TOURING THE SATURNIAN SYSTEM:
THE ATMOSPHERES OF TITAN AND SATURN
TOBIAS OWEN1 and DANIEL GAUTIER2
1 Institute for Astronomy, University of Hawaii (E-mail: [email protected])
2 Observatoire de Meudon (E-mail: [email protected])
Received 28 June 1999; Accepted in final form 17 March 2000
Abstract. This report follows the presentation originally given in the ESA Phase A Study for the
Cassini Huygens Mission. The combination of the Huygens atmospheric probe into Titan’s atmosphere with the Cassini orbiter allows for both in-situ and remote-sensing observations of Titan.
This not only provides a rich harvest of data about Saturn’s famous satellite but will permit a useful
calibration of the remote-sensing instruments which will also be used on Saturn itself. Composition,
thermal structure, dynamics, aeronomy, magnetosphere interactions and origins will all be investigated for the two atmospheres, and the spacecraft will also deliver information on the interiors of both
Titan and Saturn. As the surface of Titan is intimately linked with the atmosphere, we also discuss
some of the surface studies that will be carried out by both probe and orbiter.
1. Introduction
A1 comprehensive study of the atmospheres of Titan and Saturn has always been a
major goal of the Cassini-Huygens mission. For this review of atmospheric science
objectives, we decided to use the 1986 ESA Phase A Study as our guide. Our
purpose is to show that the spacecraft that is now on its way to the Saturn system
is indeed capable of fulfilling the original objectives established for the mission.
We also want to preserve the essence of the Phase A Study in an easily accessible
publication, as this is the document that won the approval of the mission from ESA
and NASA.
Reviewing this study, we found that a few instruments in the strawman payload
were ultimately not selected for the mission. Furthermore, our knowledge about
Saturn and especially Titan has increased significantly in the last decade. Nevertheless, the original scientific objectives remain valid, so we have kept the basic
format of the original document, revising the contents as appropriate.
During the early days of planning for this mission, well before even the Assessment Study was initiated, there was hope for a Saturn orbiter with two atmospheric
probes – one for Saturn and one for Titan. Budget constraints quickly eliminated
the Saturn probe from consideration, so the studies of Saturn’s atmosphere must be
1 Adapted from ESA/NASA Assessment Study: Cassini: Saturn Orbiter and Titan Probe, ESA Ref.
SCI (85)1 (1985)
Space Science Reviews 104: 347–376, 2002.
© 2003 Kluwer Academic Publishers. Printed in the Netherlands.
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carried out by remote sensing. It is therefore very helpful that the trajectory of the
spacecraft en route to Saturn permits a moderately close (140 RJ) flyby of Jupiter,
because data from the Galileo Jupiter Entry Probe (see Science, Vol. 272, No. 5263,
10 May 1996; J. Geophys. Res. Vol. 103, No. E10, 1998) provides a beautiful way
of calibrating the Cassini remote-sensing instruments during the flyby.
We expect a rich harvest of results from the Jupiter flyby, an enterprise that
was also envisaged in the original Phase A Study. However the main focus of this
report, mirroring that original study, is on Titan and Saturn, as it is the exploration
of the Saturnian system that is the ultimate goal of Cassini-Huygens. Our focus
remains the basic science objectives. We therefore ask the reader who is interested
in the specification of instrument capabilities to consult the relevant articles in this
issue that describe the individual instruments. We have cited these references as
volume 1 or volume 2 in the first instance each instrument is mentioned. For the
Huygens probe, additional instrument descriptions are collected in a special ESA
report (ESA SP-1177, August 1997, ed. J. P. Lebreton).
2. Titan
The encounters by Voyager 1 and 2 (summarized by Hunten et al., 1984) revealed the unique character of Titan in the solar system. Subsequent ground-based
and Earth-orbital observations have added information about Titan’s surface and
lower atmosphere, discovered additional atmospheric constituents (most notably
CO, CH3 CN and H2 O), determined isotopic ratios, provided new determinations of
atmospheric structure (see recent summary by Gautier and Raulin, 1997) provided
low resolution maps of the surface in the near IR (Meier et al., 2000) and the
first evidence for clouds (Griffith et al., 1998). We expect these additions to continue. Nevertheless, Titan will clearly remain an enigma in many respects, including the nature of its surface, the details of the complex photochemistry occurring
in its stratosphere, and the origin and evolution of its atmosphere. The CassiniHuygens Mission to Titan will address the major topics described in the following
paragraphs.
2.1. T HE ATMOSPHERIC COMPOSITION OF T ITAN
2.1.1. The Troposphere
2.1.1.1. Scientific Problem: Determine the tropospheric composition of Titan, specifically the relative abundances with altitude of CO, Ar, N2 , H2 , and CH4 . In contrast to the oxidized atmospheres of terrestrial planets, Titan’s predominantly nitrogen atmosphere is chemically reduced. The low tropopause temperature, around
70◦ K (see Figure 1) acts as a cold trap for most of the gases that could be present in
the troposphere and limits their amount in the stratosphere unless they are formed
there. Methane has been detected in the stratosphere in amounts that cannot exceed
ATMOSPHERES OF TITAN AND SATURN
349
3 or 4%, but it could have a mixing ratio as high as 10% at the 94◦ K temperature of
the surface (Courtin et al., 1995). Collision-induced absorption by nitrogen, methane and hydrogen contributes to a mild greenhouse effect. The presence of argon is
suspected for cosmogonic reasons, but it has not been firmly identified (Samuelson
et al., 1981; Owen, 1982). The upper limit on a heavy, volatile, spectroscopically
undetectable constituent such as argon has been steadily lowered until it is now
≤6% (at 3σ) (Courtin et al., 1995). This argon would be primordial 36 Ar and 38 Ar
brought to the satellite during its formation; 40 Ar produced by decay of 40 K in
Titan’s rocks should be less than 0.5% of the present atmosphere revised upward
from <0.01% (Owen, 1982) owing to escape of nitrogen (see below). Accurate
determinations of argon and other noble gases (and of their isotopes) are crucial to
test theories of the origin of the atmosphere (section 2.5).
Carbon monoxide poses a similar set of problems. At the present time (2003),
there are three different values for the CO mixing ratio: 60 ± 20 × 10−6 as determined from the 3-0 near-IR vibration-rotation band (Lutz et al., 1983; Courtin et al.,
2000) 10±5×10−6 deduced by Noll et al. (1996) from the (1-0) fundamental at 4.8
−6
µm, and 29+9
−5 × 10 , derived from microwave observations of the J (0-1), J(1-2),
and J (2-3) rotational lines (Hidayat et al., 1998). Gurwell and Muhleman (1995),
observing the J (0-1) rotational line, found a CO abundance of 50 ± 1 × 10−6 , in
good agreement with the (3-0) near-IR result. Further disagreement exists regarding the vertical distribution of CO, with Gurwell and Muhleman (1995) finding
the gas uniformly mixed, and Hidayat et al. (1998) reporting a decrease in the CO
−6
at 350 km.
concentration to a value of 5+4
−2 × 10
The CO abundance is important because of its role in atmospheric photochemistry and as a possible tracer of Titan’s original endowment of volatiles. CO is
converted to CO2 by the reaction.
CO + OH → CO2 + H
(1)
with the OH contributed by infalling H2 O (Samuelson et al., 1983). CO may be
produced by the attack of OH on CH4 , or it may be a relic from the origin of the
satellite (Samuelson et al., 1983; Lutz et al., 1983; Owen and Gautier, 1989). In
either case, reaction (1) suggests that a detectable amount of CO2 may have been
deposited on Titan’s surface during the last 4.5 billion years.
Other tropospheric constituents include molecular hydrogen, which was found
by Voyager to have a mixing ratio of about 0.1% (Courtin et al., 1995). This hydrogen is expected, it is a product of CH4 dissociation and shows that this gas
is being irreversibly destroyed in the atmosphere. Methane must condense in the
troposphere if its partial pressure reaches the limit defined by the saturation vapor
pressure law. The methane in the upper troposphere and ethane descending from
the stratosphere (see 2.1.2) would be the principal constituents of any cloud cover
that may exist below the tropopause.
The recent discovery of clouds on Titan by Griffith et al. (1998) using spectroscopy indicates that condensation indeed takes place, although the extent, fre-
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Titan's Lower Atmosphere
180
Cassini Probe,
Instruments On
160
0.003
140
Te
mp
era
t
ure
Uniform Haze
Polymers, Condensates
120
0.01
Pressure, bar
Altitude, km
100
80
0.03
60
0.1
40
0.3
20
Methane Cloud
1.0
0
1.5
Lake
80
100
120
140
Temperature, K
160
Figure 1. A cross section of Titan’s atmosphere, with various imagined surface features and a
methane rainstorm introduced by the artist.
ATMOSPHERES OF TITAN AND SATURN
351
quency, and location (are they associated with surface relief?) of these clouds
remains to be determined. AO imaging observations by Brown et al. (2002) and
Roe et al. (2003) revealed cloud systems over the South Pole of Titan. Accurate
values of the methane mixing ratio near Titan’s surface would help us to determine
the nature of the condensed materials that must exist there (see Section 2.4). The
lifetime of atmospheric methane is limited by photolytic destruction with subsequent production of other compounds (see 2.5). Strobel (1982) has estimated
that the current atmospheric abundance of methane will disappear in ∼ 20 × 106
years. Hence there must be a source of this gas on the satellite, either in the form of
outgassing from the interior or storage in a subsurface reservoir (Stevenson, 1992).
Similarly, the discovery that 15 N/14 N is 4.5 times the terrestrial value in Titan HCN
(Marten et al., 1997; Meier et al., 1997; Hidayat, 1998) suggests that an as yet unspecified amount of N2 has left the satellite (see 2.5). A determination of the exact
composition of the current troposphere as well as the time-integrated atmospheric
reservoir remain key problems. The fact that 12 C/13 C is normal (Hidayat et al.,
1997) suggests that a reservoir for CH4 exists, as does the irreversible destruction
of CH4 by solar UV (Strobel, 1982).
2.1.2. The Stratosphere
2.1.2.1. Scientific Problem: Determine the vertical and horizontal abundance distributions of trace hydrocarbons, nitriles and oxygen bearing compounds in the
stratosphere. Search for more complex organics including polymers in gaseous
and condensed phases. In addition to CH4 , six hydrocarbons C2 H2 , C2 H4 , C2 H6 ,
C3 H4 , C3 H8 , C4 H2 ) and five nitriles (HCN, HC3 N, CH3 CN, C2 N2 and C4 N2 ) have
been detected in Titan’s stratosphere (see Gautier, 1997, and Gautier and Raulin,
1997, for a review, and more recently, Marten et al., 2002). Oxygen compounds
are represented by CO, CO2, and H2 O. Except for CH4 , CO and possibly C2 H4 ,
the amounts of the detected species are much higher than would be permitted
by the saturation law at the tropopause temperature. This implies that they are
formed at stratospheric levels and above from a complex photochemistry initiated
by dissociation of methane, nitrogen and water vapor by solar UV photolysis and
bombardment by cosmic rays and high energy electrons trapped in Saturn’s magnetic field. H2 O was recently detected by ISO (Coustenis et al., 1998), its presence,
presumably from impacting particles, may be required to explain the formation
of both CO and CO2 . Photochemical models predict the formation of observed
hydrocarbons and nitriles from the radicals arising from CH4 and N2 dissociation
(Strobel, 1982; Yung et al., 1984; Toublanc et al., 1995; Lara et al., 1996). Condensation of these photochemical products and other more complex compounds
including polymers produce the aerosols that are responsible for the thick layer
of smog surrounding the satellite. These aerosols absorb sunlight, leading to a
temperature inversion and an increase of the static stability of the stratosphere
(Figure 1).
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Because of the short radiative time constant at upper atmospheric levels, the latitudinal temperature field should be distributed symmetrically about the equator at
the equinox. However, an asymmetric temperature field was observed by Voyager
at the 1 mbar level (Flasar et al., 1981). Latitudinal distributions of aerosol albedos
were also observed to be asymmetric by Voyager 1 suggesting that a complex interaction exists between thermal, chemical and dynamical systems (Sromovsky et al.,
1981). In particular, the remarkable change in the average albedo of the aerosols
that occurs just at the satellite’s equator is unique in the solar system, as is its
seasonal variation (Sromovsky et al., 1981; Caldwell et al., 1992). The distribution
and composition of trace molecules and aerosols in Titan’s stratosphere is thus a
key problem The acquisition of these data by Cassini at the local solstice compared
with Voyager data acquired at the equinox will permit a study of seasonal effects.
Of particular interest is the distribution of the “parent” molecule CH4 which may
not be uniformly mixed horizontally in the stratosphere.
One of the most important results of the Voyager encounter with Titan was
the discovery of three nitriles HCN, HC3 N and C2 N2 , subsequently joined by
CH3 CN discovered from Earth (Bézard et al., 1993) and a later detection of C4 N2
in the Voyager spectra (Samuelson et al., 1997). The synthesis of complex organic compounds from mixtures of simple gases of reducing composition has
been extensively studied in the laboratory since the famous experiments of Miller
and Urey over 45 years ago (Miller, 1955). Indeed, laboratory experiments have
demonstrated that HCN is a precursor of purines (in particular adenine) which are
among the building blocks of the nucleic acids in living systems on the Earth.
Similarly, HC3 N leads to pyrimidines which are also present in nucleic acids.
Although the composition of Titan’s atmosphere is certainly very different from
that of the primitive Earth, (it is so much colder that water vapor is only present in
the form of externally supplied traces), abiotic organic synthesis in the atmosphere
of Titan offers a test of how a Miller-Urey synthesis works at low temperature on a
planetary scale and what it produces over several billion years (Owen et al., 1992,
1997; Raulin, 1997; Raulin and Owen, 2002).
The key problem is thus to determine the abundances and distribution of organics in Titan’s stratosphere, to establish the degree of complexity these compounds
have achieved (both in the atmosphere and in subsequent reactions in any liquids
that may exist on the surface), and to determine the processes and pathways for producing them. A list of organic molecules likely to be present in Titan’s atmosphere
on the basis of laboratory experiments is given in Appendix 1.
Almost all expected organics condense at the temperatures occurring in Titan’s
lower stratosphere (Figure 1). They must form droplets or solid particles that will
precipitate with the larger aerosols to the surface where they will accumulate. At
tropospheric temperatures the organics have very low abundances in the gas phase;
they must pass through the troposphere as aerosols. The analysis of the composition of aerosols in the lower stratosphere and/or the troposphere is thus of great
importance.
ATMOSPHERES OF TITAN AND SATURN
353
2.1.2.2. Capabilities of Cassini. The Gas Chromatograph/Mass Spectrometer
(GCMS) aboard the Titan Probe will measure the elemental and isotopic composition from about 170 kilometers altitude down to the surface (Niemann et al., vol.
1). Through the use of a collector and pyrolyzer of aerosols (ACP) (Israel et al., vol.
1), it will also be able to detect condensed organics. Search for organics in gaseous
form and measurement of their vertical distributions can also be accomplished
by this versatile instrument. The far infrared-submillimeter spectrometer (CIRS)
aboard the Orbiter (Kunde et al., vol. 2) will provide, at each encounter with Titan,
the vertical distributions of HCN, HC3 N, C2 N2 and of some hydrocarbons (C4 H2 ,
C3 H4 ) at various locations on the satellite. CIRS will also determine isotope ratios
such as D/H and 15 N/14 N and will also permit the detection of (or determination of
upper limits for) new species such as other nitriles that are expected to be present
in Titan’s stratosphere (See Appendix 1). It will map the distribution of H2 O in the
upper stratosphere to determine the pattern of ice particle infall. The UV solar flux
penetration and the N2 dissociation will be investigated with the UV spectrometer
aboard the Orbiter (Esposito et al., vol. 2). The thermal structure and the abundances of CH4 and C2 H2 in the upper atmosphere will be obtained by observing the
sun or a star occulted by the limb of Titan (as successfully done by Voyager). A
combination of occultations and IR measurements should help to shed some light
on the temperature structure in this altitude range. VIMS (Brown et al., vol. 2) will
be able to study the CO fundamental at 4.7 µm and search for evidence of other
minor constituents in the long optical paths to Titan’s surface that are available
in near IR atmospheric windows. Both VIMS and the Imaging Science System
(ISS) (Porco et al. , vol. 2) will be able to search for clouds in the troposphere and
contrast features in the haze to track atmospheric motions. The Descent Imager and
Spectrometer (DISR) will also study clouds and constituent abundances along the
descent trajectory.
2.2. T HERMAL STRUCTURE AND METEOROLOGY OF T ITAN ’ S ATMOSPHERE
2.2.1. Scientific Problem: Measure winds and temperatures; investigate general
circulation and seasonal effects in Titan’s atmosphere
Voyager infrared observations of Titan’s stratospheric temperature field imply the
presence of 100 m s−1 zonal winds, super-rotating some 5–6 times faster than the
16-day rotation of the satellite body itself (Flasar et al., 1981). If this inference
is correct, then Titan’s large-scale atmospheric motions may be in a class with
those of Venus, which shows cloud-tracked wind speeds of comparable magnitude.
Unlike Venus, however, Titan has a large (27◦ ) obliquity more nearly comparable
to that of the Earth and may, therefore, show substantial meteorological variations
with changing seasons.
The super-rotating winds have been confirmed by the analysis of the occultation
of the star, zeta Sgr by Titan (Sicardy et al., 1990) and by the heterodyne infrared
based measurements of Kostiuk et al. (2000). The latter authors measured the
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difference in Doppler shifts of C2 H6 lines at the opposite limbs of Titan, demonstrating for the first time that the zonal winds in the upper atmosphere are prograde.
The general circulation of the atmosphere of Titan and its super rotation with respect to the ground have been modeled by Hourdin et al. (1995) who reproduced
rather well the wind pattern obtained by Sicardy et al. (1990).
Progress in the study of these features requires direct measurements of the zonal
winds, global mapping of temperature fields over a range of altitudes, and improved measurements of the vertical variation of clouds and aerosols. Zonal wind
measurements will confirm the assumed balance between large-scale motions and
north–south gradients in temperature previously applied to the analysis of Voyager
infrared observations. Global maps of the temperature fields at an “instant” in time
can then be used to extend the inference of thermal wind structure to other levels.
The strength of seasonal fluctuations can be inferred by a comparison of Voyager
and Cassini thermal measurements. Measured limits of the longitudinal variation
of temperature will provide important information on the role of wave dynamics in
the general circulation of Titan’s atmosphere. Global mapping of the vertical and
horizontal distribution of hydrocarbons and nitriles will provide important boundary conditions on dynamical models of the middle atmosphere. Measurements near
the surface (see Section 2.4) will provide further important clues about the possible
role of moist (hydrocarbon condensation) convection in the Titan Meteorology.
2.2.2. Capabilities of Cassini
The tropospheric thermal structure will be mapped at each encounter of the Orbiter with Titan by the far infrared submillimeter spectrometer of CIRS. The stratospheric thermal structure will be studied by means of the CIRS spectrometer
operating in the middle infrared. A large number of tropospheric and lower stratospheric temperature profiles will be provided by radio-occultation experiments
(Kliore et al., vol. 2). The Huygens Probe will measure the local temperature from
170 kilometers altitude down to the surface (Fulchignoni et al., vol. 1).
Doppler tracking of the Probe during its descent trajectory will provide a direct
measure of the zonal winds at one location (Bird et al., vol. 1). Wind velocities
will be determined to accuracies of 10 m s−1 or better at an altitude of 170 km
and 5 m s−1 near the surface. Measurements of latitudinal temperature gradients
can then be used to extend the zonal wind mapping into the troposphere. The
cloud structure along the descent trajectory will be observed by the Descent Imager
(Tomasko et al., vol. 1). Winds will also be measured directly by tracking clouds in
the lower troposphere observed by the imaging system and by VIMS on the orbiter.
The circulation in the stratosphere may be monitored if these same instruments
identify structures in the haze layer. These measurements, along with determinations of the vertical distribution of CH4 by the GCMS, will permit characterization
of convective processes in the lower troposphere. Inference of the net solar flux as
a function of height by the Descent Imager will provide information on the vertical
distribution of heat sources that drive the atmospheric circulation.
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2.3. T ITAN AERONOMY AND INTERACTIONS WITH S ATURN ’ S
MAGNETOSPHERE
The ultraviolet spectrometers on Voyagers 1 and 2 provided considerable information on Titan’s upper atmosphere through a solar occultation and studies of the
day-time airglow (see Hunten et al., 1984 for a review). The radio occultation
experiment yielded a marginal detection of an ionosphere, giving an upper limit
of 3000 electrons cm−3 a few degrees from the terminator (Bird et al., 1997). The
nature of Titan’s interaction with the magnetosphere leaves no doubt that an ionosphere exists. The upper atmosphere is mostly N2 , with several percent of CH4 and
a detectable amount of C2 H2 . H and H2 are probably present in the unusual dayglow of Titan. Although this emission is confined to the dayside of the satellite, it
is 5–10 times brighter than the glow that could be produced by the entire solar flux
below 1000 Å (angstrom). Since the phenomenon is so mysterious, its investigation
has a high priority. There is little doubt that the glow is excited by electrons with
energies between 10 and a few hundred eV. Such electrons should be sought with
in-situ instrumentation. These electrons are also the most likely source of Titan’s
ionosphere, rather than solar EUV radiation. Ionization produced by the interaction
+
of these low-energy electrons (∼100 eV) will consist primarily of N+
2 and N
as indicated from airglow observations. The principal ions in Titan’s ionosphere,
however, are expected to consist of nitrile (e.g., H2 CN+ ) and hydrocarbon (e.g.,
+
CH+
3 ) ions resulting from ion-molecule reactions between the originally formed N2
and N+ ions and CH4 present at ionospheric levels. Recombination of these ions
will be a major source of the HCN present in Titan’s atmosphere. Jupiter’s inner
magnetospheric composition and energetics are known to be dominated by heavy
ions of S and O ejected from Io. Ejecta from Titan (and other satellites) should be
similarly important to Saturn’s magnetosphere, though probably not dominant. A
neutral torus of H atoms was detected by the Voyager Ultraviolet Spectrometer and
H2 and N are expected as well along with H2 O and O from the icy satellites.
2.3.1. Capabilities of Cassini
The Orbiter will carry a full complement of remote-sensing and in-situ aeronomy
instruments (Blanc et al., vol. 1). The UV spectrometer will examine the dayglow emissions, determine the global distribution of H around Titan, and map
hydrocarbon distributions by observing stellar occultations. The ion and neutral
mass spectrometers will determine the composition of the upper atmosphere (Waite
et al., vol. 2).
Since many of the orbits pass through the upper atmosphere of Titan at altitudes
down to 800 km, close to the ionization maximum, several separate sets of insitu measurements will be accumulated over the four-year lifetime of the mission.
Repeated radio occultations and UV observations of occultations of the sun and
stars by Titan’s limb will add to this store of information. The mission design
offers an opportunity to compare inferences from remote-sensing observations with
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the “ground-truth” derived from the periodic in-situ measurements, including the
selected set carried by the Probe to the surface. Thus the global coverage available
from remote sensing can extend with confidence the localized but more detailed
coverage provided by the direct measurements.
2.4. T ITAN SURFACE AND INTERNAL STRUCTURE
2.4.1. Scientific Problem: Determine the nature and the composition of Titan’s
surface; infer the satellite’s internal structure
The interiors of Titan and Saturn have been included in the charter of responsibilities for the Atmospheres Working Group. We are adding a discussion of the
surface of Titan here for both historic and scientific reasons. This topic was certainly included in the original Phase A Report, and in fact, the surface of Titan has
a particularly important connection with the atmosphere, both as a potential source
and reservoir for volatiles that can replenish the atmosphere, and as a repository
for condensates and aerosols that must constantly precipitate from the atmosphere.
(Additional details will be found in the article by Hunine and Soderblom.)
Titan boasts the largest unexplored surface in the solar system. This mysterious domain is perpetually masked by thick layers of aerosols that prevented the
Voyager cameras from recording any surface detail. However, this ubiquitous haze
can be penetrated at near infrared, infrared and radio wavelengths, with the result
that observations of Titan from Earth have succeeded in sensing the satellite’s
surface. Radar was the first to show that the surface is inhomogeneous, with a
higher reflectivity from the leading hemisphere (Muhleman et al., 1990). Near infrared observations discovered windows in the planet’s spectrum through which the
surface could be glimpsed, again revealing a brighter leading hemisphere (Griffith
et al., 1991; Lemmon et al., 1993). Subsequent studies with the Hubble Space
Telescope and ground-based adaptive optics established the outlines of a continentsized feature that appears to be responsible for this asymmetry in Titan’s surface
albedo (Smith et al., 1996; Combes et al., 1997; Meier et al., 2000). Evidently Titan
does not have a global ocean of hydrocarbons, as initially proposed by Lunine et al.
(1983).
This poses a problem, because the stratospheric production of nitriles, hydrocarbons, and possibly other, more complex organics, with subsequent escape of
hydrogen into space, uses up methane that is not recycled. Thus it is necessary to
replenish this gas, either by outgassing from the interior (volcanism) or by evaporation from some large, subsurface reservoir, or both. The lifetime of the present
atmospheric methane abundance is only 20 × 106 years, indicating the severity of
the problem.
We expect the surface to include deposits of organic aerosols and icy grains of
CO2 and H2 O, all of which will descend through the atmosphere as microscopic
particles. Even though there is not a global ocean, there must be ponds, lakes, or
even seas of liquid hydrocarbons, presumably dominated by ethane and propane,
ATMOSPHERES OF TITAN AND SATURN
357
as these are the most abundant condensible compounds produced by the photochemistry that we know. Thus the surface of Titan contains an important record of
atmospheric processes, and may concentrate some constituents that are extremely
rare in the atmosphere itself.
The properties of Titan’s interior are very poorly understood. Assuming that
Titan consists of differentiated anhydrous chondritic rock and water ice, the density
measured by Voyager (1.8 g cm−3 ) suggests that the satellite has a bulk composition of about 50% rock and 50% ices but the internal structure remains unconstrained.
Models of the interior of Titan have been proposed by Stevenson et al. (1986)
and more recently by Grasset and Sotin (1996). These models include the probable
occurrence of a thick subsurface ocean of H2 O, CH4 and NH3 . It might be that
these species escape, continuously or randomly, up to the surface through cracks
in the crust, and replenish the atmosphere in CH4 (while NH3 , if any, and H2 O
would immediately condense at low surface temperatures). A recent model of the
subnebula of Saturn supports this assumption (Mousis et al., 2002a,b).
2.4.2. Capabilities of Cassini
The radar aboard the Orbiter will be able to map the distribution of bodies of liquid
hydrocarbons on the surface and to define surface topography (Elachi et al., vol.
2). The radar resolution (350–500 meters) will be sufficient to characterize the
surface morphology and cratering record on the solid surface, in those areas of
Titan where close radar passes occur. The VIMS instrument will be able to sense
the surface through the several atmospheric windows that occur in the 0.9–5.0 µm
(micron) spectral region, at a resolution of ∼1 km and the imaging system will
exploit the short wavelength end of this range, reaching scales of ∼15 meters.
The resulting spectrophotometry and imagery will provide additional details on
surface morphology and constrain models for the composition of surface materials,
including local variations caused by weather patterns. (One thinks here of the wind
blown deposits of dark materials on Mars and Triton and the deposition of sediment
by wind and water on Mars and Earth.) The Descent Imager (DISR) on the probe
will achieve still higher resolution over the landing area, again delivering both
images and spectroscopic data. The radar altimeter on the probe will provide local
measurements of surface roughness and reflectivity with high spatial resolution.
Comparing the probe results with the orbiter data should permit interpretations
of the latter on smaller scales than those provided directly by the remote-sensing
instrument. Gravity field measurements derived from the trajectories of the ∼40
scheduled encounters of the orbiter with Titan during the duration of the nominal
four-year mission will permit the development of significant constraints on interior
models of Titan.
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2.4.3. Conditional Measurements If a Soft Landing Occurs
If the probe survives its landing, further refinements are possible by conducting
experiments on the surface. These experiments must be carried out on a best effort
basis; they are not mandated within the mission guidelines, even though there is a
Surface Science Package explicitly included in the probe payload (see Zarenecki
et al., vol. 1). In addition to the experiments included in that package, we point out
the ability of the DISR to obtain images and spectrophotometric data that describe
conditions in the immediate vicinity of the landing site, and measurements with
the GCMS of surface composition that will be enabled by heating the projecting
inlet tube. These compositional measurements may be among the most interesting ones provided by this instrument. A landing in a drift of aerosols will enable
sampling of the most volatile constituents in a concentrated form unachievable
during descent. A wet landing would permit a determination of concentrated photochemical products in the frigid hydrocarbon equivalent of Darwin’s “warm little
pond” that saw the first steps in the chemical evolution that led to life on Earth
(Owen et al., 1992, 1997;Raulin, 1997; Raulin and Owen, vol. 1). A more difficult
scenario would be a landing on an outcrop of surface ice. If the probe survived
and the heated inlet of the GCMS made contact with the ice, it might be possible
to measure a value of D/H in Titan’s ice, an extremely interesting measurement
from the standpoint of the origin and evolution of the satellite’s atmosphere (see
following section).
2.5. O RIGIN AND EVOLUTION OF T ITAN ’ S ATMOSPHERE
2.5.1. Scientific Problem: Determine the origin of N2 , CH4 , and CO in the
atmosphere and the changes in atmospheric composition during Titan’s
lifetime
Voyager observations exclude the possibility that Titan’s atmosphere was formed
by direct retention of gases from the primitive solar nebula or from the Saturn
sub-nebula (including an extended atmosphere of the forming planet). In these
cases, the present Ne/N ratio in the atmosphere should be close to 1, which it is not
(Owen, 1982). The current upper limit on the mixing ratio of Ne set by Voyager
UV observations is 0.002 (Hunten et al., 1984). The idea that Titan’s atmosphere
might have been contributed primarily by cometary impact (Griffith and Zahnle,
1995) is weakened by the observation that D/H in cometary H2 O, the dominant
reservoir of hydrogen in comets, is 3.2 × 10−4 (see, e.g., Meier and Owen, 2000,
for a review of cometary data) whereas in Titan’s methane, it is closer to 10−4
(Coustenis et al., 1998). It is thus reasonable to think that Titan’s atmosphere was
formed predominantly by outgassing of volatiles from planetesimals that formed
the satellite. What were these volatiles?
This is one of the most interesting investigations that Cassini-Huygens will
undertake as the answers have implications not only for Titan but Triton, Pluto
and even the inner planets, as icy planetesimals (in the form of comets) must have
ATMOSPHERES OF TITAN AND SATURN
359
played a role in delivering volatiles to these bodies as well (Owen and Bar-Nun,
2001). In the interstellar cloud from which the solar system formed, approximately
70% of the nitrogen was in the form of N2 with the rest in NH3 and various organic
compounds. The carbon was predominantly in grains, as amorphous carbon and
organic compounds with up to 30% as CO and perhaps 3% as CH4 (van Dishoeck,
1998; Ehrenfreund et al., 1998). Interstellar grains that fell into the nebula disk may
have experienced significant heating due to the protosolar radiation field (Simonelli
et al., 1997) and/or passage through the accretion shock (Cassen and Chick, 1997).
In addition, surviving grains and vapor which entered into the nebula at great heliocentric distances may have been brought to the inner solar system by the systematic
inflow of gas. Some mixing of the two components was an inevitable consequence
of nebular mass and angular momentum transport (Cassen, 1994). Whatever the
composition of the solar nebula at the heliocentric distance of Saturn, the formation
of a subnebula around the forming planet may have resulted in modifications of the
composition since the initial subnebula was inevitably denser and warmer than the
surrounding solar nebula. The problem is to find the link between the composition
of the subnebula and the fact that CH4 and N2 are presently the two major components of the atmosphere of Titan. Under the assumption that the atmosphere was
formed by outgassing from the interior of the satellite, carbon and nitrogen must
have been contained in grains that agglomerated to form Titan, but in what form?
Several scenarios can be envisaged.
In the scenario advocated by Prinn and Fegley (1981), C and N are in the form of
CO and N2 respectively, in the subnebula, while most of the oxygen is in the form
of H2 O. This is based on the assumption that all matter coming from the presolar
cloud has been dissociated during the collapse of the cloud, and subsequently recombined in molecules in the hot inner part of the nebula. These molecules are
transported by turbulence out to the region of formation of Saturn. An appropriate
Saturn subnebula (namely much denser than the nebula) could permit in principle
the conversion of CO into CH4 and the conversion of N2 into NH3 . Subsequently,
when the subnebula cools down, H2 O first condenses, and at lower temperatures
NH3 , followed by CH4 .
Alternatively, as soon as H2 O condenses, water ice can form CH4 clathrates
and NH3 hydrates (Lunine and Stevenson, 1985), or trap CH4 and NH3 through
adsorption if the H2 O ice is in amorphous form. In this last case, the amounts and
proportions of trapped gases are strongly dependent on local temperature (Bar-Nun
et al., 1985). Grains containing CH4 and NH3 could subsequently form a subsurface ocean (see Section 2.4) and ultimately the primitive atmosphere of Titan. The
last step of the process requires the photolytic conversion of NH3 into N2 . To insure
that, the temperature of the early atmosphere of Titan must have been not lower
than 150 K to permit efficient evaporation and photolysis of NH3 to produce the
N2 (Atreya et al., 1978).
The difficulty with this scenario is that the conversion of N2 into NH3 , which
strongly depends on pressure, requires a subnebula that is quite dense, the validity
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of which has never been demonstrated. In fact, more recent models of the subnebula
(e.g., Coradini et al., 1989) are not as dense as that used by Prinn and Fegley
(1981). Accordingly, a second scenario considers that nitrogen was in the form of
N2 in the subnebula. The question is then to explain how N2 was efficiently trapped
in grains that formed Titan. Bar-Nun et al. (1985) and Owen and Bar-Nun (1995)
argue that trapping of this gas by water ice requires very low temperatures. Lunine
and Stevenson (1985) discuss in detail the formation of clathrates of N2 .
A solution to the difficulty of forming NH3 in the subnebula is proposed in a
third scenario (Owen, 2000). The nebula in the region of Saturn is assumed to have
conserved at least a part of the NH3 present in interstellar gas and grains. This
ammonia would then be present in the subnebula and would have subsequently
condensed as in the first scenario. Scenario 3 implies that NH3 came from the outer
part of the nebula without having been converted into N2 . It also implies that very
little conversion of NH3 to N2 occurred in the subnebula prior to the formation of
Titan.
The recent evolutionary model of a turbulent nebula elaborated by Mousis et al.
(2002a,b) supports scenario 3. The authors found that the pressure in the Saturn
subnebula must have been five orders of magnitude lower than that assumed by
Prinn and Fegley (1981). As a result, CO could not have been converted to CH4 or
N2 into NH3 in the subnebula. Moreover, the calculated structure of the subnebula
reveals that the mass of heavy elements was not large enough to form Titan where it
is today. Therefore, Mousis et al. (2002a) argued that the CH4 and NH3 present in
the cooling feeding zone of Saturn around 10 AU were trapped at low temperature
in the form of clathrate hydrate of CH4, and hydrate of NH3, respectively. These
hydrates were incorporated in planetesimals embedded in the cold external part of
the subnebula, which subsequently migrated inwards to form Titan.
How can we choose among these scenarios?
Measurements of noble gas abundances and isotopic ratios provide a means for
discriminating between the primordial and photochemical models. The presence
of a substantial amount of non-radiogenic argon (36 Ar and 38 Ar) and other noble
gases would support the idea that most of the N2 and CO we find in Titan today
were brought in by the ices, either as clathrates or as adsorbed monolayers. On
the other hand, a marked deficiency of primordial argon would support models
suggesting that the nitrogen originated in the form of condensed NH3 or other
compounds. For example, the upper limit of 6% that we already have for Ar in
the present atmosphere (Courtin et al., 1995) coupled with the isotopic evidence
for nitrogen escape (see Section 2.1.1) leads to a value of Ar/N that is much less
than solar. This means that Ar has been fractionated from nitrogen – as it has
on Earth – which could only happen if Titan originally acquired its nitrogen in
the form of compounds other than N2 . If the ice composing Titan contained the
same proportion of NH3 found in comets and in interstellar clouds – about 1% –
outgassing could have produced an atmosphere of ∼120 bars of N2 , so there is no
shortage of nitrogen with this model!
ATMOSPHERES OF TITAN AND SATURN
361
The isotope ratios will provide further clues. We need to know how the fractionation and escape of N and H has taken place in the atmosphere, leading to the
values we find today. Did the methane we find on Titan form in the subnebula?
Was it included in the ice like NH3 ? Or did it form during accretion, on the satellite
itself? The answer may be found through a study of D/H in Titan’s H2 O, which
would be the source of the hydrogen for in-situ methane production. Studies of D/H
in other compounds besides CH4 will help determine the extent of photochemical
fractionation of these important isotopes. Mousis et al. (2002b) argued that the D/H
ratio measured in Titan is representative of that in CH4 in the cooling solar nebula
at 10 AU, at the time when methane was trapped in the form of clathrate hydrates.
This scenario constrains the value of D/H in methane ices falling from the presolar
cloud into the early nebula.
2.5.2. Capabilities of Cassini
The GCMS aboard the Titan Probe will measure noble gas abundances and their
isotopic ratios. It will also determine isotope ratios for the other abundant elements.
The use of the gas chromatograph in combination with the mass spectrometer
allows a clean discrimination between molecules and their fragments that would
otherwise have overlapping charge-to-mass ratios. If the probe survives the landing, the concentration of condensed species on the surface will provide an additional increase in sensitivity. One issue to be considered then will be the possible
confusion between ice crystals coming into the atmosphere from the outside and
the ice that formed the satellite. The former may include cometary ice with a much
higher D/H than is expected on Titan. Finding evidence of both types of ice would
be extremely interesting.
3. Saturn
3.1. I NTRODUCTION
Despite the apparent similarities between the two largest planets in the solar system, the Voyager encounters have revealed that Saturn is significantly different
from Jupiter. The study of each planet will clearly benefit from comparative observations of the other. Following the Galileo mission to Jupiter in 1995–2000,
Cassini will allow us a much deeper level of comparison between the two largest
planets in our solar system. The reasons for apparent differences in internal structure and general circulation for example, can be studied in ways beyond the reach
of the Voyager spacecraft. The potential for advancing our knowledge is further
enhanced by the fact that Cassini will arrive at Saturn during a different season from
that of the Voyager encounters. Furthermore, the Cassini flyby of Jupiter (closest
encounter on 30 December 2000), en route to Saturn, allows a direct comparison
of observations of the two planets with the same suite of instruments, calibrated at
Jupiter by results from the Galileo entry Probe.
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3.2. T HE INTERIOR
3.2.1. Scientific Problem: Constrain the Extent of Interior Structure
Differentiation and Its Coupling to the Atmospheric Circulation
Voyager observations of Saturn revealed an outer envelope depleted in helium with
respect to a protosolar composition mixture and a strong internal heat source that
is larger than can be accounted for by homogeneous cooling models. Assuming a
methane-to-molecular hydrogen mixing ratio of 4 × 10−3 , Conrath et al. (1984)
found 0.06 ± 0.05 for the He mass fraction, compared with 0.238 ± 0.007 for the
He mass fraction on Jupiter, as determined by instruments on the Galileo probe
(von Zahn et al., 1998; Niemann et al., 1998). However, the Voyager observations
of Jupiter led to a lower value of 0.18 ± 0.04 for the He mass fraction on Jupiter,
suggesting the presence of a systematic error in the Voyager results, which would
affect the Saturn determination as well.
Saturn Voyager data have been reexamined by Conrath and Gautier (2000).
Using a method of retrieval of the H2 /He ratio based on the differential spectral
behavior of H2 –H2 and H2 –He collision-induced absorption coefficients (Gautier
and Grossman, 1972), these authors suggest that the mass fraction of He in the outer
envelope of Saturn is about three times higher than the value previously inferred by
Conrath et al. (1984) which was dependent on the radio occultation profile obtained
by the Voyager spacecraft.
The occurrence of internal heat in Saturn, in spite of the fact that the planet
should have lost the energy initially acquired at the time of its formation about
two billion years ago is interpreted as resulting from the demixing of helium from
metallic hydrogen in the deep interior of Saturn (Stevenson and Salpeter, 1977;
Guillot, 1999; Hubbard et al., 1999). This separation which began late in the history of Saturn and the subsequent migration of helium droplets towards its center
liberates gravitational energy responsible for the internal energy observed today
(see Figure 2).
However, this internal energy depends upon the amount of helium that migrated
from the outer envelope downwards toward the center of the planet since the formation of this object. Interestingly enough, two recent models of evolution of Saturn
by Guillot (1999) and by Hubbard et al. (1999) both conclude that the abundance
of helium in the envelope should be substantially higher than that derived by Conrath et al. (1984). The factor 3 increase deduced by Conrath and Gautier (2000)
from their reanalysis of the Voyager data is consistent with these new evolutionary
models.
3.2.2. Capabilities of Cassini
Saturn’s gravitational field will be measured throughout the Cassini orbital mission
by radio tracking of the Orbiter trajectory at a variety of inclinations (Kliore et al.,
vol. 2). The radius of Saturn at low and high latitudes will be determined by a series
of radio occultation measurements. These improvements in our knowledge of the
ATMOSPHERES OF TITAN AND SATURN
363
shape and gravitational field of the planet will provide further constraints on models for Saturn’s interior. The far infrared spectrometer will retrieve temperatures
and their variation with latitude at the top of Saturn’s convection zone. Combined
with radio-occultation retrievals, or independently, the infrared measurements will
deliver a more accurate determination of the helium abundance. Thermal emission
measurements by the Composite Infrared Spectrometer (CIRS) and observations
of reflected solar energy by the near infrared spectrometer will be used to refine
our knowledge of the global energy balance.
3.3. T HERMAL STRUCTURE AND COMPOSITION
3.3.1. Scientific Problem: Determine tropospheric and stratospheric
temperatures and their vertical variation at all latitudes. study seasonal
variations of temperature. Determine ortho-para ratios of molecular
hydrogen at various locations on the planetary disk. Infer vertical
distributions of convection tracers. Infer vertical and latitudinal
distributions of hydrocarbons in the stratosphere. Measure the helium
abundance and D/H ratio. Search for new atmospheric constituents
Voyager IRIS and Radio Science measurements have provided retrievals of vertical
temperature structure at pressure levels between 1 and 1000 mbar. The Voyager
IRIS data revealed latitudinal gradients in temperature in Saturn’s upper troposphere that are strongly correlated with the cloud-tracked winds, but could not be
measured at higher levels. A north-south hemispheric asymmetry of temperature
was also observed, indicating a seasonal response (at a time near the equinox)
with moderate thermal inertia. The altitude range of the retrieval of tropospheric
profiles from Voyager IRIS data was limited by the lack of measurements below
200 wavenumbers (λ > 50 µm). Ground-based microwave observations also show
strong latitudinal variations in tropospheric temperatures at pressure levels of a
few bars, including a pronounced warm band at northern mid-latitudes. These
are difficult to interpret as kinetic temperatures owing to uncertainties in the NH3
abundance, but may be diagnostic of vertical motions. IRIS determinations of the
ortho-para ratio in Saturn’s upper troposphere have provided evidence for variations of this ratio with latitude, as recently demonstrated by Conrath et al. (1998).
IRIS determinations of the CH4 abundance in the stratosphere indicate an enrichment of carbon amounting to approximately four-and-one-half times the solar
abundance value (Courtin et al., 1984), which was supported by ground-based
observations that indicate an enrichment of 4 ± 1 × solar (Karkoschka and Tomasko, 1992). However, more recently, Kerola et al. (1997) found an enrichment
of 3 + / − 1× solar. It is an important issue, as models for the origin of giant
planet atmospheres predict that abundances of heavy elements should be greater on
Saturn than on Jupiter (e.g., Pollack and Bodenheimer, 1989; Owen and Bar-Nun,
1995). However, Hersant et al. (2003) propose a scenario in which CH4 and NH3
were trapped by clathration and hydration in the feeding zone of Saturn, thereby
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producing in the planet a carbon enrichment of 2.5, and a nitrogen enrichment of
2.0. Both values are less than those observed in Jupiter, in which Ar, Kr, Xe, C,
N, S are all enriched by a factor 3 ± 1 compared to solar abundances tabulated by
Anders and Grevesse (1989). What are thus the abundances of species other than
methane observed in Saturn? (Owen et al., 1999). What is the enrichment of NH3
on Saturn? Limited information on the deep tropospheric abundance of NH3 on
Saturn has been obtained from ground-based microwave measurements as well as
from IRIS spectra. Indications of an NH3 abundance near the saturation limiting
value at 500–700 mbar were provided by Voyager radio occultation data. Various
hydrocarbon constituents, such as C2 H2 , C2 H6 , C3 H4 , and C3 H8 have been detected
in Saturn’s stratosphere, but with only moderate spectral resolution and signalto-noise ratio, which has prevented an accurate assessment of their abundances.
(reviewed by Atreya [1986], Noll and Larson [1991] and Atreya et al., 2003). There
are indications of significantly higher abundances of some constituents on Saturn
that are also present on Jupiter, especially PH3 which may be enhanced by an order
of magnitude (Noll and Larson, 1991). ISO has found CO2 , H2 O and CH3 in the
stratosphere (Encrenaz et al., 1999). Several other as yet undetected molecules
are expected to be present in Saturn’s atmosphere. If detected they would offer
important clues to Saturn’s atmospheric chemistry.
It is widely recognized that an accurate determination of D/H can provide useful
constraints on models for the origin and evolution of planetary atmospheres. As H2
dominates all other hydrogen-bearing constituents by a large margin on Jupiter
and Saturn, the value of D/H in the H2 in these atmospheres should be almost
identical to the value in the hydrogen that dominated the interstellar cloud from
which the solar system formed. Neither of these giant planets is sufficiently massive
to produce internal temperatures and pressures that would permit nuclear reactions
converting D to 3 He. The only disturbance of the original ratio would therefore
come from other hydrogen compounds that were brought to the forming planet as
solids (primarily in ices) in sufficient amounts.
The value of D/H on Saturn was given as D/H = (1.7 + 1.9/ − 1.0) × 10−5
by Gautier and Owen (1989) from an analysis of observations of bands of CH4
and CH3 D (with correction for fractionation from hydrogen) by several different
observers. More recently, Griffin et al. (1996) have reported D/H = (2.3 + 1.2 −
0.8) × 10−5 from observations of HD absorption lines by the Infrared Space Observatory (ISO). From the short wavelength spectrometer (SWS) of ISO Lellouch
et al. (2001) derived D/H = 1.85 + 0.85/ − 0.6) × 10(−5) from HD, and a somewhat lower value from CH3HD. Using Jupiter for comparison, we point out that
Mahaffy et al. (1998) derived D/H (2.6 ± 0.7) × 10−5 from direct measurements
of HD and H2 in Jupiter’s atmosphere. This is significantly higher than the value of
1.6 ± 0.2 × 10−5 found in local interstellar hydrogen today (Linsky, 1996), as expected from models for galactic evolution that predict the destruction of deuterium
with time as a result of nuclear “burning” in stars.
ATMOSPHERES OF TITAN AND SATURN
365
Figure 2. Conventional views of the interiors of Jupiter and Saturn, with a three-layer structure
including the ice/rock core, the molecular and the metallic regions. A transition zone, assumed to lie
between 1 and 3 mbar, is represented but experiments are still unclear as to whether the separation
between the molecular and metallic regions is sharp or not. The planets are shown to scale, with their
observed oblateness and obliquity (from Guillot, 1999).
Within the errors of the current D/H measurements, Jupiter and Saturn may
exhibit the same value of D/H, in spite of the fact that Saturn appears to contain
more heavy elements, and thus presumably more ices than Jupiter (Guillot et al.,
1994). However, under the assumption that volatiles were trapped in the form of
clathrate hydrates in the feeding zones of both Jupiter and Saturn, Hersant et al.
(2003) predict that the enrichment in oxygen in Saturn should be almost twice less
than in Jupiter (Gautier el. 2002). Unfortunately, the O/H ratio is still unknown in
both planets. Obviously it is very important to increase the precision of the D/H
measurement on Saturn in order to decide whether or not the hydrogen reservoir in
the planet has been enriched in deuterium from the condensed matter (predomin-
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antly in the form of icy planetesimals) that formed the core and contributed excess
heavy elements to the atmosphere. Coupled with more accurate determinations of
the abundances of carbon, nitrogen and phosphorous, the deuterium measurement
can then be used to estimate the mass of heavy elements contained in the planets
and to characterize the composition of the planetesimals that delivered them.
3.3.2. Capabilities of Cassini
Spectral measurements from 10 to 1400 cm−1 (7 to 1000 µm) will permit the
determination of atmospheric composition and thermal structure between 100 and
1000 mbar as well as between 0.1 and 10 mbar. The planned 2004 arrival date
will permit the measurement of atmospheric conditions near the solstice, providing
an important seasonal comparison with Voyager observations. Good horizontal
mapping and some vertical information on the ortho-para hydrogen ratio will be
available from CIRS, partly by assessment of the relative strength of H2 –H2 and
H2 -He collision – induced absorption lines. The hydrogen dimer (H2 )2 , detected
by Voyager, will also be spectrally analyzed at 16 and 28 µm, providing a very
sensitive evaluation of its latitudinal variation. These determinations will significantly constrain the adiabatic lapse rate at deeper levels and also provide useful
diagnostics of vertical motions. Latitudinal variations of methane emission will be
measured in limb-sensing mode at levels up to 1 microbar, which will yield the
kinetic temperature with a vertical resolution of one scale height. Vertical and horizontal distributions of C2 H2 , C2 H6 , and other hydrocarbons can also be obtained
by limb sounding in the middle infrared.
Radio occultation measurements will provide an independent retrieval of thermal
profiles at selected locations, permitting a determination of He/H2 in conjunction
with the infrared measurements. This ratio will also be directly retrieved by CIRS
from IR measurements at long wavelengths.
Ammonia (NH3 ) cloud optical depths will be mapped using spectral regions
near 200 and 1000 wavenumbers and in the near infrared. The near infrared spectrometer will yield further important constraints on cloud properties including particle sizes and relative heights of different cloud layers.
Making many measurements of additional HD lines over four years at several
locations on the disk of the planet (thereby permitting an improved evaluation
of the cloud opacity), CIRS will easily be able to surpass the accuracy of the
ISO determination. Improved modeling and calibration will further enhance this
effort. In addition, D/H in methane can also be inferred from measurements of the
CH3 D/CH4 ratio. Comparing this D/H with D/H in H2 may permit us to evaluate the
isotopic enrichment factor in methane in the upper troposphere of Saturn (Lecluse
et al., 1996). A precise determination of this factor would help us to discriminate
among the various models of dynamics of the deep atmosphere, models that are
presently only poorly constrained.
Information on the thermal structure and composition of the upper atmosphere
will be acquired by observing the sun or a star at the planetary limb with the
ATMOSPHERES OF TITAN AND SATURN
367
ultraviolet spectrometer. The high-speed photometer will measure high-altitude
atmospheric densities. Soundings of the deep atmosphere from 2–10 bar will be
obtained by microwave radiometry. The opportunity for variable viewing geometry
and, by use of the telecommunications dish, for high spatial-resolution observations
will aid in separating kinetic temperature and variable abundance effects.
3.4. ATMOSPHERIC DYNAMICS AND GENERAL CIRCULATION
3.4.1. Scientific Problem: Perform cloud-tracked wind and correlated
temperature measurements at all latitudes with sufficient time and
longitude coverage to analyze eddy-mean flow exchanges. Constrain
global scale temperature-density gradients at sub-cloud levels. Measure
vertical wave propagation at selected latitudes
Voyager imaging observations of large-scale motions on Saturn revealed a superrotating equatorial jet and at higher latitudes an axisymmetric and apparently longlived pattern of counter-flowing jet streams. Although these features of the atmospheric circulation are qualitatively similar to observed motions on Jupiter, the
Saturn jet streams are a few times stronger, somewhat wider, and more dominantly
prograde, measured with respect to the planet’s rotation period. Two extreme model
interpretations have been suggested for the general circulation of both giant planets.
In one view the observed jet streams are the manifestation of counter rotating
cylinders of convection, concentric with the planetary spin axis, and extending
deep into the interior of the molecular hydrogen envelope. In the other view the
observed motions are confined to a relatively shallow layer extending no more
than a few scale heights below the cloud tops and supported by strong latitudinal
gradients in temperature and composition (Figure 3). Further measurements will
be required to distinguish between the two models. The Voyager observations of
Saturn’s high wind speeds and its enormous equatorial jet provide an even greater
challenge to theoretical understanding than the observations for Jupiter and have
served to emphasize the distinction between the two extreme models illustrated in
Figure 3. The diagnostic assessment of wave – mean flow exchange and cloud –
temperature correlations are also more problematic for Saturn. Present knowledge
cannot provide definitive answers to the most fundamental questions about the
nature of the Saturn meteorology. How deep do the jet streams extend into the
interior? How does the deep internal convection couple to the motions at the cloud
tops? Do small scale horizontal eddies feed the large-scale jets or do the jets support
the eddies? What is the role of moist latent heat-release and ortho-para hydrogen
conversion in the global circulation balance?
Progress in answering these questions will require more elaborate measurements extended to longer time frames and other vertical levels of the atmosphere
than were available from the flyby Voyager encounter. Voyager infrared measurements of north-south gradients in temperature just above the cloud tops imply a
reduction of the wind speeds with altitude toward a net west-to-east directed flow.
OWEN AND GAUTIER
Wea
the
r la
ye
368
r
Cloud
top winds
Liquid
"metal"
core
Free
convection
(a) Shallow weather layer
Cloud
top winds
Cylinders
Liquid
"metal"
core
Fluid
envelope
(b) Deep convective cylinders
Figure 3. Two alternative models for the general circulation of Saturn’s atmosphere: (a) suggests that
the circulation is driven in a rather shallow layer in the visible part of the atmosphere; (b) invokes
deep convection following a pattern of concentric cylinders. (After Michael Allison.)
ATMOSPHERES OF TITAN AND SATURN
369
The extension of these measurements to higher levels of the upper stratosphere
will offer important clues to the circulation balance at deeper levels. Long wave
microwave sounding combined with radio occultation measurements of vertical
temperature profiles will place important constraints on the horizontal gradients in
temperature and composition. Measurements of the ortho-para hydrogen ratio and
its variation with latitude will provide important information on the thermodynamic
state of the wind layer and will be diagnostic of the vertical motion field. Measurements of zonal mean and eddy scale cloud velocities over several months of
observations will characterize the zonal momentum balances. Studies of the motion
of small-scale cloud features and waves will serve as an additional diagnostic of
the wind structure at deeper levels.
3.4.2. Capabilities of Cassini
CIRS will provide horizontal temperature gradients in the upper troposphere and
stratosphere that can be used in thermal wind analyses of the Saturn jet streams.
This will permit characterization of the zonal wind field in both the troposphere
and stratosphere, including inference of the vertical jet decay at high levels. Zonal
thermal structure will be used to analyze tropospheric waves, while a combination
of zonal structure and vertical structure with a resolution better than one scale
height obtained from limb sounding and radio occultations will provide information on stratospheric wave propagation. These results, combined with a determination of possible latitudinal variations of hydrocarbon and ammonia abundances, can
be used to define stratospheric circulation and transport. Extended visible imaging
observations during dayside petal orbits of the Saturn disk will provide measurements of zonal winds in the upper troposphere and the localized motion of
atmospheric waves and vortices (Porco et al., vol. 2).
Radio occultation measurements will provide data on the figure of Saturn, which,
combined with improved determinations of the gravitational moments, will place
important constraints on the rotational state of Saturn’s interior.
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3.5. S ATURN ’ S IONOSPHERE
3.5.1. Scientific Problem: Determine from radio occultations the diurnal
variation of ionization. Establish the role of plasma transport as well as
chemical processes associated with infalling of H2 O
Radio occultation observations from Pioneer 10 and Voyager 1 and 2 have established the existence of an ionosphere with a peak density of ∼ 2 × 104 cm−3 . In
addition, the Voyager Planetary Radio Astronomy experiment (PRA) has provided
an indirect measure of the diurnal variation of the peak ionospheric density through
its observations of Saturn Electrostatic Discharges (SED), thought to originate
from lightning in the lower atmosphere. The SED frequencies (which necessarily
exceed the electron plasma frequency) also constrain the allowable diurnal variation to a maximum of greater than 105 cm−3 around local noon and a minimum of
less than 103 cm−3 near dusk. Although these variations are suggestive of ion pair
production by solar extreme ultraviolet radiation, there is an indication of rising
electron densities during the night that is inconsistent with a simple photochemical
explanation. It has been suggested that ion loss processes arising from the conversion of H+ to H2 O+ and H3 O+ , may be driven by an injection of H2 O from the
icy rings. The possible role of plasma transport processes in controlling the height
of the ionospheric peak also remains an unsolved problem. The recent discovery
by ISO of CO2 in addition to H2 O in Saturn’s upper atmospheric (Encrenaz et al.,
1998) strengthens this interpretation and adds another molecule to be considered
in this analysis.
3.5.2. Capabilities of Cassini
Repeated radio occultations of Saturn by the Cassini Orbiter are expected to provide
a wealth of new ionospheric data. Ion chemistry (possibly including H2 O transport)
can be studied by the UV spectrometer. CIRS will be able to map the destination of
incoming H2 O or CO2 . SED’s containing ionospheric information will be observed
by the plasma/radio wave experiment (Gurnett et al., vol. 2). Measurements of the
polar wind and possible field-aligned flows from the ionosphere by the plasma
instrument can provide direct information about the ion composition of Saturn’s
ionosphere at high latitudes. Taken together, these data will contribute substantially
to understanding the origin and variability of Saturn’s ionosphere.
4. Summary
It should be evident from this review that the capabilities of the spacecraft now on
its way to Saturn and Titan are extraordinarily powerful and diverse. Additional
details about the instruments on the probe, their scientific objectives and more
background science for Titan can be found in the ESA Special Report (SP-1177)
HUYGENS: Science, Payload and Mission (ed. J. P. Lebreton, ESA Publications,
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371
TABLE I
Organic Compounds Expected to Be Present in Titan’s Atmosphere
Compounds
Expected Mean Mole Fraction
Hydrocarbons
10−7 –10−8
C3 H6 (propene) a
a
C3 H4 (allene)
10−7 –10−8
C3 H6 (cyclopropane) a
10−8 –10−10
a
C6 H2 (triacetylene)
10−8 –10−10
a
C8 H2 (tetraacetylene)
10−9 –10−10
C 4 H4
10−8 –10−10
1,2–C4 H6
10−8 –10−10
1,3–C4 H6
10−8 –10−10
C 4 H8
10−8 –10−10
n C4 H10
10−7 –10−9
iso–C4 H10
10−7 –10−9
C 6 H6
< about 10−9
(related to C2 H2 polymerization)
Nitriles
C2 H5 –CN
CH2 = CH–CN
CH3 –CC–CN
n–C3 H7 –CN
iso–C3 H7 –CN
cyclo–C3 H6 –CN
HC5 N
< 5 × 10−7 b
< 2 × 10−7 b
< 2.5 × 10−7 b
< 6 × 10−8 c
< 2 × 10−8 c
< 1.5 × 10−8 c
< 1.5 × 10−8 c
O-Compounds
HCO
H2 O
CH3 OH
10−10 –1012
10−10 –1012
10−10 –1012
a May play an important role in Titan’s photochemistry and
needs to be studied.
b Likely to be in the range 10−8 –10−9 .
c Likely to be in the range 10−8 –10−10 .
Noordwijk, The Netherlands, 1997) 364 pp. While the early hopes for a joint
US/European mission to the Saturn System have been fully realized (Owen et al.,
1983; Gautier and Ip, 1984; Owen et al., 1986), we regret that the planned follow-
372
OWEN AND GAUTIER
up of a Saturn atmospheric probe (Owen et al., 1983, 1986; Swenson et al., 1987)
is not in any current mission plan. Experience from the Jupiter probe has stressed
the necessity of in-situ experiments if we are to achieve the level of precision
required for meaningful comparative studies of the giant planets. In fact, it is now
clear that we need missions with multiple probes to sample these atmospheres, so
we can achieve a proper global picture. Multi-probe missions to all four giants,
including Uranus and Neptune orbiters, are thus the logical next steps to follow
Cassini-Huygens. Cooperation among the world’s space-faring nations will make
such ambitious missions affordable, just as it has for Cassini-Huygens.
Acknowledgements
We thank our co-leader, Wing Ip and the other members of the Joint Science
Working Group for the Cassini-Huygens mission, all of whom participated in the
preparation of the science sections in the original Phase A Study Report: M. Allison, S. Bauer, M. Blanc, S. Calcutt, J. Cuzzi, M. Fulchignoni, D. Hunten, T.
Johnson, H. Masursky, P. Nicholson, R. Samuelson, F. Scarf, E. Sittler and B.
Swenson.
Appendix
Table I lists organic compounds not yet detected in Titan’s atmosphere but expected to be present at concentrations detectable from Cassini Huygens (Gautier and
Raulin, 1997).
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