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 H2 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 CH4 . 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. H2 vertical distribution
The H2 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 H2 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 specλ
tra for the analysis are calculated at a resolution δλ
∼ 500000 to an accuracy of ∼ 5% for
those transitions that contribute measurably to the absorption spectrum. The transmission
spectra are fitted using separate vibrational vectors of the ground state H2 X(v:J) structure.
Details of H2 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. H2 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 C2 H2 diffuse temperature sensitive (C̃ – X̃) 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 H2 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 H2 model structure, and the resulting vertical H2 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]/[H2 ] =
0.12 mixing ratio at 1 bar (Shemansky & Liu 2008). The [He]/[H2 ] 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 C2 H2 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, CH4 , C2 H2 , and C2 H4 . 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 CH4 (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 H2 densities (see Figure 3). The
CH4 distributions are anchored at the 1 bar level with the value ([CH4 ]/[H2 ] = 5.1 X 10−3 )
established by Flasar et al. (2005). As noted above the temperature dependent C2 H2 (C̃ –
X̃) 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 C2 H2 (C̃ – X̃) bands.
The hydrostatic model calculation of the structure confined by the measured H2 profile at
higher altitudes, and the radio occultation results at altitudes below 400 km, is dependent on
the [He]/[H2 ] 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 C2 H2 (C̃ – X̃) 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 H2 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 H2 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 H2 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 H2 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, [ehν ](photoelectrons), [ea ] (ambient electrons), [H+
3 ],
+
+
[H ], [H2 ], [H], and [H2 ], 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 H2 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 H2 vertical density distribution, which the model code must match,
the measured [ea ], 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 H2 X vibrational populations for several altitudes, indicating a very large deviation from LTE. The primary limiting factor in the H2 X(v) distribution is the population of the ambient electrons. As the ambient electron density rises in this
system, the H2 X(v) population tends toward equilibrium with the electron temperature.
At extremely high ambient electron densities, the H2 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 RS 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 RS . 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.
5.1.
Discussion and conclusions
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 H2 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 H2 O as a quenching agent. The primary
reactions controlling rates in pure hydrogen are:
1 +
e + H2 X 1 Σ+
g (vi : Ji ) ↔ e + H2 X Σg (vj : Jj )
(1)
–9–
H + + H2 X (v : J) → H + H2+ X (v : J)
(2)
H2+ X (v : J) + H2 X (v : J) → H3+ + H
(3)
ea + H3+ → H ∗ + H ∗ + H ∗
ea +
H3+
→ H2 X
1
Σ+
g
(v : J) + H
es /ehν (Ei ) + H2 X(v : J) → H2 b + es /ehν (Ej ),
(4)
∗
(5)
(6)
H2 b 3 Σ+
→ H∗ + H∗
u
(7)
H + H2 X(vi : Ji ) ↔ H + H2 X(vj : Jj )
(8)
, where ea refers to the ambient electron population, es /ehν (Ei ), 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 H2 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 H2 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 H2 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 H2 X(v:J) and H+
2
X(v:J) be calculated at the rotational level. The activation state of H2 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
H2 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 H2 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 H2 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 C2 H2 in the
VUV-UV region, J. Geophys. Res , 106, 7629
This preprint was prepared with the AAS LATEX macros v5.2.
– 11 –
Data statistical quality
Saturn UVIS star occultations to date
12
11
10
9
8
7
6
5
4
3
2
1
0
signal rate quality
photometric cutoff
-40
-20
-0
Latitude
Fig. 1.— UVIS star occultation quality.
20
40
– 12 –
1.2
1.0
Alt: 929 km
T = 240 K; C (H2) = 1.3x1020 cm-2
Vib. Distribution: F13
I/I0
0.8
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 –
CH4 UVIS Lat -40o
H2 UVIS Lat -40o
H2 V2 UVS Lat 3.4o
CH4 V2 UVS Lat 3.8o
C2H2 UVIS Lat -40o
C2H4 UVIS Lat -40o
He V2 UVS Lat 3.8o
He UVIS Lat -40o
2000
1800
1600
h (km)
1400
1200
1000
800
600
400
V2 UVS 1981 DOY 238 δδ-Sco Lat 3.8o
200 UVIS 2005 DOY 103 εε-Ori
Lat -40o
ε
0
0
2
4
6
8
10
12
14
16
18
20
-3
Log{[n] (cm )}
Fig. 3.— Stellar occultation derived densities of H2 , He, CH4 , C2 H2 , and C2 H4 from UVIS
2005 DOY 103 (Lat -40◦ ) and H2 , He, CH4 from the V2 Voyager 1981 DOY 238 (Lat 3.8◦ )
events.
– 14 –
CH4 UVIS Lat -40o
H2 UVIS Lat -40o
H2 V2 UVS Lat 3.4o
CH4 V2 UVS Lat 3.8o
C2H2 UVIS Lat -40o
C2H4 UVIS Lat -40o
He V2 UVS Lat 3.8o
He UVIS Lat -40o
CIRS_Fletcher 2007
2000
1800
1600
h (km)
1400
1200
1000
CIRS P total Lat -40o
800
600
400
200
0
V2 UVS 1981 DOY 238 δδ-Sco Lat 3.8o
UVIS 2005 DOY 103 εε-Ori
Lat -40o
ε
10-11 10-10 10-9 10-8 10-7 10-6 10-5 10-4 10-3 10-2 10-1 100
101
102
103
P (mb)
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
UVIS FUV 2005 DOY 103 583 km
model _042e_02
-2
-ττ
-3
-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 CH4 , C2 H2 , C2 H4 , and upper limits to C2 H6 and several other species(Shemansky
& Liu 2008). The temperature of the diffuse C2 H2 (C̃ – X̃) 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 –
V2_lat4_00_reanal
Cassini UVIS_rev6 _08_13_1
Hubbard et al. 1997
CIRS Fletcher 2007
2000
1800
V2 UVS 1981 DOY 238 δδ-Sco Lat 3.8o
1600
UVIS 2005 DOY 103 εε-Ori
Lat -40o
ε
Hubbard et al. (1997) Earth based stellar:
equator model
h (km)
1400
1200
1000
800
600
400
Fletcher et al.(2007) CIRS Lat -40o
200
0
0
100
200
T (K)
300
400
500
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 –
0.07
0.06
(a)
FUV
0.05
0.04
B (10-6 c s-1 cm-2)
0.03
0.02
0.01
0.00
1150
0.05
1250
1350
1450
HA H2 Model
REV7 Observation
Atomic Hydrogen Model
SM H2 Model
(b)
0.04
0.03
1550
EUV
0.02
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. H2 discrete band and continuum emission, H Lyα and H Lyβ emission
are included in the fit.
– 18 –
2200
ehv
H3+
H
H2
H2+
H+
UVIS H2
ea
2000
Altitude (km)
1800
1600
1400
1200
1000
800
600
10-4
10-2
100
102
104
106
108
1010
1012
1014
Density (cm-3)
Fig. 8.— Comparison of the HA modeled species distributions to the Cassini UVIS rev6
observation-derived H2 density profile.
Photoelectron Distribution (cm-3)
– 19 –
10-1
h = 1750 km
h = 1050 km
h = 850 km
h = 950 km
10-2
10-3
10-4
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 –
100
H2 X (v) / H2 X
10
2000 km
1650 km
1250 km
850 km
T = 320 K
-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
H2 Vibrational Energy Level
Fig. 10.— Selected HA model atmosphere H2 X (v) normalized density distributions as a
function of altitude. Model distributions are compared to a normalized thermal distribution
at 320 K.
– 12 –
Fig. 115.— 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 RS . 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◦ .