter through the atmosphere. Figure 1 shows a quantitative evaluation plot of the current stellar occultation events. Half of the occultations, as shown in Figure 1 fall below the level at which an analysis can be

3 carried out with full spectral capability, limiting accuracy and the number of hydrocarbon species that can be measured to CH 4. By far the highest quality occultation occurred at lat 40, 2005 DOY 103. The remaining occultations at southern latitudes are poor quality and can be reduced only photometrically as indicated in the figure. Of the occultations at northern latitudes 6 can be evaluated spectroscopically. Results from the analysis of the 2005 DOY 103 occultation, and a reanalysis of the Voyager 2 (V2) UVS stellar occultation are described here. 2.1.1. H 2 vertical distribution The H 2 component is obtained in the EUV Channel. An example of the transmission spectrum is shown in Figure 2. The analysis of the transmission spectra is carried out through forward modeling, simulating the instrument response in detail. This process is required in order to account for the point spread function (psf) of the UVIS spectrographs. The transmission spectrum of a model atmosphere using a modeled stellar source spectrum is applied iteratively in simulated instrument output to obtain an optimal fit to the observations. The EUV spectral resolution in the core of the psf is δλ = 0.9 Å FWHM. The EUV transmission spectrum at Saturn is uncontaminated hydrogen over the entire altitude range to the point of complete extinction. The H 2 model used in the analysis was developed over the past several years successively at the University of Southern California (USC) and Space Environment Technologies (SET). The temperature dependent absorption cross section spectra for the analysis are calculated at a resolution λ 500000 to an accuracy of 5% for δλ those transitions that contribute measurably to the absorption spectrum. The transmission spectra are fitted using separate vibrational vectors of the ground state H 2 X(v:J) structure. Details of H 2 physical properties are described by Hallett et al. (2005a,b) for the non-lte environment that develops in the excited atmosphere. The fundamental physical quantities (coupled state structure, absorption probabilities, predissociation probabilities) used in these calculations are referenced by Shemansky et al. (2008); Hallett et al. (2005a,b). The calculation methodology assumes that rotational structure is in thermal equilibrium at the gas kinetic temperature, but vibrational population is non-lte (Hallett et al. 2005a). Detailed model calculations (Shemansky et al. 2008; Hallett et al. 2005a) show that in actuality there is some deviation from thermal populations in rotation, but evidently not enough to significantly affect the absorption spectral properties. H 2 X vibrational populations, however, deviate significantly from thermal and critically affect the state of the ionospheric plasma (Shemansky et al. 2008; Hallett et al. 2005a,b). It has been found that the spectral fitting process is more sensitive to rotational temperature than vibrational population distribution. In the Cassini UVIS observations, therefore, the atmospheric temperature is derived through

4 both iterative determination of rotational temperature, and the shape of the vertical abundance distribution in the hydrostatic model calculations (Shemansky et al. 2005). At lower altitudes the kinetic temperature is also constrained by the measurements of the absorption structure of the C 2 H 2 diffuse temperature sensitive ( C X) bands (Shemansky & Liu 2008; Wu et al. 2001). Figure 2 shows the observed extinction spectrum at impact parameter 929 km against the modeled non-lte calculation. Derived rotational temperature and H 2 abundance are given in the figure. The 2005 DOY 103 occultation was a dayside observation at solar phase 38, with subsolar latitude 18.1. Figure 3 shows the forward modeled hydrostatic density distribution anchored at the 1 bar level to 400 km using the Lindal, et al. (1985) results. Shemansky & Liu (2008) have reanalyzed the Voyager 2 (V2) UVS stellar occultation using the current H 2 model structure, and the resulting vertical H 2 profile is also shown in Figure 3. The V2 occultation was on the darkside (Smith et al. 1983) at a latitude of 3.8. The differences in density at a given altitude evident in the figure are mainly the consequence of the gravitation scale at the different latitudes of the two occultations (Shemansky & Liu 2008). Figure 3 also shows the modeled Helium distribution anchored at a [He]/[H 2 ] = 0.12 mixing ratio at 1 bar (Shemansky & Liu 2008). The [He]/[H 2 ] mixing ratio affects the modeled temperature structure in the vicinity of the mesopause, and the applied ratio is limited by the temperature dependence of the C 2 H 2 cross section (Shemansky & Liu 2008). In the upper thermosphere density differences are extended by the fact that the top of thermosphere temperatures at V2 and Cassini differ by more than 100K. The vertical partial pressure profiles from Voyager UVS and Cassini UVIS compared to the total pressure derived from CIRS at latitude -20 (Fletcher et al. 2007) are shown in Figure 4. The vertical temperature profile is discussed in Section 2.1.3. 2.1.2. Hydrocarbon vertical distribution The analysis of the UVIS FUV spectrograph stellar occultation data yields the hydrocarbon densities shown in Figure 3. Although there is an indication of the presence of other species in the transmission spectra (Shemansky & Liu 2008), the species for which vertical distributions are obtained are those shown in Figure 3, CH 4, C 2 H 2, and C 2 H 4. The evidence for other species is discussed by Shemansky & Liu (2008). The only species profile that can be reliably extracted from the V2 stellar occultation is CH 4 (Shemansky & Liu 2008) ( see Figure 3. The hydrocarbon homopause is just above 600 km in the UVIS occultation and at 900 km in the V2 occultation. In the reanalysis of the V2 occultation data by Shemansky & Liu (2008) it was concluded that altitudes above the 1 bar level cited by the original Smith

5 et al. (1983) work were consistent with the reference level used by Shemansky & Liu (2008) within a few km. The vertical displacement of the hydrocarbon homopause levels in the two cases is consistent with the vertical displacement of the H 2 densities (see Figure 3). The CH 4 distributions are anchored at the 1 bar level with the value ([CH 4 ]/[H 2 ] = 5.1 X 10 3 ) established by Flasar et al. (2005). As noted above the temperature dependent C 2 H 2 ( C X) bands are used to limit kinetic temperatures below 700 km. An example of the UVIS FUV transmission spectrum is shown in Figure 5 against a model fit. 2.1.3. Vertical kinetic temperature profiles The derived temperature profiles from UVIS 2005 DOY 103, V2 1981 DOY 238 (Shemansky & Liu 2008), the Hubbard, et al. (1997)(multiple Earth based) stellar occultations, and the CIRS profile at latitude -20 (Fletcher et al. 2007) are shown in Figure 6. The UVIS result at lat -40 shows a distinct mesopause at 545 km near a temperature of 120 K. The mesopause temperature is limited by the measured structure of the C 2 H 2 ( C X) bands. The hydrostatic model calculation of the structure confined by the measured H 2 profile at higher altitudes, and the radio occultation results at altitudes below 400 km, is dependent on the [He]/[H 2 ] mixing ratio. The mixing ratio (0.12) adopted by Shemansky & Liu (2008) is at the upper limit in this model calculation, because higher ratios drive the mesopause temperature above 120 K, in disagreement with the observed C 2 H 2 ( C X) absorption structure in this region. The UVIS result is the only existing derivation from the sunlit atmosphere. The 460 K thermosphere obtained from the 1981 occultation is 140 K higher than the result in 2005. 3. Dayglow emission and the subsolar activated atmosphere For the first time since the Voyager encounters in 1979-80, EUV/FUV spectra of the Saturn dayglow have been obtained with the Cassini UVIS. The UVIS spectra are the first observations of the excited atmosphere at solar minimum. The emission properties are qualitatively different from the Voyager observations. An analysis of the dayglow properties observed using UVIS has been reported by Shemansky et al. (2008). These results show the dayglow in H 2 band emission is accounted for entirely by solar fluorescence photoelectron excitation. The emission intensity also varies with the magnitude of the solar EUV flux. The H 2 band emission has been modeled in a 1-d calculation based on a detailed hydrogen physical chemistry model constrained by the Saturn atmosphere model discussed in Section 2, and the ionospheric measurements from the Cassini RSS observations (Nagy et al. 2006).

6 The model is a detailed calculation with physical chemistry rate coefficients established individually for each rotational level built into a single architecture ( 3400 states). The model calculation establishes testable state populations, and because of the fact that the photon radiation field is necessarily intrinsic part of the calculation, all emission transitions in the system are predicted from radar frequency to the EUV, and can be tested against observation (Hallett et al. 2005a,b; Shemansky et al. 2008). The emission observations were obtained at southern latitudes. Figure 7 shows the predicted spectrum of the H 2 Rydberg band system compared to UVIS EUV and FUV spectra, calculated using the Shemansky et al. (2008) 1-d model. The Shemansky et al. (2008) model calculations establishing the spectrum in Figure 7 propogate the depostion of solar photons through the top of the atmosphere to the point of extinction. The general properties of the pure hydrogen physical chemistry model as described by Hallett et al. (2005a) establishes an important intrinsic property that holds over a wide range of excitation and gas densities. This property, caused by a combination of bootstrapping reactions, is the strong tendency of the ion population to be dominated by H + 3. This results in a very short lived plasma, having relaxation lifetimes typically of order 2000 sec in the activated Saturn atmosphere. For this reason Shemansky et al. (2008) argue that atmospheric dynamics has little impact on the structure of the ionosphere. It is known that the anti-solar mid latitude atmosphere shows no evidence of a measurable ionosphere. Deep exposures on the darkside show no evidence for atmospheric emission other than the passive scattering of the LISM H Lyα radiation field. The Nagy et al. (2006) averages of the RSS occultations show definitive evidence that the ionosphere is a dayside phenomenon that builds during the 5 hour day from dawn to dusk, with observations in the 1200 km to 2000 km region showing factor of 15 dusk/dawn differences in electron density. Although Shemansky et al. (2008) have not calculated a time constant for establishing a steady state, a plausible explanation for the dawn/dusk effect is that in order for the plasma to develop, the H 2 X(v:J) population must develop into high vibrational levels before thresholds for high production rates of H + 3 are obtained. Figure 8 shows the calculated vertical profiles calculated for Cassini (HA) conditions. In order to understand the constraints on the hydrogen system model calculations, it must first be understood that if a pure hydrogen volume is excited by electron forcing without the loss of mass from the volume, the gas will relax to an [H]/[H + ]/[e] end-point, with a very rapid transition to [H + ]/[e] or [H] as a function of the forcing electron temperature. The partitioning of the species in Figure 8, [e hν ](photoelectrons), [e a ] (ambient electrons), [H + 3 ], [H + ], [H + 2 ], [H], and [H 2 ], therefore depends on the loss and acquisition of mass in the volume as the excitation process continues. The primary volumetric loss in this system is H, which is produced kinetically hot in a number of reaction processes. Differential ion diffusion is

7 neglected here because of the intrinsically rapid recombination process. Mass loss in the model calculation is H, which is replaced by inward diffusion of H 2 into the volume. The state of the gas in the 1-d calculation is then determined by the penetration of solar flux, constrained by the known H 2 vertical density distribution, which the model code must match, the measured [e a ], and the calculated relaxed steady state multiply scattered photoelectrons. The steady state photoelectron distribution is calculated in a multiple scattering (elastic and inelastic) relaxation system that feeds the population of relaxed ambient electrons. H + 3 is the major ion throughout the vertical profile shown in Figure 8. The energy distribution of steady state multiply scattered photoelectrons feeding the ambient electron population is shown in Figure 9, as calculated for several altitudes shown for the range 3 to 31 ev. It is this distribution that determines the excitation rates for the system, and obviously constitute the core loop in the calculation, as all rates and state populations in the system are heavily coupled to the electron differential energy distribution. Figure 10 shows the calculated H 2 X vibrational populations for several altitudes, indicating a very large deviation from LTE. The primary limiting factor in the H 2 X(v) distribution is the population of the ambient electrons. As the ambient electron density rises in this system, the H 2 X(v) population tends toward equilibrium with the electron temperature. At extremely high ambient electron densities, the H 2 X(v) distribution is forced to thermal equilibrium by the high rate of opposing electron deactivation and excitation collisions. 4. Direct observational evidence for the upper thermospheric heating process Cassini UVIS maps of the Saturn magnetosphere have revealed distinct atomic hydrogen distributions in the region inside 4 R S of planet center showing the gas escaping the top of the thermosphere (Melin et al. 2008). Figure 11 (Melin et al. 2008) shows the image of Saturn in H Lyα emission from a spacecraft viewing angle edge-on to the rings in a surface contour plot. The subsolar latitude is 23 with the sun on the right side of the image. The spatial pixel size is 0.1 R S. The figure shows a bright feature extending outward from the sunlit thermosphere at an angle about 8 below the ring plane. The H Lyα brightness of the peak contour is about 1000 R. The contour lines off the sunlit southern latitudes show atomic hydrogen escaping at all latitudes below the auroral regions. The anti-solar side of the planet shows an asymmetric distribution consistent with a combination of an orbiting and ballistic hydrogen source on the sub solar thermosphere, and consistent with the conclusions drawn from images obtained using the Voyager UVS (Shemansky & Hall 1992). The escape of atomic hydrogen from the top of the atmosphere requires a translational energy ranging from 5.5 ev at the equator to 7.2 ev at the poles. This evidently indicates heating of the

8 top of the thermosphere by the atomic hydrogen generated in dissociation processes. The forcing must be electron impact (Section 5.2), but the mechanism for inserting energy into ionospheric electrons at the top of the atmosphere is not evident. 5. Discussion and conclusions 5.1. Neutral atmosphere vertical structure The Cassini UVIS occultation measurements reveal a substantially colder thermosphere than the previous Voyager results. A reanalysis in which the Voyager results were established on the same model basis as those of Cassini UVIS has verified the original work (Shemansky & Liu 2008). The Saturn thermospheric temperature may be correlated with the major solar cycle. The Cassini UVIS temperature profile shows two distinctive properties: 1) A distinctive mesosphere with a temperature near 120 K (Figure 6), 2) A temperature peak near 1200 km at latitude 40 indicating localized heat deposition in the sunlit atmosphere over a range of 1 to 2 scale heights. The reanalysis of the Voyager results shows a much less distinctive mesosphere (Shemansky & Liu 2008) (Figure 6) similar to that derived by Hubbard, et al. (1997). 5.2. Activated atmosphere vertical structure The Cassini UVIS dayglow observations show a spectrum that is explained entirely by solar radiation deposition in both spectral content and absolute brightness (Shemansky et al. 2008). The Voyager observations required a dominant high altitude electron excited source to explain the spectrum in 1981 (Shemansky & Ajello 1983). The solar deposition manifesting in the observed emission in 2005, however, does not account for the upper thermospheric temperature (Shemansky & Liu 2008). The non-lte model calculations fitting the observed Cassini spectrum predict a short lived ( 3000 sec) plasma population dominated by H + 3 (Shemansky et al. 2008). The Shemansky et al. (2008) model calculation providing the fit to the data shown in Figure 7 departs radically from previous ionosphere calculations such as Moore et al. (2006), in which large amounts of H 2 O are required to quench the ionosphere. Shemansky et al. (2008) (see Hallett et al. (2005a)), however, point out that the previous model calculations compensate for seriously flawed hydrogen physical state calculations through the device of injecting H 2 O as a quenching agent. The primary reactions controlling rates in pure hydrogen are: e + H 2 X 1 Σ + g (v i : J i ) e + H 2 X 1 Σ + g (v j : J j ) (1)

9 H + + H 2 X (v : J) H + H + 2 X (v : J) (2) H + 2 X (v : J) + H 2 X (v : J) H + 3 + H (3) e a + H + 3 H + H + H (4) e a + H + 3 H 2 X 1 Σ + g (v : J) + H (5) e s /e hν (E i ) + H 2 X(v : J) H 2 b + e s /e hν (E j ), (6) H 2 b 3 Σ + u H + H (7) H + H 2 X(v i : J i ) H + H 2 X(v j : J j ) (8), where e a refers to the ambient electron population, e s /e hν (E i ), is the multiply scattered photoelectron population in state i (Figure 9), (v : J) refers to vibration/rotation state, and refers to kinetically hot. The state and lifetime of the hydrogen plasma depends critically on the level of activation in H 2 X. The critical limiting reaction (1) ( 9600 electron excitation/deactivation transitions) is not present in any of the earlier work (see Hallett et al. (2005a), Shemansky & Liu (2008)). Reaction (2) limits the population and lifetime of H +. The reaction chain (2 5) constitutes a bootstrapping process by generating highly excited H 2 X(v:J) that would end in a runaway reaction loop were it not for the limiting reaction (1). The reactions (4 7) generating kinetically hot atomic hydrogen deposit heat and excite H 2 X(v:J) further through reaction (8). The rates for these reactions with the exception of (8) are fully established and applied in Hallett et al. (2005a,b) and Shemansky et al. (2008). The state of the hydrogen weakly ionized plasma cannot be resolved without the inclusion of reactions (1 7) which require that the populations of H 2 X(v:J) and H + 2 X(v:J) be calculated at the rotational level. The activation state of H 2 X is also affected by solar photon fluorescence (Liu et al. 2007). H + 3 is the dominant ion throughout the Saturn ionosphere to 3000 km in the Shemansky et al. (2008) calculations and the introduction of H 2 O as a quenching agent is unnecessary in fitting the observed vertical structure of the ionosphere. Upper atmospheric heating is explained by the reactions (4 8) inferred to take place within 1-2 scale heights of the exobase by electron forcing. The direct observational evidence for this is the outflow of atomic hydrogen (Figure 11) from the sunlit atmosphere, which requires a minimum of 5.5 ev/atom at the equator.

10 REFERENCES Esposito, L. W., et al. 2004, The Cassini Ultraviolet Imaging Spectrograph Investigation, Space Sci. Rev., 115, 299-361 Flasar, F. M., et al. 2005, Temperatures winds and composition in the Saturnian system, Sci., DOI: 10.1126/science.1105806, 1247-1251 Fletcher, L. N., et al. 2007, Characterizing Saturn s vertical temperature structure from Cassini/CIRS, Icarus, 189, 457 478 Hallett, J. T., D. E. Shemansky, and X. Liu. (2005a), A Rotational-Level Hydrogen Physical Chemistry Model for General Astrophysical Application, Astrophys. J., 624, 448 461. Hallett, J. T., D. E. Shemansky, and X. Liu (2005b), Fine-structure physical chemistry modeling of Uranus H 2 X quadrupole emission, Geophys. Res. Lett., 32, L02204, DOI:10.1029/2004GL021327. Hubbard, W. B., et al., 1997, The structure of Saturn s mesosphere from the 28 Sgr occultations, Icarus, 130, 404-425 Lindal, G. M., D. N. Sweetnam, & V. R. Eshleman 1985, The atmosphere of Saturn: An analysis of the Voyager radio occultation measurements, AJ, 90, 1136-1146 Liu, X., D. E. Shemansky, J. T. Hallett, & H. A. Weaver 2007, Extreme non-lte H 2 in comets C/2000 WM1(Linear)AND C/2001 A2(Linear) 169, 458-471 Melin, H., D. E. Shemansky, & D. T. Hall 2008, The distribution of neutral gas in the magnetosphere of Saturn, Icarus, xx, xxx-xxx Moore, H., et al. 2006, Cassini radio ccultations of Saturn s ionosphere: Model comparisons using a constant waterflux, Geophys. Res. Lett., 33, L22202,doi:10.1029/2006GL027375 Nagy, A. F., et al. 2006, First results from the inospheric radio occulations of Saturn by the Cassini spacecraft, J. Geophys. Res., 111, A06310, oi:10.1029/2005ja011519, A06310 Shemansky, D. E., & J. M. Ajello 1983, The Saturn Spectrum In The EUV - Electron Excited Hydrogen, J. Geophys. Res., 88, 459-464 Shemansky, D. E., & D. T. Hall 1992, The distribution of atomic hydrogen in the magnetosphere of Saturn, J. Geophys. Res., 97, 4143 4161 Shemansky, D. E., A. I. F. Stewart, R. A. West, L. W. Esposito, J. T. Hallett, & X. M. Liu 2005, The Cassini UVIS Stellar Probe of the Titan Atmosphere, Science, 308, 978 Shemansky, D. E., X. Liu, & J. T Hallett 2008, Modeled vertical structure of the Saturn dayside atmosphere constrained by Cassini observations, Icarus, xxx, xxx Shemansky, D. E. & X. Liu, 2008, Vertical atmospheric structure of Saturn from Cassini UVIS occultions, Icarus, xxx, xxx Smith, G. R., D. E. Shemansky, J. B. Holberg, A. L. Broadfoot, B. R. Sandal, & J. C. McConnell 1983, Saturns upper atmosphere from the Voyager 2 solar and stellar occultations, J. Geophys. Res., 88, 8667-8678 Wu, C. Y. R., F. Z. Chen, D. L. Judge 2001, Measurement of temperature-dependent absorption cross sections of C 2 H 2 in the VUV-UV region, J. Geophys. Res, 106, 7629 This preprint was prepared with the AAS L A TEX macros v5.2.

11 Data statistical quality 12 11 10 9 8 7 6 5 4 3 2 1 0 Saturn UVIS star occultations to date signal rate quality photometric cutoff -40-20 -0 20 40 Latitude Fig. 1. UVIS star occultation quality.

12 1.2 1.0 Alt: 929 km T = 240 K; C (H 2 ) = 1.3x10 20 cm -2 Vib. Distribution: F13 0.8 I/I 0 0.6 0.4 0.2 0.0 900 950 1000 1050 1100 1150 λ (A) Fig. 2. UVIS EUV stellar occultation transmission spectrum obtained 2005 DOY 103 at effective impact parameter 929 km. The rotational temperature is iteratively determined assuming LTE. The vibrational population distribution is non-lte determined iteratively by fitting separate vibrational vectors into the model for optimal match to the spectrum.

13 h (km) 2000 1800 1600 1400 1200 1000 800 600 400 200 0 V2 UVS 1981 DOY 238 δ-sco Lat 3.8 o UVIS 2005 DOY 103 ε-ori Lat -40 o CH 4 UVIS Lat -40 o H 2 UVIS Lat -40 o H 2 V2 UVS Lat 3.4 o CH 4 V2 UVS Lat 3.8 o C 2 H 2 UVIS Lat -40 o C 2 H 4 UVIS Lat -40 o He V2 UVS Lat 3.8 o He UVIS Lat -40 o 0 2 4 6 8 10 12 14 16 18 20 Log{[n] (cm -3 )} Fig. 3. Stellar occultation derived densities of H 2, He, CH 4, C 2 H 2, and C 2 H 4 from UVIS 2005 DOY 103 (Lat -40 ) and H 2, He, CH 4 from the V2 Voyager 1981 DOY 238 (Lat 3.8 ) events.

14 h (km) 2000 1800 1600 1400 1200 1000 800 600 400 200 V2 UVS 1981 DOY 238 δ-sco Lat 3.8 o 0 UVIS 2005 DOY 103 ε-ori Lat -40 o 10-11 10-10 10-9 10-8 10-7 10-6 10-5 10-4 10-3 10-2 10-1 10 0 10 1 10 2 10 3 P (mb) CH 4 UVIS Lat -40 o H 2 UVIS Lat -40 o H 2 V2 UVS Lat 3.4 o CH 4 V2 UVS Lat 3.8 o C 2 H 2 UVIS Lat -40 o C 2 H 4 UVIS Lat -40 o He V2 UVS Lat 3.8 o He UVIS Lat -40 o CIRS_Fletcher 2007 CIRS P total Lat -40 o Fig. 4. The equivalent of Figure 3 on a pressure scale. The CIRS results giving total pressure at latitude -20 (Fletcher et al. 2007) are included.

15 0-1 -2 UVIS FUV 2005 DOY 103 583 km model _042e_02-3 -τ -4-5 -6-7 1300 1400 1500 1600 1700 1800 1900 λ (A) Fig. 5. Cassini UVIS FUV stellar occultation optical depth spectrum against a model fit containing CH 4, C 2 H 2, C 2 H 4, and upper limits to C 2 H 6 and several other species(shemansky & Liu 2008). The temperature of the diffuse C 2 H 2 ( C X) band cross section in this calculation is empirically determined at 120 K. Weak contributions from dayside airglow appear in the short wavelength deep optical depth region of the spectrum, and in the 1500-1600 Å region. The effective impact parameter of this spectrum is 583 km.

16 h (km) 2000 1800 1600 1400 1200 1000 V2_lat4_00_reanal Cassini UVIS_rev6 _08_13_1 Hubbard et al. 1997 CIRS Fletcher 2007 V2 UVS 1981 DOY 238 δ-sco Lat 3.8 o UVIS 2005 DOY 103 ε-ori Lat -40 o Hubbard et al. (1997) Earth based stellar: equator model 800 600 400 200 Fletcher et al.(2007) CIRS Lat -40 o 0 0 100 200 300 400 500 T (K) Fig. 6. Saturn temperature profiles derived from the UVIS, V2, Hubbard, et al. (1997), occultations, and the Cassini CIRS result (Fletcher et al. 2007) at latitude -20.

17 B (10-6 c s -1 cm -2 ) 0.07 0.06 0.05 0.04 0.03 0.02 0.01 (a) 0.00 1150 1250 1350 1450 1550 0.05 (b) HA H 2 Model 0.04 REV7 Observation EUV 0.03 Atomic Hydrogen Model SM H 2 Model 0.02 FUV 0.01 0.00 900 950 1000 1050 1100 1150 Wavelength (D) Fig. 7. Model fits to the FUV and EUV Saturn dayglow spectra for the model atmosphere constrained by the Cassini UVIS rev6 occultation results, and the observed Cassini RSS ionosphere profile. H 2 discrete band and continuum emission, H Lyα and H Lyβ emission are included in the fit.

18 Altitude (km) 2200 2000 1800 1600 1400 1200 e hv H 3 + H H 2 H 2 + H + UVIS H 2 e a 1000 800 600 10-4 10-2 10 0 10 2 10 4 10 6 10 8 10 10 10 12 10 14 Density (cm -3 ) Fig. 8. Comparison of the HA modeled species distributions to the Cassini UVIS rev6 observation-derived H 2 density profile.

19 Photoelectron Distribution (cm -3 ) 10-1 10-2 10-3 10-4 h = 1750 km h = 1050 km h = 850 km h = 950 km 10-5 3 7 11 15 19 23 27 31 Photoelectron Energy (ev) Fig. 9. Photoelectron density distributions as a function of energy for selected altitudes produced from the rotational-level hydrogen chemistry and attenuated solar flux model, constrained by the Cassini UVIS Rev6 occultation-derived atmosphere.

20 H 2 X (v) / H 2 X 10 0 10-1 10-2 10-3 10-4 10-5 10-6 10-7 10-8 10-9 0 2 4 6 8 10 12 14 H 2 Vibrational Energy Level 2000 km 1650 km 1250 km 850 km T = 320 K Fig. 10. Selected HA model atmosphere H 2 X (v) normalized density distributions as a function of altitude. Model distributions are compared to a normalized thermal distribution at 320 K.

12 Fig. 5. Cassini UVIS image in a surface contour plot in H Lyα emission showing the escape of atomic hydrogen in a non uniform asymmetric distribution from the top of the Saturn atmosphere. The image pixel size is 0.1 0.1 R S. The edge-on view of the rings is indicated. The sun is off the right side of the plot with a sub-solar latitude of 23. Auroral emission is apparent at the poles extending over the terminator. Solar phase is 77.