1
The Saturn hydrogen plume
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D. E. Shemanskya , X. Liua and H. Melina
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a
Planetary and Space Science Division
Space Environment Technologies
320 N Halstead St., Suite 110
Pasadena, CA 91107, USA
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Images of the Saturn atmosphere and magnetosphere in H Lyα emission during the
Cassini spacecraft pre and post Saturn Orbit Insertion (SOI) event obtained using the
UVIS experiment FUV spectrograph have revealed definitive evidence for the escape of
H I atoms from the top of the thermosphere. An image at 0.1 × 0.1 Saturn equatorial
radii (RS ) pixel resolution with an edge-on view of the rings shows a distinctive structure
(plume) with FWHM of 0.56 RS at the exobase sub-solar limb at ∼−13.5◦ latitude as
part of the distributed outflow of H I from the sunlit hemisphere, with a counterpart
on the antisolar side peaking near the equator above the exobase limb. The structure
of the image indicates that part of the outflowing population is suborbital and re-enters
the thermosphere in an approximate 5 hour time scale. An evident larger more broadly
distributed component fills the magnetosphere to beyond 45 RS in the orbital plane in
an asymmetric distribution in local time, similar to an image obtained at Voyager 1 post
encounter in a different observational geometry. It has been found that H2 Rydberg
EUV/FUV emission spectra collected with the H Lyα into the image mosaic show a distinctive resonance property correlated with the H Lyα plume. The inferred approximate
globally averaged energy deposition at the top of the thermosphere from production of
the hot atomic hydrogen accounts for the measured atmospheric temperature. The only
known process capable of producing the atoms at the required few eV/atom kinetic energy
appears to be the direct electron excitation of non-LTE H2 X 1 Σ+
g (v:J) into the repulsive
,
although
details
of
the
processes
need
to
be
examined
under
the constraints imH2 b 3 Σ+
u
posed by the observations to determine compatibility with current knowledge of hydrogen
rate processes.
29
1. Introduction
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Images of the Saturn magnetosphere in H Lyα emission have been obtained utilizing
system scans using the Cassini UVIS FUV experiment (Esposito et al., 2004) from pre
orbital insertion (SOI) in 2003 and 2004 and the early post SOI period in 2005. In
the latter period an image obtained with the spacecraft viewing direction aligned with
the ring plane, with image pixel size 0.1 × 0.1 RS , shows remarkable structure in the
region extending outward from planet center to ±5 RS indicating the outflow of atomic
hydrogen in a distinctly asymmetric distribution from the sub solar atmosphere. This
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D. E. Shemansky
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high quality image from a unique observation geometry provides definitive evidence of
the atomic escape process at the top of the Saturn atmosphere. The implication of the
observed phenomenon and inferred flux is that the energy contained in the hot hydrogen
is sufficient to account for the high temperature at the top of the thermosphere.
The argument that the heating required to account for the temperature at the top of
the thermosphere arises in the physical chemistry of H2 is not new. Imaging obtained
with Voyager 1 UVS experiment system scans revealed a broad distribution of H I with
a local time asymmetry, showing a broad maximum abundance at dusk extending to the
orbit of Titan, and a minimum in pre-dawn longitudes (Shemansky & Hall , 1992). This
distribution was explained by Shemansky & Hall (1992) as sub-solar ejection of H I from
the top of the sub-solar thermosphere, primarily the result of electron impact dissociation
of H2 . Shemansky & Hall (1992) further argued that the magnetosphere, based on the
ambient plasma temperature, must contain neutral species, OH in particular, with suitable
radiative cooling properties to provide the needed plasma quenching. Shemansky et al.
(1993) discovered OH in the magnetosphere the following year in an inferred orbiting
torus distribution centered near 3.5 RS . It is now clear that the Saturn magnetosphere is
dominated by neutral gas, with the Cassini UVIS observation of atomic oxygen (Esposito
et al., 2005; Melin et al. , 2009). H I is much more broadly distributed in the Voyager 1
image, extending out to the orbit of Titan, and latitudinally to ±16 RS above and below
the orbital plane. The Voyager UVS image, however, was limited by the 1 RS X 1 RS pixel
size and a low signal to noise ratio (S/N). The images reported here have much higher
spatial resolution and S/N, and show orbiting H I extending beyond ±45 RS expanded
latitudinally ±20 RS surrounding the orbital plane.
The Cassini UVIS H I images allow a direct estimate of the energy contained in the
flux of gas out of the thermosphere, based on the shape of what can be described as a
plume originating in a south latitude range in the sub-solar atmosphere surrounding −
13.5◦ . The observed structure is definitive, but the explanation of the physical processes
is not evident and will require a substantial research effort.
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2. The Cassini UVIS Saturn system images
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The full extent of current Cassini UVIS maps of the Saturn magnetosphere are described
by (Melin et al. , 2009). The image considered here is part of the sequences of system scans
that began in December 2003 prior to the insertion of Cassini into the Saturn system.
Figure 1 (Melin et al. , 2009) shows the image of Saturn in H Lyα emission in a surface
contour plot from a spacecraft viewing angle edge-on to the rings. The H Lyα emission is
dominated by solar photon flux fluorescence everywhere in the system except in the planet
polar regions where emission is forced mainly by particle impact. The bright regions near
the planet are moderately optically thick (∼1). This image was accumulated in the period
2005 DOY 74 through DOY 86 over a spacecraft(S/C)-planet range of 24 to 44 planet
radii (defined as the equatorial radius)(RS ). The subsolar latitude is −22.3◦ with the sun
on the right side of the image at a phase of 77◦ . The spatial pixel size in the accumulation
matrix is 0.1 RS × 0.1 RS . The figure shows a bright feature extending outward from
the sunlit thermosphere at an angle of 13.5◦ below the ring plane, as outlined by the
image brightness contours. The H Lyα brightness of the peak contour is about 1100
The Saturn hydrogen plume
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R. The contour lines off the sunlit southern latitudes show atomic hydrogen escaping at
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). Quantitative brightness
slices of the Figure 1 image along lines of constant planetocentric latitude are shown in a
series of figures beginning with Figure 2 from planet center to 4.8 RS . Figure 2 contains
data at selected latitudes from pole to pole on the sub-solar side. The polar plots in Figure
2 encounter the auroral oval at a single point in the image matrix, indicating the projected
oval width is less than 6000 km wide. The north polar peak is a factor of 6 weaker than
the brightness of the south pole peak. The polar peaks occur at 0.9 RS , within a pixel
width of the polar nominal 1 bar radius. At a radial position of 1.1 RS , as shown in Figure
2, the polar plots show the brightness reaches a near constant value extending to 4.8 RS .
The signal in this region, as will be discussed below, is mainly emission from the extended
distribution of long lived orbiting atomic hydrogen, and not part of the gas population in
the near Saturn environment, defined here as the region inside 5 RS . The plot for latitude
−90◦ shows a dayglow brightness on the planet body of ∼800 R, compared to ∼200 R
for the latitude 90◦ . Given that the signal in the north polar plot is primarily from the
broader magnetosphere foreground, emission from the planet body along the north polar
line is below the detection limit, consistent with all previous observations of the Saturn
darkside below the auroral ovals. Figure 2 includes plots at latitudes, 0◦ , −13.5◦ , −27◦ ,
and 27◦ . The Figure 2 plot at −13.5◦ lies along the peak of the plume that constitutes
the most prominant feature of Figure 1. The Figure 2 plots at −27◦ and 27◦ indicate that
the main source of ballistic and escaping atomic hydrogen from Saturn is contained near
the central −13.5◦ latitude. The brightness distribution for 0◦ (Figure 2), falls almost a
factor of 2 below the plume peak in local hydrogen abundance at 1.5 RS and then merges
with the −13.5◦ distribution above 2.2 RS . This property is indicative of the broadening
of the plume with radial distance from planet center, and the confinement of the source in
the atmosphere. The Figure 2 plot for latitude 27◦ shows that the source of escaping and
ballistic hydrogen at this north latitude is very much weaker than the plume region. At the
south latitude, −27◦ shown in Figure 2, the peak emission near the limb is actually brighter
than at the plume core latitude, but drops well below the latitude −13.5◦ brightness in the
region above 1.3 RS . The full width at half maximum (FWHM) of the plume structure
is shown in Figure 4, indicating an approximately linear increase in width from 0.56 RS
at 1 RS from planet center to 1.7 RS at 4 RS from planet center. The brightness image
cannot be treated as in a uniform linear relationship with the line-of-sight (los) abundance
of gas. This is because the emission in the auroral oval is for practical considerations
entirely produced by precipitating electron excitation. These objects cannot be reduced
to abundances without undertaking a detailed analysis of the entire hydrogen emission
spectrum, beyond the scope of the present paper. Moving downward in latitude from
the auroral oval, the emission source transitions to dayglow, which on the planet below
the exobase is a combination of solar photon fluorescence, photoelectrons, and inferred
electrodynamically heated electrons. Above the exobase the hydrogen abundance can be
obtained by assuming the emission is entirely conservative scattering of the solar H Lyα
emission line. As indicated in Figure 2, the emission brightness shows a transition to a
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D. E. Shemansky
much lower slope in the regions above the poles, which is background/foreground to the
local emission, that extends to high latitudes. The background/foreground is interpreted
here as long lived orbiting H I that forms part of the extensive distribution described by
Shemansky & Hall (1992).
Figure 3 shows the approximate latitudinal boundaries of the auroral and dayglow
regions. It should be noted that the plot lines of constant latitude are limited by definition
to the region above the nominal 1 RS location. Because of los projection effects, the pixels
along radial lines in the region r < 1 RS cross lines of latitude on the planet surface, except
at 0◦ . As in Figure 2, Figure 3 includes the pixel brightness along the rotational axis as
a point of reference. The plot labeled 76.5◦ in Figure 3 marks the latitudinal limit to the
north auroral zone. The plot line in Figure 3 at latitude 63◦ marks the first indication of
dayglow sourced atomic hydrogen emission and first evidence of ballistic/escaping atoms.
At south latitude the Figure 3 plot at −81◦ marks the outer edge of the auroral oval. Note
that the auroral region brightness distributions show no evidence of measurable atomic
hydrogen in escape or ballistic trajectories above the background/foreground of the long
lived orbiting atoms, beyond 1.4 RS . It is assumed that the depth of auroral deposition
combined with the increased energy required for escape does not produce a measurable
atomic escape phenomenon. The plot in Figure 3 at −81◦ shows a peak emission at 1.1
RS suggesting high altitude electron excitation of atomic hydrogen, where this species
is the dominant neutral component. This phenomenon does not appear in the north
polar region. The line plot in Figure 3 at latitude −40.5◦ marks the onset latitude of
measurable ballistic and escaping atoms. The distinctive feature of the −40.5◦ line is the
very bright emission leading up to the 1.0 RS location from planet center. Here again is a
suggestion that electrodynamics may be driving electron excited emission at moderately
lower altitudes below the exobase. The latitude 0◦ line on the subsolar side is included
in Figure 3 for comparison. Figure 5 shows 4 E-W image slices of the Figure 1 image
at −0.2 and at ±1.0 RS , above and below the orbital plane, and one slice through the
plane. The slices at ± 1 RS pass through the polar auroral regions and show sharp peaks
close to the rotational axis, that merge into the broad background/foreground of orbiting
H I. The background/foreground signal is almost constant at ±1 RS (N-S) in the E-W
range ± 5 RS because of the large extent of the gas in this dimension, compared to the
N-S dimension. The slice in Figure 5 passing through the center of the planet shows the
local gas above the background/foreground signal with a large sharp peak just beyond 1
RS on the sub-solar side. This central slice shows a sharp dip in signal at −1 RS caused
by shadowing of the H I gas close to the planet. It is evident from Figure 1 that the
plume feature on the sub-solar side of the planet has a counterpart on the anti-solar side
indicating a consistent off-center tilt to the orbital plane of about 13.5◦ characterizing
this feature. Figure 6 shows H Lyα brightness along radial lines on the antisolar side of
the planet, including the equatorial subsolar and antisolar distribution identical to the
the 0.0 RS slice in the Figure 5 plot. The radial line at −76.5◦ marks the outer edge of
the auroral oval, to be compared to −81◦ on the subsolar side shown in Figure 3. The
Figure 6 plot line at −67.5◦ marks the first appearance of ballistic atomic hydrogen on
the antisolar side.
Figure 7 shows an image of the extended atomic hydrogen distribution obtained by
Cassini UVIS in 2003 providing a measurement reaching ±45 RS in the orbital plane.
The Saturn hydrogen plume
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The reader is reminded that the image in Figure 7 is a projection on a plane at the planet
tilted 13.5◦ relative to the image plane in Figure 1. The Figure 7 image shows the local
time asymmetry in the distribution of orbiting gas, similar to the result obtained using
the Voyager 1 UVS (Shemansky & Hall , 1992) in system scans at a 25◦ elevation to
the orbital plane (compared to the Cassini observations at low elevation). It is evident
that the magnetosphere contains a substantial population of orbiting long-lived atomic
hydrogen as well as the short-lived ballistic distribution inside 5 RS of planet center. Table
1 shows the derived total number of H I atoms (see Killen et al., 2009, for methodology
and solar model) in the system, and the estimated loss rate to escape and photoionization.
Table 1 includes current status for other neutral species in the system, O I, OH, H2 , H2 O,
and N I. The loss rate for H I in Table 1 is calculated for photoionization and escape, and
does not include the ballistically limited population. The tabulated loss rate therefore
does not represent the production rate at the source in the Saturn atmosphere. The rate
of H2 loss from Titan (Cui et al., 2008) is 0.3% of the hydrogen photoionization loss in
the system, and is therefore a negligible component. Based on the distinct structure of
the H I distribution inside 4 RS of planet center, and the physical limits of the activated
H2 source (Section 3), this component of the total population is considered to be mainly
ballistic and therefore short lived.
2.1. H2 band emission properties
The imaging of the hydrogen emission described here includes the entire (restricted)
spectrum of the UVIS instruments. In earlier dayglow measurements (Shemansky et al.,
2009) where the range to the planet and image pixel size was substantially larger, it was
assumed that the electronic H2 band system emission was uniform over the central dayside atmosphere. It has now been found that the emission brightness of the H2 bands as
well as the spectral properties change significantly with location on the sunlit hemisphere
correlating with the distinctive features in the atomic hydrogen image discussed above.
Spectra of the emission corresponding approximately to the latitudinal H Lyα plots in
Figures 2, 3, 5, and 6 have been extracted to further the physical interpretation. Figure 8
shows the FUV spectrum in the 1175 Å – 1375 Å region, at locations 0.9 RS and 1.1 RS
on the H Lyα plume latitude, −13.5◦ . The spectra shown here are limited by Cassini data
volume restrictions in which the spectral vectors are windowed and compressed, reducing
the spectral resolution below instrument capability. The spectra inside and outside the 1
RS boundary show H2 band emission that is unlike normal electron excited band structure. The spectrum shows strong features described here as resonances superposed on the
broadly distributed complex of overlapped electronic band transitions. The phenomenon
of resonance effects in non-LTE H2 has been recently reported in FUSE comet observations (Liu et al. , 2007). The spectrum of the region outside the 1 bar radius centered at
1.1 RS shown in Figure 8 shows the resonance features significantly reduced relative to the
underlying continuum of bands, compared to the spectrum inner region. The resonance
features are most prominent at the plume latitude, where they constitute about half the
total brightness of the bands. Figure 9 shows the H2 pre SOI mean dayglow spectrum
from Shemansky et al. (2009) compared to the spectrum at 0.9 RS at −13.5◦ latitude.
The resonance features are not evident in the mean dayglow spectrum in Figure 9, but
some features are visible in the original higher resolution version of the same spectrum
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D. E. Shemansky
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shown by Shemansky et al. (2009). Figure 10 shows the H2 spectrum at the edge of the
south auroral region at 1.1 RS and latitude −81◦ compared to the spectrum at 0.9 RS
at −13.5◦ latitude. The −81◦ spectrum shows the presence of the resonances at 1183,
1255, 1337, and 1352 Å, but the features in the spectrum at −13.5◦ at 1237, 1262, 1317,
and 1363 Å are not evident in the −81◦ spectrum. The broadly distributed H2 bands at
−81◦ are significantly brighter than the spectrum at −13.5◦ . Figure 11 shows the spectrum at 0.9 RS , −13.5◦ latitude, compared to the model calculation by Shemansky et al.
(2009) in the 100 km wide vertical segment at altitude 1950 km in the 1d model, scaled
upward by a factor of 33. The model calculation in Figure 11 contains solar forcing only,
and contains the resonance features associated with pumping by solar H Lyβ line. Other
resonances with solar discrete lines are not visible in this model because excitation of H2
X state populations are limited by the excitation process. Possible identifications of the
resonance transitions are indicated on the Figure 11 plot. There are numerous levels of
the H2 X state connecting to electronic states in resonance with the solar H Lyα line, that
could be responsible for pumping emission in the observed features, but rate processes
affecting energy deposition into the H2 X state are not fully developed in the present research program. Figure 12 shows the south pole aurora in comparison to the spectrum at
0.9 RS on the latitude −13.5◦ peak in the plume phenomenon. The auroral spectrum in
Figure 12 shows the presence of resonances at two (∼1190 and ∼1337 Å) of the 7 spectral
locations recognized in the plume spectrum, indicating the apparent physical complexity
of the excitation processes affecting the H2 X state vibration-rotation populations. The
relationship and absolute brightness of these features are discussed below in Section 3.
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3. Physical interpretation
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The H I plume described in Section 2 is an unanticipated feature of the outflow from
the top of the thermosphere. The correlated variation in H2 Rydberg emission properties
was also previously not anticipated from past observations. The physical chemistry of
activated hydrogen is known to produce kinetically hot H I products, as discussed below,
so the primary physics can be addressed. The production of energetic atomic hydrogen,
however, must take place within a scale height of the exobase for the majority of the
upward going atoms to appear as the ballistic and escaping components observed in the
magnetosphere. Otherwise the exothermic energy released in the activation would be
delivered directly into thermospheric heat without showing observable evidence in the
magnetosphere. The challenge in interpreting these observations is as follows:
1. A large flux of atoms is confined to a plume structure centered at latitude −13.5◦
showing a FWHM source region of 0.55 RS at the exobase limb. The outward
FWHM expansion is approximately linear with radial distance (Figure 4). In addition the peak brightness of the plume decreases proportionally to the FWHM,
indicating that the total number of atoms confined to the plume structure in any
given N-S image slice is approximately unchanged from the 1 bar limb (1 RS ) to at
least 4 RS . If, as the observations indicate, the source of atomic flux in the plume
is confined to the sunlit atmosphere, the explanation of the confinement in latitude
needs to be developed.
The Saturn hydrogen plume
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2. The energy contained in the atomic flux out of the atmosphere, as discussed below, is at least a factor of 10 larger than the direct deposition from solar photon
flux, and furthermore most of the solar energy is deposited far below the exobase.
The required energy deposition process and mechanism within a scale height of the
exobase needs exploration.
3.1. Properties of ballistic atoms launched from the exobase
The escape of atomic hydrogen from the top of the thermosphere requires a translational
energy ranging from 5.5 eV at the equator to 7.2 eV at the poles. The forcing must be
electron impact (Section 3.3), but the mechanism for inserting energy into ionospheric
electrons at the top of the atmosphere is not evident. At the equator approximately
0.5 eV in translational energy is provided by atmospheric rotation, and has an enhanced
latitude dependent effect on the ballistic property, in addition to the gravitational effect
of Saturn’s oblate shape. The lifetime of a vertically launched sub-escape atom for two
selected energies is shown in Figure 13 as a function of launch latitude. The lifetime
shortens rapidly with increasing latitude, as shown in Figure 13, in which a 3.06 eV atom
has a ballistic lifetime of 10 h at the equator (apogee = 3 RS ) compared to 3.6 h at 40◦
latitude (apogee = 2 RS ) and 2.2 h at 85◦ latitude (apogee = 1.7 RS ). For this reason,
depending on the energy distribution of the atomic hydrogen product in the activated
gas, a significant latitudinal restriction in the flux distribution of atoms can develop apart
from possible spatial variation of the activation of the gas. Figure 14 shows the vertical
launch energy required to achieve ballistic lifetimes of 5 and 10 h as a function of latitude.
A 10 h lifetime at latitude 85◦ requires a launch energy of 5.8 eV, compared to 3.06 eV
at the equator.
3.2. Energy content of ballistic H I
The shape and abundance of the H I distribution inside 5 RS and the ballistic properties
discussed in Section 3.1, allow a rough calculation of the mean flux and energy bound into
this system. Table 2 lists the relevant quantities. The total number of H I atoms in the
magnetosphere inside 5 RS contained in the ballistic component is calculated to be 5.7 ×
1034 . Based on the finding that most of the H I atoms in the inner magnetosphere are
concentrated in the dusk region of local time (Shemansky & Hall , 1992), a mean lifetime
of 5 hours is applied to arrive at a rate of 3 × 1030 atoms s−1 for the ballistic component.
Assuming that half of the total energy is contained in the sub-escape population of atoms,
the globally averaged energy deposition rate at the top of the Saturn atmosphere is ∼0.1
erg cm−2 s−1 . This is the approximate magnitude of the energy deposition rate into kinetic
energy required to explain the temperature at the top of the atmosphere. It is known
that the upper thermosphere temperature has a strong latitude dependence (Shemansky
& Liu , 2009a) and a global mean temperature is not established. It is found that the high
upper thermosphere, 407 K, found by Shemansky & Liu (2009a), at low latitude (15.2◦ )
correlates with the low latitude concentration of energy deposition by atomic hydrogen,
and the temperature found at −42.7◦ , 320 K, is a region where much weaker fluxes of
escaping atoms are observed.
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3.3. Hydrogen physical chemistry and the ionosphere
This section discusses the limits of physical conditions necessary to provide a plausible
explanation for the observed hydrogen atomic flux out of the the top of the thermosphere.
The fact that H2 is homonuclear, and the ion products in an activated system develop a
feedback system that tends to increase the rate of dissociation, makes the realistic calculation of the state of the gas for any kind of forcing complicated to develop. Hallett et al.
(2005a) (see Hallett et al., 2005b) describe a detailed method that calculates the non-LTE
state of the gas using reaction rates specific to the individual molecular vibration-rotation
levels. This has allowed the development of a 1-d ionosphere model for Saturn (Shemansky et al., 2009) that is used here as a reference, in comparison to extensive observations
using the Cassini RSS occultation experiments as well as from earlier spacecraft encounters (Nagy et al., 2006). The Shemansky et al. (2009) work is limited to forcing by solar
photon deposition only, which accounts for < 10% of the energy required explain the
observed flux of atoms injected into the magnetosphere or the measured temperature at
the top of the thermosphere.
The Cassini UVIS dayglow mid latitude observations show a spectrum that is explained
entirely by solar radiation deposition in both spectral content and absolute brightness
(Shemansky et al., 2009), with the exception of distinct features in the higher resolution
spectra in the pre SOI period described by Shemansky et al. (2009), that indicate some
H2 X vibrational levels are highly populated. It was assumed prior to examining the
high spatial resolution spectra in contained in the image of Figure 1, that the dayglow
was relatively uniform. This, however, is not the case as demonstrated in the results
described in Section 2.1. The Voyager observations required a dominant high altitude
electron excited source to explain the spectrum in 1981 (Shemansky & Ajello, 1983).
A comparison of the Cassini and Voyager quantitative results is given in Table 3. The
non-LTE model calculations fitting the observed Cassini mid latitude mean spectrum
predict a short lived (∼1 hour) plasma population dominated by H+
3 (Shemansky et al.,
2009). The solar forced ionospheric model, however, cannot predict the observed H2
band emission in the atomic hydrogen plume region, as shown in Figures 9 and 11. The
Shemansky et al. (2009) model calculation providing the fit to the data 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. (2009) (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. Limits set for the mixing ratio of H2 O by UVIS stellar occultations
and the actual reported observed vertical abundance fall 1 or more orders of magnitude
below the calculated mixing ratios of the H2 X(v>3) population responsible for quenching
H+ (Shemansky & Liu , 2009b), as described in the following discussion. 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
H + H2 X Σ+
g (vi : Ji )
1 +
H2+ X 2 Σ+
g (v : J) + H2 X Σg (v : J)
ea + H3+
ea + H3+
H2+
2
Σ+
g
(1)
(vj : Jj )
(2)
→ H3+ + H
→ H∗ + H∗ + H∗
∗
→ H2 X 1 Σ+
g (v : J) + H
(3)
(4)
(5)
→ H+
X
The Saturn hydrogen plume
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3 +
es /ehν (Ei ) + H2 X 1 Σ+
g (v : J) → H2 b Σu + es /ehν (Ej )
3
Σ+
u
∗
∗
H2 b
→ H +H
1 +
H + H2 X 1 Σ+
g (vi : Ji ) ↔ H + H2 X Σg (vj : Jj )
H + + H ∗ → H +∗ + H
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(6)
(7)
(8)
(9)
, where ea refers to the ambient electron population, es /ehν (Ei ), is the multiply scattered
photoelectron population in state i, (v : J) refers to vibration/rotation state, and ∗ refers
to kinetically hot. The full list of reactions in the model calculations are described by
Hallett et al. (2005a). 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). Reaction (2) limits the population and lifetime of H+ , but is exothermic
and rapid only with H2 X(v>3) states. The reaction chain (2 – 5) along with reaction
1 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 (4). 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
established and applied in Hallett et al. (2005a,b) and Shemansky et al. (2009). The
state of the weakly ionized hydrogen 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 at least 2000 km in the Shemansky et al. (2009) calculations. Upper
atmospheric heating is produced in the exothermic reactions (4 – 8) that must take place
within 1 - 2 scale heights of the exobase by electron forcing if the observed outflow of H
I into the magnetosphere is to be explained.
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, replaced by inward diffusing H2 , the gas will
relax to an [H]/[H+ ]/[e] end-point, with a very rapid transition to [H+ ]/[e] if the forcing
electron temperature is higher than 10000 K. The partitioning of the species in Figure
15, [ehν ](multiply scattered photoelectrons above 1 eV/electron), [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 reactions 4, 5, and 7. Differential ion
diffusion is neglected here because of the intrinsically rapid recombination process. Mass
loss in the model calculation is H, which is replaced by 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 [ehν ]. The steady state photoelectron distribution is calculated
in a multiple scattering (elastic and inelastic) relaxation system that feeds the population
of relaxed ambient electrons (Shemansky et al., 2009). H+
3 is the major ion throughout
the modeled vertical profile shown in Figure 15. The energy distribution of steady state
multiply scattered photoelectrons feeding the ambient electron population determines the
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excitation rates for the system, and 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. As noted above, solar photo deposition cannot account for the required heat
input at the top of the atmosphere. Figure 15 includes the weighted statistical averages of
the Cassini RSS experiment results in the dawn/dusk measurements (Nagy et al., 2006).
Nagy et al. (2006) report significant variations in the electron density profiles, and the
statistical variance in density from the average are typically factors of 2. The model
calculations shown in Figure 15 are limited to an altitude of 2000 km above the 1 bar
level, where the ability to match observed H2 density with the electron density forced by
observation fails. The observed profiles show substantial electron densities measurable to
5000 km (Nagy et al., 2006) . It is not possible for the Shemansky et al. (2009) model to
fit the observed structure above 2000 km with a pure hydrogen system using solar forcing
alone. Table 4 shows the rates and Table 5 shows atmospheric rate quantities used in
the present calculations to explore energy deposition and the state of the gas. Note that
the globally averaged energy deposition rate in Table 5 differs slightly from the value
given in Table 1 because the computational basis is different in the two cases. At 2000
km the electron/H2 mixing ratio ([e]/[H2 ]) is in the range (0.8 – 6) × 10−6 , depending
on whether the measured dawn or dusk averaged data is used, and 0.9 × 10−6 if the
Shemansky et al. (2009) model value is applied (see Table 4). Above this altitude, the
[e]/([H] + [H2 ]) projected mixing ratio rapidly rises, inferring the need for substantial
electron heating and energy deposition in model calculations in order to match the state
of the gas. This has not been explored in detail in using the present (Shemansky et
al., 2009) model. The mixing ratio [H]/[H2 ] is also a rapidly rising quantity at 2000 km
([H]/[H2 ] = 0.01) in the model calculation, and the values given in Table 5 are uncertain
for the ambient H I profile because they depend on the limitation to solar forcing in the
Shemansky et al. (2009) model. Projection of the H2 model to 3500 km based on the
Cassini UVIS occultation measurements (Shemansky & Liu , 2009a) is shown in Figure
15 with the electron profile from Nagy et al. (2006). [e]/([H] + [H2 ]) ∼ 5 × 10−4 (dusk)
to 9 × 10−4 (dawn) in Table 5 at 3500 km, but the density of H I below the exobase has
not been extracted from the Cassini UVIS occultations to date. [H]/[H2 ] = 0.03 in Table
5 at 3500 km, where [H] is equally partitioned between ambient and hot atoms. The
kinetically hot H I produced as part of the hydrogen physical chemistry diffuses upward
and downward from the reaction volume, and in the Shemansky et al. (2009) model
the H I density is constrained by the volumetric diffusive loss parameter determined by
iteratively adjusting diffusive loss until the observed H2 density is matched by the model
calculation. The production of hot atomic hydrogen in the solar forced activation of
the gas is several orders of magnitude below the rate inferred from the observed outflow
into the magnetosphere. Furthermore the recombination of the dominant ion, H+
3 , in
reactions 4 and 5 also falls short of the required deposition energy at any ambient electron
temperature (see Table 5). The only possible source for the hot H I, if it is to be explained
through hydrogen physical chemistry exothermicity, is the excitation of the repulsive H2
b state by electrons in reactions 6 and 7. The solar forced activation of the hydrogen gas,
however, produces an ambient population of electrons close to the temperature of the H2
because of the large sink of electron energy in momentum transfer and the large cross
sections for excitation/deactivation of vibration-rotation in the H2 X state in reaction
The Saturn hydrogen plume
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1. The rapidly rising ratio of electrons to neutrals above 2000 km (Figure 15, Table 5)
indicates that the neutral population above the exobase is depleted by ionization at rates
much higher than solar input can provide. The implication is that the Saturn H2 density
profile is overestimated above the exobase because the hydrostatic model takes no account
of dissociation loss. The region between 2000 km and 3500 km is evidently a region of
+
transformation from an H+
3 dominant ionosphere to H dominance, and an intrinsically
longer lived plasma (see Table 5). The exobase at latitude 15.2◦ is 2210 km, below which
the solar flux controls the formation of the ionosphere. In the kinetic collisional region
below the exobase, the plasma is dominated by H+
3 according to the model, and the
lifetime of the plasma is of the order of 1 hour. This plausibly explains the fact that
the dawn electron profiles at and below 2000 km are depleted on average by an order of
magnitude relative to the dusk profiles(Figure 15). At 3000 km the dawn/dusk electron
densities are similar, but at 4000 km the dawn average is consistently factors of roughly
2 above the densities at dusk (Table 5; Nagy et al., 2006). At the peak of the observed
plume the H I density is 1.4 × 104 cm−3 (at an altitude of ∼4000 km) assuming a base
radius of 0.56 RS . This density refers to the kinetically hot component which according to
the approximate conservation of atoms in the plume is the dominant atomic population
at ∼4000 km.
3.4. Energy distribution of the atomic products of dissociation
The recombination of H+
3 (reactions 4 and 5) produces discrete energies/atom because
of the state specific products. The dominant product is complete dissociation to atoms
(reaction 4), which produces 1.59 eV/atom for H+
3 X(v = 0). Reaction 5 produces atoms
at energies ranging from 3.18 eV/atom to 6.15 eV/atom. The probability distribution
of the discrete energies is plotted in Figure 16. H+
3 is assumed to be in the ground
vibrational state in these calculations. The energy deposition rate in H+
3 recombination
is not significant in the present calculations (Tables 4,5).
The reactions 6 and 7 produce continuum distributions of kinetically hot atomic hydrogen. As discussed in Section 4, the electron temperature required to produce the
observed flux of atoms through these reactions in the present simple calculations results
in a prediction of emission in the H2 Rydberg systems that that when combined with solar
forcing, agrees approximately with observation. It is clear that the partitioning of these
two quantities depends on the level of activation of the H2 X state. The H2 X state is expected to be significantly non-LTE if the main deposition process is near the exobase, and
this is indicated by the observed spectra in the plume region (Figure 8). If the electron
temperature is in the order of 20000 K, the populations in H2 X will be forced toward
equilibrium with the electrons through the reactions 1. The relative rates of electron
excitation of the H2 Rydberg systems and the dissociation process will depend on the H2
X(v:J) distribution. The rates given in Table 4 refer only to excitation from H2 X(0). The
rates for the excitation of the Rydberg states are accurately known and modeled. The
rates for reaction 6 for H2 X(v>0) are not determined because of the unknown electronic
form factor, although Franck Condon factors for the continuum distribution have been
calculated to provide the energy distribution of the atomic products. These calculations
are shown in Figure 17 for electron impact on H2 X(v<6) at asymptotic energies. Photon
excitation of this transition is electron spin forbidden. Energies limited to about 7 eV are
12
D. E. Shemansky
466
indicated for electron impact on the H2 X(v<6) levels. The distributions shown for the
levels in Figure 17 are normalized, and relative H2 X(v) strengths are not established.
467
4. Discussion and conclusions
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The images from the Cassini UVIS system scans reported here show that atomic hydrogen is flowing out of the top of the Saturn atmosphere in a mix of ballistic, orbiting,
and escaping components. Analysis of Voyager 1 UVS system scan observations also arrived at this conclusion (Shemansky & Hall , 1992), but the spatial resolution is 10 times
higher in the UVIS image of the near planet region, giving much greater detail. A plume
of hydrogen is observed at the sub-solar limb centered at ∼-13.5◦ with FWHM 0.56 RS in
the latitudinal plane, expanding outward to FWHM 1.7 RS at 4 RS from planet center.
At 1.05 RS the H I density in the plume feature is 1.4 × 104 cm−3 assuming the effective
path in the orbital plane is 0.56 RS , and 1.5 × 103 cm−3 at 4 RS (Table 5 gives mean
values over the plume fan). On January 17, 2006 the measured H I density between the
spacecraft (x,y,z= −5.50, −1.13, 0.562 RS ) and Enceladus (x,y,z= −3.22, −2.225, 0.441
RS ) was 1.6 × 103 cm−3 at a Saturn solar phase of 61◦ . The density of H I at Enceladus
will depend on Saturn local time of day according to these observations, and will vary
from ∼450 cm−3 to ∼1.6 × 103 cm−3 . The measured density of O I in the January 17,
2006 observation event is 530 cm−3 . H I is flowing out of the sub-solar atmosphere at
other latitudes also in an asymmetric distribution and there is an evident broader distribution filling the entire Saturn system to well beyond the bow-shock. The amount of
energy in this process, an estimated globally averaged deposition rate of ∼0.1 erg cm−2
s−1 , is sufficient to maintain the temperature at the top of the thermosphere (Shemansky
et al., 2009).
4.1. Kinetically hot atomic hydrogen source processes
The only plausible mid latitude source of the observed several eV/atom atomic hydrogen
entering the magnetosphere appears to be in the rate processes that take place in an H2
dominated gas activated by electrons. The analysis of the Saturn dayglow from the Cassini
UVIS experiment (Shemansky et al., 2009) shows H2 Rydberg emission (319 R)that is
explained by solar deposition alone. These results differ from the Voyager encounter
results which obtained an H2 dayglow brightness of 916 R (Shemansky & Ajello, 1983) ,
dominated by electron excitation, when the estimated upper thermosphere temperature
was 440 K (Shemansky & Liu , 2009a). The Saturn dayglow in H Lyα emission is ∼1100
R in the Cassini observation, compared to 4900 R at the Voyager encounter (Shemansky
& Ajello, 1983). These comparative results are shown in Table 3. The insolation deposits
a globally averaged energy rate of ∼0.01 erg cm−2 s−1 . The required deposition rate
at the top of the atmosphere is ∼0.12 erg cm−2 s−1 to maintain a mean 320 K upper
thermospheric temperature. This is the energy flux involved with the present observed
magnetospheric atomic hydrogen, within uncertainties in the calculation. The H2 Rydberg
system emissions corresponding to the H I plume structure differ from the mean mid
latitude spectrum in containing much larger resonance components that are not currently
modeled, as discussed in Section 2.1. These features double the total H2 band emission
brightness relative to the mid latitude spectrum. The calculations shown in Table 5 gives
the electron excited H2 band emission rate required for excitation by 20000 K electrons
The Saturn hydrogen plume
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corresponding to the rate of energy deposition from the measured H I distribution in
the near Saturn environment. This calculation is based on excitation of H2 X(v=0) into
the repulsive H2 b state. This is therefore only a crude estimate of what is required for
electron temperature in the exobase region, that would provide an explanation for the
hot H I production. The current detailed physical chemistry model that predicts H2 band
emission properties does not predict the observed spectrum in the H I plume region, and
will clearly require more research to investigate reactions that may be responsible for
the distinctive H2 X(v:J) population distribution responsible for the emission properties.
A physically plausible result is, however, obtained here because the current combined
H2 solar forced model and the electron excited component required to produce the hot
H I atoms gives a total predicted emission rate (for 20000 K electrons) just below the
observed H2 band emission rate in the H I plume structure. The physical considerations
that require further investigation beyond the scope of the present work are as follows:
1. The suggestion that the existence of the atomic hydrogen plume, confined to the
near equatorial region and orbiting in a displaced latitudinal plane at ∼ −13.5◦ , is
at least partially explained by the combined centrifugal effect of rapid rotation and
the oblate shape, needs further investigation. Empirical simulations of the observed
distribution need to be made to definitively establish the source distribution in the
atmosphere. The activation process is not understood, and is almost certainly non
uniform. More reduced UVIS system scans are needed to establish a 3-d image of
the atomic hydrogen distribution.
2. The subsolar upper thermosphere is subject to apparent electron impact excitation,
delivering 10 times the energy flux of insolation. The forcing of this system is not
understood. Rough calculations of emission in the H2 Rydberg systems that would
be expected along with the energy required to produce the kinetically hot H I, give
a total emission brightness that corresponds to the observations. The non-LTE
H2 X(v:J) populations required to reproduce the observed plume region spectrum
certainly are not predicted by the rate processes in the current model, and this
is the main uncertainty presented by the observations described here. Collision
strengths need to be developed particularly for reaction 8, which are not contained
in the architecture developed by Hallett et al. (2005a) (see Hallett et al., 2005b).
In addition absolute cross sections for reaction 7 are needed for electron impact
on H2 X(v>0). If the phenomenon is to be explained by the reactions outlined in
Section 3.3, it is necessary for substantial energy to be accumulated internally into
H2 X in order to predict the necessary energy/atom to match the observations. The
viability of this process as an explanation will hinge on the range of states of the
gas and detailed energy budget that can be established in physical chemistry theory
with known rate processes. The mechanism for the forcing of the strong resonance
structure evident in the H2 plume emissions is not apparent. Further laboratory
experimental work with hydrogen is necessary to guide rate process development.
3. The Cassini RSS results (Nagy et al., 2006) reporting the measured ionospheric
electron vertical profile provide important insight into ionization rates, although
more detailed model calculations are necessary. The ionosphere model calculations
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D. E. Shemansky
by Shemansky et al. (2009) (Figure 14) indicate that the observations can be modeled using solar deposition alone up to a limit of about 2000 km, although this is
somewhat misleading at the top of the atmosphere because of the clear presence
of a large flux of atomic hydrogen. In this calculation H2 O plays no role because
the mixing ratio cannot compete with reaction 2 (Shemansky & Liu , 2009b), and
the ionospheric density is limited by reactions 4 and 5. The lifetime of the ionosphere is 1 hour or less at and below 2000 km. The (Shemansky et al., 2009) model,
limited to solar input, cannot account for the observed ionosphere above 2000 km.
At 3500 km the neutral atmosphere is dominated by atomic hydrogen according to
the combination of observation and model, and the plasma is about 2% of the gas
density (Table 4). At 4000 km the electron density is still about 400 cm−3 in the
dawn average, and ∼200 cm−3 at dusk. The depleted average electron density at
dawn relative to dusk below 2000 km is compatible with the model calculation of
a ∼ 1 hour lifetime. It is certain that above 2000 km the dominant ion transitions
+
from H+
3 to H because of the rapidly rising [H]/[H2 ] mixing ratio. This results in
an unavoidable dramatically increased intrinsic plasma lifetime, while on average
the electron density drops by a factor of 2 during the 5 hour day above 3000 km.
The only reasonable explanation for a plasma lifetime less than 5 hours at 4000 km
is transport loss. Under ordinary circumstances this would be very hard to explain
given the confinement of the rotating magnetic field. A rough calculation, however,
gives a vertical flux of H I of 2 × 1010 cm−2 s−1 . This flux charge captures with
ionospheric H+ (reaction 9; Table 4) at a calculated rate of 2. × 108 cm−2 s−1 or
volumetrically about 0.02 captures cm−3 s−1 . The newly created H+ will have an
energy of roughly 3 eV on average, but a fraction of the population will be near
flat pitch angle with the magnetic field, and loss will be confined to the part of the
population at higher pitch angle. The vertically extended charge capture process
implies that the H+ population will be exchanged out with a time-constant of ∼6 h
(Tables 4, 5), so this is a plausible mechanism. The H+ transport loss time-constant
would then be slower than the charge capture rate by some indeterminate factor.
If that is the case, the upward diffusing atomic hydrogen would be converted to
new ions at a rate ∼7.5 × 10−9 cm3 s−1 by electron impact. This rate requires an
electron temperature of Te ∼13 eV. If a plausible time-constant of 10 h is assumed,
an electron temperature of Te ∼ 8.6 eV would be required. An electron temperature
this high requires an electrodynamic acceleration process. At 4000 km, ionosphere
electrons at 8 to 13 eV would not generate measurable H2 emission from excitation
of ground state H2 (see Tables 4, 5). These very rough calculations need verification with detailed model calculations on a global scale to explore the transport loss
process inferred from the measured ionosphere properties.
It is evident that our understanding of the phenomenon of the observed escape of kinetically hot atomic hydrogen from the Saturn atmosphere is at a primitive level that needs
extensive theoretical model exploration to limit the possible processes.
The Saturn hydrogen plume
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The Saturn hydrogen plume
Table 1
Neutral gas populations in the Saturn magnetosphere
Species Density (3 – 4 RS )
Total
Loss rate
(atoms cm−3 )
(atoms)
(atoms s−1 )
OI
500
3. × 1034
∼ 1029
OH
700
∼ 4. × 1034
∼1029
a
35
HI
450
2. × 10
3. × 1030b
H2 O
∼200c
–
–
NI
minor
–
–
a
Local time variable
Estimated escape rate
c
Theoretical orbiting at Enceladus and other sources
b
Table 2
Energy invested in atomic hydrogen
Totala
Energy rateb Global averagec
(atoms)
(erg s−1 )
(erg cm−2 s−1 )
34
19
5.7 × 10
5.1 × 10
0.10
a
Total number of atoms inside 5 RS of system center,
constituting 25% of total system population
b
Assuming 5 eV/atom, 5 hour mean lifetime
c
Upper thermosphere deposition rate
17
18
D. E. Shemansky
Table 3
Comparison of Cassini and Voyager Saturn dayglow/auroral observations
Quantity
UVISa
UVISb
UVISc
UVISd
I(H Lyα) (R)
1140.
1127.
1026.
1199.
I(H Lyβ) (R)
2.9
–
–
–
I(H2 Rydberg) (R)
319.i
606.
270.
468.
I(H Lyα)/I(H2 Rydberg)
3.60
1.86
3.80
2.56
a Cassini Apr - May 2004 mid latitudes (Shemansky et al., 2009)
b Post
SOI −13.5◦ latitude at 0.9 RS
c Post
SOI −13.5◦ latitude at 1.1 RS
d Post
SOI −45◦ latitude at 0.9 RS
e Post
SOI −45◦ latitude at 1.1 RS
f Post
SOI −81◦ latitude auroral peak at 1.1 RS
g Post
SOI −90◦ latitude auroral peak at 0.9 RS
h Voyager
i Emission
UVISe
633.
–
156.
4.06
UVISf
1398.
–
547.
2.56
UVISg
1781.
–
1096.
1.63
UVSh
4900.
10.
916.
5.35
1 dayglow mid latitudes 1980; H2 foreground (molec cm−2 ) = 5.× 1014 (Shemansky & Ajello, 1983)
attributed to solar deposition (Shemansky et al., 2009)
Table 4
Hydrogen rate processes
Te a
H2 (b)b
−12
(K) (10
cm−3 s−1 )
350
–
500
–
1000
–
5000
–
10000
0.85
15000
32.1
20000
208.
22300
340.
25000
648.
30000
1390.
a Electron temperature
b Reactions
cH
2
e + H+d
cm−3 s−1 )
–
–
–
6.82
4.18
3.15
2.51
–
–
–
(10−13
E(H∗ )f
(eV)
0.2
1.0
2.0
4.0
–
–
–
–
–
–
6 & 7 product rate; includes cascade from a,c,d,e states; derived from Khakoo et al. (1987)
9 recombination rate (Osterbrock, 1974)
e Reactions
4 & 5 product rate (2.5 atoms); derived from Larsson et al. (1993)
translational kinetic energy
g Charge
e
e + H+
3
−3
cm
s−1 )
92.5
79.0
56.5
11.6
8.5
–
–
–
–
–
(10−9
(B,C – X) emission rate; approximation to Liu et al. (1998)
d Reaction
f H∗
H2 (B,C – X)c
(10−12 cm−3 s−1 )
–
–
–
–
0.00124
0.183
2.32
4.70
10.9
31.1
capture rate; Barnett et al. (1990)
H∗ + H+g
cm−3 s−1 )
2.91
5.70
7.50
9.90
–
–
–
–
–
–
(10−9
The Saturn hydrogen plume
19
Table 5
Atmospheric partitioning and rates
hb
(km)
1570.
2000.
2500.
3000.
3500.
4000.
6000.
[edn ]c
(cm−3 )
490.
684.
928.
935.
500.
436.
–
[edk ]d
(cm−3 )
3329.
4838.
2009.
679.
270.
199.
–
[Ha ]e
(cm−3 )
3.5+07a
8.5+06
8.1+05
7.7+04
8.0+03
8.8−02
–
[Hh ]f
(cm−3 )
–
7.0+03
6.9+03
6.8+03
6.6+03
6.4+03
5.7+03
Flux and emission rates
F (Hh k )
F(Hhr )l
F(Hhb )m
(atoms cm−2 s−1 )
1.70+10
1.1+07
1.75+10
a Read 3.5+07 as 3.5 × 107 .
b Altitude
I(H2 (B,C))n
(R)
195.
[H2 ]g
(cm−3 )
7.8+09
8.8+08
7.2+07
6.1+06
5.4+05
5.0+04
–
[H]/[H2 ]h
–
4.5−03
9.7−03
1.1−02
1.4−02
2.7−02
1.5−01
–
above 1 bar at 15.2◦ latitude.
RSS mean dawn electron density (Nagy et al., 2006) equatorial.
d Cassini
RSS mean dusk electron density (Nagy et al., 2006) equatorial.
ambient H density from Shemansky et al. (2009) 1-d model.
f Measured
gH
2
[edk ]/[N]j
–
4.2−07
5.5−06
2.8−05
1.1−04
4.9−04
3.5−03
–
Ehr o
Ehb p
erg cm−2 s−1
4.2−05 9.8−02
c Cassini
e Nominal
[edn ]/[N]i
–
6.3−08
7.7−07
1.3−05
1.5−04
9.0−04
7.7−03
–
mean hot H density (2. – 7 eV) in the plume structure; scale height 19500 km; present work.
density at 15.2◦ latitude (Shemansky & Liu , 2009a) Hydrostatic model does not account for loss to dissociation at
high altitude.
h Mixing
ratio total nominal H density, lower limit values.
i Dawn
electron mixing ratio to total neutrals, lower limit values.
j Dusk
electron mixing ratio to total neutrals, lower limit values.
k Calculated
l Flux
flux from H Lyα image data.
from reactions 4 and 5 for Te = 20000 K. See table 4.
m Flux
from reactions 6 and 7 for Te = 20000 K. See table 4.
n Predicted
emission rate from electron excited (Te = 20000 K) H2 Rydberg systems from 1 scale height below exobase
(2210 km). See table 4.
o Calculated
energy deposition rate for reactions 4 and 5.
p Calculated
energy deposition rate for reactions 6 and 7.
20
653
654
655
D. E. Shemansky
Acknowledgements
This research was supported by the University of Colorado Cassini UVIS Program
contract 1531660 to Space Environment Technologies.
The Saturn hydrogen plume
21
2
Rs 0
2
4
hlya_post_01tt
, fft
2
0
Rs
2
4
Figure 1. 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. Image accumulated 2005 DOY 74 – 86 at S/C-planet range of 24 –
44 RS . The image pixel size is 0.1 × 0.1 RS . The edge-on view of the rings is indicated;
sub-S/C latitude is 0◦ . Range in the virtual image is indicated at the frame of the image
in units of RS , where 0,0 is the position of the planet center. Contour lines of constant
brightness are shown on the plot. Quantitative brightness values are plotted in following
figures. The sun is off the right side of the plot with a sub-solar latitude of −22.3◦ .
Auroral emission is apparent at the poles extending over the terminator. Solar phase is
77◦ .
22
D. E. Shemansky
1800
lat 90o
lat -90o
lat 0o subsolar
lat 27o subsolar
lat -27o subsolar
lat -13.5o subsolar
1600
1400
(R)
1200
1000
800
600
400
200
0
0
1
2
3
4
5
r (RS)
Figure 2. Selected latitudinal slices of the image in Figure 1 on the west (sub-solar)
side of the planet at radial distances ranging from planet center to 4.8 RS , showing the
distribution of the expanding plume of H I atoms into the magnetosphere. The vertical
scale is H Lyα brightness in Rayleighs. The horizontal scale is radial position relative
to planet center, where negative latitudes refer to the southern hemisphere. The plots
show emission brightness along lines of constant planetocentric latitudes above 1 RS .
Some pixels are empty because of incomplete coverage of the image matrix. The local
interplanetary medium has a brightness of about 40 R in this pointing direction. The
rings are edge-on at latitude 0◦ , and show no measurable reflectivity. The brightness data
is plotted as filled circles for north latitudes, open circles for south latitudes. The latitudes
of the plotted curves are given in the plot legend (see text). The polar latitude plots and
the north 27◦ latitude plot terminate at an approximate constant value of about 300 R
just above 1 RS . This signal is the forground/background of extended atomic hydrogen
orbiting the system mainly beyond 4.8 RS , as may be observed in the large scale image
of Figure 7.
The Saturn hydrogen plume
23
1800
lat 90o
lat -90o
lat 0o subsolar
lat 76.5o subsolar
lat 63o subsolar
lat -81o subsolar
lat -40.5o subsolar
1600
1400
(R)
1200
1000
800
600
400
200
0
0
1
2
3
4
5
r (RS)
Figure 3. Selected latitudinal slices of the image in Figure 1 on the west (sub-solar)
side of the planet (see Figure 2), showing the approximate latitudinal boundaries of the
aurora and dayglow. The latitudes of the plotted curves are given in the plot legend (see
text). The plot line at −81◦ marks the approximate south auroral outer boundary on
the subsolar hemisphere, and the line at −40.5◦ marks the onset of the dayglow process
measurably ejecting atomic hydrogen. The plot line for 63◦ marks the onset of measurable
atomic hydrogen ejection at northern latitudes.
24
D. E. Shemansky
1.8
1.6
FWHM (RS)
1.4
1.2
1.0
0.8
0.6
0.4
1
2
r (W) (RS)
3
4
Figure 4. FWHM of the atomic hydrogen plume on the sub-solar side of the planet,
derived from the data in Figure 2. The expansion slope is 0.36.
The Saturn hydrogen plume
25
1400
-0.2 RS
0.0 RS
-1.0 RS
1.0 RS
LISM
1200
(R)
1000
800
600
400
200
0
-5
-4
-3
-2
-1
0
1
E-W (RS)
2
3
4
5
Figure 5. East-west slices of the image in Figure 1 passing through the north and south
polar aurorae (open circles), the center of the planet and at -0.2 RS where the peak of the
H I plume occurs at the limb. The latter slice falls off the northern edge of the plume at
about 3 RS (see figure 2) and merges into the brightness profile of the r = 0.0 slice. The
±1.0 RS slices plateau off the auroral peaks to a background of about 250 R. Solar flux
is from right to left. The interplanetary background is indicated by a dashed line.
26
D. E. Shemansky
1800
lat 90o
lat -90o
lat 0o subsolar
lat 0o antisolar
lat -63o antisolar
lat -76.5o antisolar
lat -67.5o antisolar
1600
1400
(R)
1200
1000
800
600
400
200
0
0
1
2
3
4
5
r (RS)
Figure 6. Selected latitudinal slices of the image in Figure 1 on the east (anti-solar) side
of the planet (see Figure 2), showing the approximate latitudinal boundaries of the aurora
and dayglow. The latitudes of the plotted curves are given in the plot legend (see text).
The plot line at −76.5◦ shows the aurora peaking at 1.1 RS as it does on the sub-solar side
at −81◦ . The plot line at −67.5◦ marks the approximate south auroral outer boundary
on the sub-solar hemisphere with a peak at 1.0 RS , and the line at −63◦ marks the onset
of the dayglow process measurably ejecting atomic hydrogen. The plot line through the
north pole shows the distinctive effect of crossing the terminator onto the dark atmosphere
where the combined atmospheric emission and foreground fall below the brightness of the
foreground/background in the region above 1 RS .
The Saturn hydrogen plume
27
Figure 7. Contour plot of H Lyα in the pre SOI image of the Saturn system obtained Dec
2003, showing the extent of atomic hydrogen occupying the magnetosphere. The solar
phase is 62.5◦ , subsolar latitude 23.1◦ , and sub-S/C latitude 13.2◦ . Image pixel size is 1.4
× 1.4 RS . The sun is located above the upper frame of the image. The image plane is
tilted 13.2 degrees to the rotational axis. The scale is given on the plot frame in units of
RS . The 20 RS wide ridge in the center of the image surrounds the orbital plane. The
major peak in the profile is slightly off planet center because it is concentrated on the
sub-solar hemisphere. The zero base plane is indicated in brown. Measurable emission is
obtained to ±30 RS above and below the orbital plane. The image in the E-W direction
extends ±30 RS . Measurable emission from the system extends beyond 45 RS downstream
from the sun. Local time asymmetry is evident. The small peak above the orbital plane
is a star.
28
D. E. Shemansky
1.00
UVIS fuv lat -13.5o 0.9 RS
UVIS fuv lat -13.5o 1.1 RS
0.90
S (c s-1 px-1)
0.80
0.70
0.60
0.50
0.40
0.30
0.20
0.10
0.00
1175 1195 1215 1235 1255 1275 1295 1315 1335 1355 1375
λ (A)
Figure 8. UVIS FUV spectra of the atomic hydrogen plume image pixels located at
latitude −13.5◦ and radial positions 0.9 RS and 1.1 RS (Figures 1 and 2). The spectra are
significantly degraded in spectral resolution and spectral range relative to normal UVIS
function because of spacecraft limits on data volume during the observation sequences.
The superposed spectra are identified in the plot label. The spectra contain the H Lyα
line and blended H2 electronic band transitions. The strong resonance features have not
been previously observed in Saturn spectra, and are stronger at the base of the plume
than at any other location on the planet.
The Saturn hydrogen plume
29
1.00
0.90
S (c s-1 px-1)
0.80
UVIS fuv lat -13.5o 0.9 RS
UVIS fuv mid lat dayglow pre SOI
0.70
0.60
0.50
0.40
0.30
0.20
0.10
0.00
1175 1195 1215 1235 1255 1275 1295 1315 1335 1355 1375
λ (A)
Figure 9. Comparison of mean mid latitude FUV dayglow spectrum from Shemansky et
al. (2009) with the spectrum at the base of the plume at 0.9 RS (see figure 8). Resonance
features are not visually evident in the mid latitude dayglow spectrum shown here, but
the un-degraded version of the spectrum shown by Shemansky et al. (2009) does contain
visible resonance features that do not appear in the solar forced model.
30
D. E. Shemansky
1.00
0.90
UVIS fuv lat -13.5o 0.9 RS
UVIS fuv lat -81o 1.1 RS
S (c s-1 px-1)
0.80
0.70
0.60
0.50
0.40
0.30
0.20
0.10
0.00
1175 1195 1215 1235 1255 1275 1295 1315 1335 1355 1375
λ (A)
Figure 10. Comparison of the FUV emission spectrum at the auroral peak in the sunlit
atmosphere at 1.1 RS ,latitude −81◦ , with the spectrum at the base of the plume at 0.9
RS . The strong high altitude auroral H2 spectrum shows evidence of the same resonances
as the −13.5◦ latitude spectrum at ∼1187, ∼1336, and ∼1354 Å but not for the other
strong features in the plume spectrum.
0.20
0.10
RB 2(6,7) 1365.21
PB 1(6,7) 1365.66 PB 4(11,9) 1365.36
RB 2(0,4) 1335.12
RB 2(5,6) 1335.6
PB 1(0,4) 1335.87
PB 2(9,8) 1354.58
0.30
RB 2(6,6) 1315.56 PB 4(14,9) 1314.43
0.40
PB 3(5,5) 1290.15 PB 1(10,1) 1289.0
PB 1(10,7) 1290.3 RB 1(2,4) 1289.27
0.50
RB 6(1,3) 1271.
0.60
PB 1(6,5) 1265.7
0.70
QW 6(8,12) 1257.8
S (c s-1 px-1)
0.80
H Lyα
α
0.90
PB 8(1,1) 1183.3
QW 6(8,9) 1188.0
1.00
31
PB 8(1,2) 1237.9
PB 7(15,7) 1235.97
QW 6(8,11) 1238.75
The Saturn hydrogen plume
0.00
1175 1195 1215 1235 1255 1275 1295 1315 1335 1355 1375
λ (A)
Figure 11. Comparison of the model calculation of the solar forced Saturn emission
spectrum calculated for the 1d 100 km vertical segment at 1950 km altitude (magenta line),
with the spectrum at the base of the plume (green line) at 0.9 RS . The model spectrum
is scaled upward by a factor of 33. The model shows resonance peaks corresponding
to fluorescence of the solar H Lyβ line but there is no evidence of the other resonances
observed in the plume spectrum. The plot shows possible H2 transition identifications
corresponding to the resonances, using standard nomenclature, but the correspondence is
not inclusive of all possible contributors.
32
D. E. Shemansky
1.00
0.90
UVIS fuv lat -13.5o 0.9 RS
UVIS fuv lat -90o 0.9 RS
S (c s-1 px-1)
0.80
0.70
0.60
0.50
0.40
0.30
0.20
0.10
0.00
1175 1195 1215 1235 1255 1275 1295 1315 1335 1355 1375
λ (A)
Figure 12. Comparison of the FUV emission spectrum at the auroral peak in the sunlit
atmosphere at 0.9 RS , latitude −90◦ , with the spectrum at the base of the plume at 0.9
RS . The strong auroral H2 spectrum shows evidence of the same resonances as the −13.5◦
latitude spectrum at ∼1187, ∼1336, and ∼1354 Å, but not for the other strong features
in the plume spectrum, as in the spectrum at −81◦ (Figure10).
The Saturn hydrogen plume
33
10
Evert= 3.06 eV
τ (h)
8
6
Evert= 2.55 eV
4
2
0
10
20
30
40
50
60
Lat (deg)
70
80
90
Figure 13. Calculated lifetime of vertically launched hydrogen atoms at 3.06 and 2.55 eV
as a function of latitude.
34
D. E. Shemansky
6
W= 10 h
Evert (eV)
5
W= 5 h
4
3
2
0
10
20
30
40
50
60
70
80
90
lat (deg)
Figure 14. Kinetic energy of vertically launched hydrogen atoms to achieve ballistic
lifetimes of 5 and 10 h, as a function of latitude.
The Saturn hydrogen plume
Altitude (km)
3500
2500
35
ehv
H3+
H
H2
H2+
H+
UVIS H2
ea
e RSS dusk
e RSS dawn
1500
500
10-4 10-2 100 102 104 106 108 1010 1012 1014
Density (cm-3)
Figure 15. 1-d model calculation of the Saturn ionosphere (Shemansky et al., 2009) based
on solar forcing, constrained by the range of the measured electron and H2 profiles. The
model fails to converge on measured electron density above 2000 km. Electron density
[ehν ] is the multiply scattered solar photoelectron population with energy above 1 eV. The
Electron population labeled ea is the ambient relaxed population. The temperature of the
ea population is close to the atmospheric temperature because of the large cross sections
for reaction 1. H+
3 is the dominant ion up to 2000 km. The modeled ionosphere lifetime
is 1 hour at 2000 km. The mean electron densities from the dawn and dusk Cassini RSS
experiment (Nagy et al., 2006), identified on the legend, are plotted to 3500 km although
the RSS data extends to 5000 km.
36
D. E. Shemansky
100
7
5
4
3
p (atom-1)
2
10-1
7
5
4
3
2
10-2
7
5
4
3
2
10-3
1
2
3
4
5
6
7
E (eV)
Figure 16. Discrete probability distribution for dissociative recombination of H+
3 . The
dominant reaction path is 4, giving 1.56 eV atoms. Recombination to H2 X(v) + H
produces the distribution shown at higher energies. Based on the measurements and
calculations of Strasser et al. (2001).
The Saturn hydrogen plume
37
Normalized Probability (eV-1)
0.5
V=0, J=1
V=1, J=1
V=2, J=1
V=3, J=1
V=4, J=1
V=5, J=1
0.4
0.3
0.2
0.1
0.0
0.0
1.0
2.0
3.0
4.0
5.0
E (eV/atom)
6.0
7.0
8.0
Figure 17. Probability distributions of H(1s) kinetic energy of reactions 6 and 7 for
electron impact at asymptotic energy on H2 X(v = 0 -5, J=1). Note that excitation
from v = 6-14 levels is not shown. The calculation is based on the X state potential
energy curves of Schwartz & Le Roy (1987) and b state potential energy of Staszewska &
Wolniewicz (1999).