Titan: Recent developments

Titan: Recent developments

vistas in Astronomy, Vol. 34, pp. 11-50, 1991 Printed in Great Britain. All rights reserved. 0083--6656/91 $0.G0 + .50 ~) 1991 Pergamon Press plc. T...

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vistas in Astronomy, Vol. 34, pp. 11-50, 1991 Printed in Great Britain. All rights reserved.

0083--6656/91 $0.G0 + .50 ~) 1991 Pergamon Press plc.

TITAN: RECENT DEVELOPMENTS Athena Coustenis D6partement de Recherches Spatiales, Observatoire de Paris (Section de Meudon), 92195 Meudon, France

Abstract Considerable and rapid progress has recently been achieved in Titan science, as a result of Voyager data analyses and ground-based observations leading to more complete and precise modeling of the satellite's atmosphere. Valuable sdentific information has thus been added to our previous knowledge of Titan. In particular, data sets unexploited up to our days were investigated, new models developed to account for these observations, theories extended and controversed or sustained, and measurements re-analysed. The aim of this paper is to give a brief overview of our current understanding of the sateURe's nature and composition. In the interest that has risen in view of future space missions to the saturnian system, such a synthesis could be useful in better making apparent future assignements. After a brief survey of the pre-Voyager rather rudimentary view of Titan's atmosphere and of the post-encounter first-years period -which submerged the scientific community with large amounts of data-, this study will concentrate on the Titan science developments that have occurred in the past five years or so.

Address for correspondance: Dr Athena Coustenis DESPA (Bat. LAM) Observatoire de Meudon 92195 MEUDON Cedex FRANCE

J I~*'A 34 s I 12-11


A. Coustenis




Italian, French and English astronomers are the people to whom we owe the bulk of satellite discoveries from 1610 to 1851. Galileo observed the four biggest jovian satellites (Io, Europa, Ganymede and CaUisto), while Herschel and Cassini shared the satellites of Saturn. Those of Uranus were, in their majority, left to the observational skill of Lassell. Yet, the one satellite which was for many years believed to be the biggest one (until Voyager measured Ganymede's radius with precision and found it to be the largest in our solar system), was discovered on March 25 th 1655 by a novice Dutch astronomer: Christiaan Huygens. This satellite has certainly deserved its important name of Titan (on Herschel's proposal in 1847), not only because of its great size (comparable to that of Ganymede and Callisto), but also because of another characteristic: its substantial atmosphere similar to the Earth's. This property, at present not found elsewhere, remained hidden until almost 300 years after Titan's discovery.


Pre-Voyager knowledge

Before the Voyager missions, ground-based observations of the satellites had brought very scarce information, so much so that few of them -as registered today- were known and information lacked on the ones already discovered. With respect to Titan, the main scientific concern was focused on its atmosphere, identified and found to contain methane first by Kuiper (1944).

From 1972 to 1979 several

scientists (Trafton, Lutz, Owen, Cess, Fink, Larson and others) concentrated their efforts on an estimation of the methane abundance and of the pressure conditions in the atmosphere from observations made in the -,- 1 to -,- 2 micron spectral region. The results varied a great deal because of the uncertainties and the pre-Voyager models did not agree on the main composition or on the temperature/pressure conditions in Titan's atmosphere. However, there were two principal models: the first one favored methane as the main component and predicted surface conditions of T=86 K for 20 mbar, as well as a temperature inversion in the higher atmospheric levels illustrated by the presence of emission features of gases in the infrared spectrum of Titan (Danidson, 1973; Caldwell, 1977); the second one (Hunten, 1977) was based on the assumption that ammoniac dissociation should produce molecular nitrogen (transparent in the visible and infrared spectrum) in large quantities and held that the surface temperature and pressure would be quite high (200 K for 20 bars).

Titan: Recent Developments


On the other hand, albedo measurements suggested the presence of aerosols in appreciable amounts in form essentially of CH4 clouds and haze (Veverka, 1973). Some of the absorbers expected to be formed by methane photolysis (ethane, acetylene, ethylene and mono-deuterated methane) were already identified in spectra obtained mainly by Oillett (1975-1977). For a more detailed description of the pre-Voyager understanding of Titan's atmosphere and surface, see Hunten (1977) in "Planetary Satellites" (J.A. Burns, ed., Univ. of Arizona Press). 1.2




The Voyager 1 encounter with Titan occurred on November 12 *h 1980 at a closest approach distance of --~ 4000 km from the sateUite's surface. Voyager 2 followed nine months later passing a hundred times further away. The bulk of the information we possess on Titan through space missions is therefore due to Voyager 1. A large number of review papers, as well as several chapters in books related to the saturnian system are devoted to Saturn's satellites. The immediate post-encounter analyses of the data recorded by the various Voyager instruments are included in special issues of Science (1981: 212, pp. 159-243; 1982: 215, pp. 499-594) and of Nature (1981: 292, pp. 675-755). Also, there are review papers on Titan, which describe and compare Voyager and pre-Voyager results (e.g. Owen, 1982; Strobel, 1982 etc.). There follows a brief overview of the information concerning Titan gained from Voyager and ground-based observations in the first few years that foUowed the encounter. The answer to the question as to Titan's nature came finally as a combination of the two pre-Voyager models. Nitrogen (detected by the UV spectrometer) is by far the major component of the atmosphere (.,~ 90%), with methane as the next most abundant molecule (2-8 %) and traces of H2 and other more complex organic gases; a temperature inversion occurs near the tropopanse (--~ 42 km) which manifests itself in the infrared spectra in the presence of emission bands of the minor components; finally, the surface conditions are 94 + 2 K for a pressure of ~ 1.5 bar (Tyler et al., 1981; Hans1 e t a l . , 1981). The radio-occultation experiment provided density profiles as a function of altitude near the equator, afterwards converted (Lindai et al., 1983) into temperature/mean molecular weight vertical profiles in the atmosphere (see 1.3.1). A precise value of Titan's radius was derived: 2575 + 2 kin. The main properties of Titan as established at the time are listed in Table I from Hunten (1984).

A. Coustenis


Table I

Properties of T i t a n Surface radius.' R-r Mass. M GM

Surface gravity, gs Mean density Rock/ice ratio (by mass) Distance from Saturn from Sun Orbit period around Sun Obliquity Temperature. K ' Su~e ~t'ective Tropopause (42 kin) Stratopau~e (200 kin) Exobase (1600 km) Bond albedo Solar flux Surface pressure

2575 km 1.346 x 1026g ffi 0.022 x Earth 8.976 x l0 la cm ~ s--" 135 cm s -2 1.881 g cm -3 - 52:48 1.226 x 106 km ffi 20 9.546 A U 15.95 day 30 yr 267 (assumed)

86 71.4 170 186 - 20 0.20 1.1% x Earth 1496 *-. 20 mbar

a Atmospheric temperatures below 200 km assume mean molecular weight 28 and scale with this value (Lindal et al. . 1983). Table from Hunten et al. (1984). Under the surface conditions mentioned above, liquid methane is expected to exist and could even form an ocean, while in the troposphere, methane clouds (formed by saturation of methane gas) might cause rains which would resupply the ocean. However, the saturation degree in the lower atmosphere is unknown.

The CH4 abundance is therefore difficult to

determine. Samuelson e t aL (1981) suggested a mea~ value of 3% above the condensation level and up to (]% below. The mean molecular weight value is around 28 ainu sustaining the N2 majority, but m a y be as high as 29.4 ainu, implying the presence of a heavier component than nitrogen in Titan's atmosphere. On cosmological basis, the presence of argon was suggested (Owen, 1982; Strobel, 1982), but was not detected which suggests that its abundance must be lower than 6% in the higher atmosphere: Broadfoot et al. (1981). The images taken by the cameras aboard Voyager 1 showed that Titan was completely covered by thick haze which allows no visibility of the surface. Also a pole-to-pole asymmetry

Titan: Recent Developments


(due probably to compositional/seasonal reasons) appeared, further confirming the presence of aerosols in the atmosphere. To our current knowledge, the aerosols axe mainly produced in the upper atmosphere through methane-nitrogen photochemistry by UV photolysis or dectron-impact reactions. They axe believed to form polymers diffusing downwards in the 50-100 km region where the lowest temperatures occur. The hydrocarbons and other complex molecules then condense and become solid suspended particles responsible for Titan's characteristic orange color. These particles slowly descend further down, collide with each other, develop into aggregates, which fall rapidly, and finaUy axe deposited on the surface. The infrared spectrometer (IRIS) aboard Voyager 1 also studied the structure of Titan's atmosphere. The infrared spectra cover the 200 to 1500 cm -1 spectral region with a resolution of 4.3 cm -1 and at relatively high spatial resolution. These spectra confirmed the presence of nitrogen as the main constituent and of other minor components, and allowed the identification of several other gases not yet detected. Thus, the presence of CH4, C~H2, C2H4 and C2Hs in Titan's atmosphere was confirmed (Hanel et al., 1981), while in the infrared spectra the signatures of HCN (important prebiotic molecule) and two other nitriles (HCsN and C2N2: Kunde et al., 1981) were found. Also idewtified by comparison with laboratory spectra were some hydrocarbons whose abundances were estimated: C4Ha (Kunde et al., 1981), C3H4 and C3Hs (Maguire et al., 1981) and CH3D (Kim and Caldwetl, 1982). Finally, Samuelson et al. (1983) found carbon dioxide at 667 cm -~ and determined its column abundance.




During the rush that usually follows space mission encounters scientists are set to work in narrow time-schedules to extract as much information as possible from masses of data in order to inform a public eager for knowledge. Once this is done, long-term thorough analyses are undertaken, demanding improved laboratory measurements and involving interaction with other studies. Our understanding of Titan was described twice in books that appeared in 1984 and 1986 and which contain a succint but complete overview of the information that was available at the time: Part VI of "Saturn", (Gehrels and Matthews Eds., 1984) pp. 671 to 759 by Hunten et al., is devoted to Titan. In a narrower frame, Titan is discussed in "Satellites" (Burns and Matthews, Eds., 1986), pp. 786 to 794, in a chapter by Morrison et al. which deals with Saturn's satellites (whose number almost doubled since the 1980 discoveries). The following subsections serve as an introduction to Sections 2 and 3. They contain a summary of some of the knowledge acquired on Titan during the first half of the 1980s.


A. Coustenis


Atmospheric composition

Prior to its detection, the presence of CO in Titan's atmosphere was advocated both on cosmogonical reasons and because of the identification of CO2 in the IRIS spectra and the attendant photochemical implications. CO was discovered by Lutz et al. (1983) from groundbased observations in the near infrared and its abundance was estimated to be ,-~ 1 x 10 -4 in agreement with photochemical model predictions. For those gases which were either observed directly or inferred to be in the atmosphere of Titan, relative abundances were tabulated (see e.g. Table IV in Morisson et al. , 1986). Some of them (N2, Ar, CO and H~) are expected to be uniformly mixed throughout the lower atmosphere without undergoing phase changes.

Methane should also be uniformly

mixed globally, but since it may undergo phase changes in the troposphere, it may not be vertically uniform. Thus, its mixing ratio is supposed to be constant above the condensation level, below which methane is saturated and may form clouds. Such clouds seem necessary to account for the observed opacity between 200 and 400 cm -1. The methane abundance can therefore increase with temperature down to the surface. After methane, the next most abundant hydrocarbons are ethane, acetylene and propane. Each of these organics and the less abundant ones must condense at some level in the lower stratosphere and precipitate out, effectively reducing its gas phase abundance to smaller amounts below this level. Although no systematic quantitative study was performed at that stage, latitudinal variations were observed for the least abundant molecules, while the most abundant gases appeared to be approximately constant with latitude. HCsN and C2N2 could not be identified below -,~ 60°N and seasonal effects were suggested as the origin of this variation. 1.3.2

Thermal structure

a) Troposphere-stratosphere Following a concise analysis of the ingress radio-occultation data by Tyler et al. (1981), the measurements made by this experiment (RSS) during the ingress and egress of Voyager 1 near Titan's equatorial region were extensively exploited by Lindal et al. (1983). They derived information on the troposphere, stratosphere and ionosphere of Titan. Two vertical refractivity profiles were generated and converted, under the hypothesis of a pure N~ atmosphere, into vertical density distributions. The latter were then used to calculate vertical pressure profiles and finally to obtain temperature profiles as a function of altitude (0 to 200 kin) in Titan's equatorial region. In the troposphere the temperature is bound between the surface value and a minimum

Titan: Recent Developments


of -~ 71 K at the tropopause. It therefore remains lower than the condensation temperature of N2: it is thus improbable that nitrogen clouds could form. On the other hand, condensation of methane is possible if the CH4 stratospheric mixing ratio exceeds 1.6%.

The surface

temperature and the tropospheric composition influence the properties of an ocean on the surface. The surface pressure and temperature conditions suggest the presence of methane in liquid form. However, the maximum temperature lapse rate observed near the surface is compatible with the adiabatic value expected for a dry N~ atmosphere --indicating that methane saturation may not occur in this region. A lapse rate transition observed near 3.5 km can be interpreted in two different ways: either as a mark of a boundary between a convective region near the surface and a higher radiative equilibrium zone, or as a mark to the bottom of a CH4 cloud and a perfect N2-CH4 mixture below this level. A detailed examination of the temperature profile in the last kilometers near the surface tends to exclude the presence of a global methane ocean on the surface (Eshleman et ai. , 1983; Flasar, 1983) on the ground of dynamics. On the other hand, ethane is one of the main photochemical products of methane in Titan's atmosphere and also a gas which could be found as a liquid on the surface. The presence of a global ethane ocean, 1-kin deep, and in which methane in appreciable quantities and traces of nitrogen and other atmospheric condensates could be found, was suggested by Lunine et al. (1983) in a theory compatible with the observed temperature lapse rate. The uncertainty on the temperature towards 200 km is --- 10-15 K, because of the dependance of these profiles on the initial assumptions: atmospheric composition (which we know to be different than pure nitrogen), lack of calibration, perfect gas law deviations etc. Thus, the RSS data are reliable up to -.- 100 km. At higher altitudes, the thermal structure becomes more and more sensible to the initai conditions adopted in the integration of the refractivity profiles (Hunten et a l . , 1984). Above 200 kin, data obtained by the UVS experiment during a solar occultation are available. Thus, a density value of 2.7 -t-0.2 × l0 s cm -3 was measured at 1265 km for a corresponding temperature of 186 =t= 30 K which suggests an average temperature of 165 K in the 200 to 1265 km altitude range (Smith et al. , 1982). This same experiment allowed the detection of a methane luixing ratio of 8 + 3% towards 1115 km and placed the homopause level at around 925 + 70 kin. The thermal profiles by Lindal et al. (1983) are used as reference in most of the studies of Titan's atmosphere. However, they do not include a detailed analysis of the associated uncertainties.

The exact knowledge of these uncertainties is essential because it reflects

directly into errors on the atmospheric/oceanic composition and affects our comprehension


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of the physico-chemical processes in the atmosphere. b) Thermosphere The thermal balance of the thermosphere of Titan was first studied by Friedson and Yung (1984). They resolved simultaneously the equations of heat transfer and of hydrostatic equilibrium by using the UV occultation measurements of density and temperature as boundary conditions. The heat budget of the atmosphere takes into account the sources of energy (essentially solar radiation and a contribution from electron precipitation in the magnetosphere) and the losses (cooling through emission of the minor components -mainly acetylene- in a non-LTE environment). The heat is transported downwards by molecular conduction. Friedson and Yung (1984) thus modeled the thermosphere profile down to the mesopanse ( ~ 736 km and 110 K). 1.3.3

Opacity structure

Samuelson (1983) developed an analytical radiative equilibrium model that was vertically homogeneous and considered the radiation balance in three broad spectral intervals. His results confirmed pre-Voyager models in which the temperature inversion observed on Titan was the result of strong stratospheric absorption of solar UV and penetration to near the surface of longer wavelength solar visible radiation. Samuelson suggested the presence of enhanced opacity in the regions near 20 and 65 kin, which he attributed to possible condensation clouds of methane and C~HrC2He-C3Hs, respectively,, He also identified the region from 400 to 600 cm -1 as the thermal infrared "window" region, a spectral region of low optical depth that limits the greenhouse effect on T i t a n . However, the existence of condensate clouds in Titan's troposphere was not established and some disagreement also remained on the nature and extent of any such clouds. The high altitude haze observed on Titan could be photochemically produced organic solid material ("tholin') (see, e.g. Khare etal., 1982; Sagan and Thompson, 1984). Rages

etal. (1983) and Rages and Pollack (1983) investigated the properties of the aerosols from high-phase-angle Voyager images and found the particle radius to be between 0.2 and 0.5 pro. These "smog" particles form a ~ayer that enshrouds the entire globe of Titan and stretches / from the surface to an altitude of about 200 kin. These authors also estimated the optical thickness of the haze and the vertical distributions of the aerosols; they found a local peak in the extinction coefficient at vertical optical depth ~ 0.01 (340-360 km of altitude). The Voyager pictures show indeed a thin, high layer of haze about 100 km above tile main haze layer. On the other hand, a radius distribution was derived from polarimetry data by Tomasko

Titan: Recent Developments


and Smith (1982) and the two ranges were incompatible when applied to spherical particles. West et al. (1983) summarize several results in the Literature for scattering by nonspherical particles. A variety of calculations and measurements of di/ferent shaped particles, seems to indicate that large, compact, irregular dark particles may have many of the properties required of Titan's aerosols. 1.3.4

Surface and interior

The presence of aerosols in Titan's atmosphere implies scattering occuring from the aerosol and whatever methane clouds may be present at and below the tropopause. This susgests that the ground-based near-infrared spectroscopic measurements may have never corresponded to an optical path that reached the satellite's surface. Thus, very few statements can be made concerning the latter. Since most of the trace constituents in the atmosphere will condense out and precipitate to the surface, Titan's atmosphere must be in constant evolution. Methane is released into the atmosphere from the surface which may be partly composed of an ethane-methane ocean. Some of the atmosperlc methane is cycled back to the surface as "rain", and some may pass through the cold trap at the tropopause and diffuse upward to replenish the stratosphere. There it is decomposed by photolysis and particle bombardment. The ions and radicals thus produced recombine chemically to form higher hydrocarbons. Additional chemistry produces nitriles from N2 and oxygen compounds from water (not yet detected but originating presumably from meteoritic debris). These products diffuse down to levels where they condense and finally precipitate to the surface, forming a mixture of tars and snow (Hunten et al., 1984). Exceptions are propane and ethane, which remain liquid at the surface. Except for methane, vapor pressures are too low to permit recycling back into the atmosphere of the other gases. Molecular hydrogen, formed as a photochemical byproduct, diffuses upward and escapes to space. Eventually the atmosphere will become almost totally depleted in methane, leaving behind on the surface -~ 1 km of frozen or liquid residue (Hunten et aL, 1984). Little can be said with certainty about Titan's interior. The mean density of 1.86 g cm -S and radius of 2575 km show that the satellite lies somewhere between Ganymede and Callisto (see Table IV Hunten et al., 1984). These two satellites of Jupiter are believed to consist of rock (silicates and iron) and water ice (25-50% by mass). The similarities between the three satellites may suggest that they have similar interior properties, but, if Titan formed in lower temperatures than Ganymede and Callisto it may have incorporated other ices having similar


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densities to water ice (NHs-H~O and CH4-nH20). All three satellites are silicate-enriched relative to cosmic abundances, which predict 52% water ice vs 40% rock by mass (Morrison et al. ,


Titan seems to be an intermediate object between terrestrial planets (where the atmosphere consists of very minor constituents of the planet) and of the giant planets (where the atmosphere is intimately related to the deep interior). The bulk composition of Titan is probably determined by the range of condensates that formed in a dense, gaseous disk around proto-Saturn (Stevenson, 1982). Examples of these sequestrated ingredients would be: water ice, ices and silicates, ammonia, clathrates and other volatile ices and liquids. 1.3.5


The discovery of an N2 atmosphere on Titan by the combined infrared, ultraviolet and radio occultation experiments brought new light in our understanding of Titan's photochemistry. The formation of ethane and acetylene through CH4 photolysis, as well as the further catalytic dissociation of methane by C2H2 producing polyacetylenes was advocated even before the encounter (Allen et a l . , 1980), but the important effects of N~ on hydrocarbon chemistry were not considered. One of the most important issues of molecular nitrogen being the dominant background gas is the photochemical reactions with CH4 in which it can be involved, producing primarily HCN and other nitriles (HCsN, C2N2 etc.). The discoveries of CO2 and CO in Titan's atmosphere may indicate its formation with significant CO (input of meteoritic material) or sputtering of oxygen off Saturn's icy satel]ites (Lutz et al. , 1983). The existence of oxygen in Titan's reducing atmosphere also leads to some interesting photochemistry. Yung et al. (1984) proposed a model of Titan's photochemistry based on the Voyager 1 observations and on the prevailing photochemical reactions in a N2-CH4 atmosphere. According to their model, the atmosphere in the beginning contains only tile mother molecules: molecular nitrogen, methane (contained in the volatiles trapped at the time of Titan's formation) and H~O (from a small but regular meteoritic source). The interaction of these molecules with ultraviolet radiation, energetic particles and cosmic rays in the thermosphere and mesosphere produces firstly important species such as HCN, CO, CO2, C2H2, C~H4 and CHsC2H. N2 dissociation by magnetospheric electron impacts in the thermosphere, is, according to this model, the major source of HCN, while the hydrocarbons mainly result from methane photolysis. The organics thus formed are afterwards transported downward and

Titan: Recent Developments


act as the precursors of stratospheric chemistry by catalysing CH¢ dissociation to form more complex molecules. 1.3.6

Origin and evolution

There are two main possible origin scenarios for Titan's atmosphere: accretion from an unfractionated proto-Saturnian nebula or devolatilization of the ices and rocky material that accreted to form the satellite. Titan's enrichment in silicates and its depletion in neon compared with an object that condensed from an unfractionated portion of the solar nebula, indicate that Titan's atmosphere probably formed after the planet's accretion, from outgassing of the volatiles trapped as hydrates in Titan's surface (Hubbard and MacFarlane, 1980). The presence of nitrogen on Titan calls for more discussion.

The nitrogen found in

the outer solar system is generally assumed to originate in the solar nebula as ammonia. However, we cannot teU if the nitrogen we see now in Titan's atmosphere was originally in the form of N2 (incorporated in Titan's atmosphere by accretion in the form of hydrates) or NH3 (captured, instead of N2, in the hydrates). A distinction between the two cases would have great implications on the formation conditions of the sateUite (Owen, 1982). Atreya et al. (1978) proposed a scenario postulating the existence of a warm period at the beginning of Titan's evolution, during which ammonia outgassed from the interior and into the atmosphere where it photodissociated into molecular nitrogen. A surface temperature of >__ 150 K is needed to account for the amount of N2 observed today in the atmosphere of Titan. On the other hand, the absence of significant argon or neon and the UV low upper limit found for CO, lead Strobel (1982) to propose a different case for the formation of Titan's N2 atmosphere. In this scenario, Titan formed in cold accretion (nebular temperature <_ 60 K) in which case N2 sublimates preferentially as a solid clathrate hydrate. This mechanism could explain the upper limits set for CO and the rare gases. Strobe/(1982) concludes that N2 outgassing from Titan's surface as it warmed up to its present 95 K is a possible source of the satellite's atmosphere. Another important chemical tracer in the outer Solar System is deuterium, found on Titan in the form of CH3D. Monodeuterated methane was identified in the Voyager infrared spectra by Kim and Caldwell (1982) who retrieved a relative abundance of D / H ~ 4.2 × 10 -4 suggesting that Titan is enriched in deuterium compared to the giant planets.



A. Coustenis

Thermal structure and composition of the atmosphere

This and the following section are devoted to the progress accomplished and to the new developments achieved in Titan science in the past five years. This goal is, however, much too axnbitious in itself, since one is bound to omit relevant contributions. All the more, the author is certainly not sufllciently acquainted with all the different aspects of Titan to be able to succesfully conduct a complete overview of all the physical mechanisms active on the satellite. As a consequence, a major part of this paper is concentrated on the temperature and composition of Titan and in particular on some recent analyses performed by the author and colleagues, as a complement to parts 1.3.1 and 1.3.2. Section 3 gives an, undoubtedly, rough outline of recent progress in other aspects, as a complement to parts 1.3.3 to 1.3.6 of the Introduction.


Re-analysis of Voyager RSS data

Although central to our understanding of the Titan atmosphere, the Voyager IRIS spectral observations of Titan have been woefully underanalyzed. What follows is a brief report on an attempt by the author and colleagues to correct this situation. The results exposed hereafter are our major contribution in this vein. The thermal flux emitted by Titan depends on both the thermal profile and on the mixing ratios of the absorbers as related by the radiative transfer equation. In order to determine the abundances of the gases present in the infrared spectra of T i t a n , one has to start by defining as well as possible the thermal structure of the satellite. To do so, as a first step a combined analysis of the radio-occultation (RSS) and of the infrared (IRIS) data permitted us to constrain the thermal structure at the equator by yielding two extreme temperature profiles from the ground up to 400 km. These profiles were associated to limit cases of atmospheric composition. The abundances of the major components (CH,, Ar and N2 indirectly) were also bracketed by extreme values. Still at the equator, a nominal temperature profile was derived from the best agreement obtained with three different sets of infrared data. It corresponds to a given composition of the atmosphere (-~ 98% of nitrogen, -~ 2% of methane and 0.2% of hydrogen). As a second step, this study was extended into a model of the thermal structure in the thermosphere (0-1300 km), based on the RSS and IRIS data and also on tile solar occultation (UVS) measurements. Once the temperature structure is defined in all possible locations in Titan's atmosphere, the chemical composition of the satellite can be derived as a function of latitude, from the best fit of different spectral averages. Thus, the abundances of the minor species in the atmosphere of Titan (C2He, C~H2, C4H~, CsH4, C~H4, C3Hs, CHsD,

Titan: Recent Developments


HCN, CO~, HCsN, C=N2, HCN) with the assodated uncertainties were retrieved with high precision. 2.1.1

T e m p e r a t u r e a n d m a i n c o m p o s i t i o n o f t h e lower a t m o s p h e r e

In principle, the IRIS data can provide both the reference required to calibrate the radiooccultation profile and independant thermal profiles. In actual fact, we need to know the distribution of the opacity sources. Methane is the only gas involved and, although we may assume that it is uniformly mixed in the stratosphere, its abundance is not directly measurable but has to be deduced from physical considerations. Another difflcnlty is that we dispose of only one methane band in the infrared spectrum (at 1304 cm-1), which allows us to probe the 0.01 to 10 mbar (,,, 100 to 400 kin) atmospheric region. Thus, the thermal structure derived through inversion of the radiative transfer equation, using methane as an opacity source, does not exactly join the radio-occultation temperature profile in the zone where the latter is reliable. The connection between the two profiles remains therefore uncertain. A normalization procedure, practically independent of the exact nature of the opacity sources, was developed in 1981 by Samuelson et al. Contrary to the other spectral regions, the IRIS data recorded at low latitudes from 200 to 540 cm -1 show little dependence on the emission angle. This indicates that either the emissions from the troposphere or the stratosphere are almost equal, or that the atmosphere is transparent at those wavenumbera and the emission observed originates essentially from the surface. The radiative transfer analysis shows in fact that the troposphere is opaque at 200 cm -1, with a temperature brightness temperature minimum of 74.0 -t- 0.5 K. This implies an upper temperature limit at the tropopause of the equator: < 74.5 K (Saznuelson et al., 1981). By comparing with the radio-occultation profile, we reach an upper limit of the mean molecular weight (29.36 ainu) which suggests the presence of a heavier component than nitrogen in Titan's atmosphere. The most probable candidate is argon (Owen, 1982). Samuelson et al. (1981) advocated an almost transparent troposphere at 540 cm -1, which allowed them to evaluate a lower limit of the surface temperature: 94.1 K. Taking into account a 0.5 K difference which may exist between the surface temperature and the gas temperature immediately above --due to turbulence-, Hunten et al. (1984) derived a lower limit of the atmospheric temperature above the surface (93.6 K) which implies a mean molecular weight value of _> 27.8 ainu. This lower limit is compatible with CH4-N2 (and no At) atmosphere, if the methane abundance does not exceed 1.6% at the surface (Gautier, 1985). On the other hand, ground-based observations in the 1.2 to 1.6 pm region (Fink and Larson, 1979; De


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Bergh et a l . , 1988) suggest a greater CH+ abundance in the lower atmosphere, which would favor the presence of argon. However, Toon et al. (1988) find the H2 and aerosol opacity in the 500 to 600 cm -t region to be too high to allow probing of the surface and consider all previous determinations of surface temperatures as incorrect. The presence of argon is then again an open question. As a consequence, the lower limit of the mean molecular weight may not be reliable. The thermal structure in the 0.01 to 10 mbax region (100 to 400 kin) can be estimated in two different ways, both of which require extrapolation toward the radio-occultation profile. The first method consists in reproducing as well as possible the emission observed in the v4 methane band at 7.7 pro, with radiative transfer programs incorporating different thermal profiles. The temperature profile leading to the best fit of the observations is adopted. The second approach uses the relation between the brightness temperatures measured at 1304 and 1260 cm -1 and the associated emission levels, that is the pressure (or altitude) levels corresponding to the maxima of the respective weight functions. In general, the results of the two methods - b o t h of which need to assume a stratospheric methane mole fraction, since it cannot be directly measured- are in good agreement. Using the first method, the refractivity data obtained during the radio-occultation experiment were lately re-analysed using different main atmospheric compositions (Lellouch et a l . , 1989). With no available observational constraints on the methane abundance, the CH4

vertical distribution used in that paper is defined as follows: below the "cold trap" ( ~ 30 kin), the m a x i m u m CH4 abundance is determined at each level by the local saturation law; above, the methane abundance is taken to be constant and equal to the minimum mixing ratio of the "cold trap". Thus, the CH4 stratospheric mole fraction depends on the "cold trap" temperature, which, in turn, is a function of the composition. The influence of the tropospheric methane clouds is not accounted for, because there is no information on the abundance of the methane condensate or its refractivity. This procedure is justified by the authors. Beside the uncertainties on the molecular mass, this analysis considers also the uncertainties on the refractivity profile due to noise. More in detail, starting with the Lindal et al. (1983) results, Lellouch et al. generated different temperature profiles by varying the main chemical composition (N2, CH4, Ar) whereas Lindal et al. had calculated their thermal profiles for a pure N2 atmosphere. The new thermal profiles were constrained to satisfy the condition: T(inversion) _< 74.5 K. For a given Ar/N2 ratio, the m a x i m u m CH4 abundance is found by iteration. According to the first method mentioned above, synthetic spectra were generated in the 1270 to 1340 cm -t spectral region for each thermal profile and compared to the observations until, by trial and error, the best

Titan: Recent Developments


possible fit of the emission observed in the v4 methane band, centered at 1304 cm -1, was obtained. The associated thermal profile was then considered as acceptable. Using the limit composition cases, extreme thermal profiles were also obtained. LeUouch et al. tested a great number of thermal profiles before deriving the "envelope" temperature profiles shown in Fig. 9 of their paper and the nominal profile proposed for a stratospheric methane abundance of 1.5%. The following conclusions were also derived: the Ar/N2 ratio in the atmosphere is lower than 27%, but could be nil. As a consequence, the presence of argon in Titan's atmosphere is not demonstrated, but an upper limit was inferred. the stratospheric mole fraction of methane lies between 0.5 and 3.4% (1.0 and 1.7% if argon is absent). The maximum CH4 abundance in the troposphere is 21% at the surface level (12% if there is no argon). The new uncertainties on the thermal profile are q- 2 K at the tropopause and q- 4 K between 150 and 200 kin. The surface temperature lies between 92.5 and 101 K (or 95.7 K if there is no argon). The "real" temperature profile is probably warmer than the one suggested by Lindal et al. (1983), especially in the stratosphere. with respect to the properties of a putative ocean on Titan's surface, this analysis shows that it could be almost entirely composed of either ethane (and be relatively shallow --700 m - - at 92.5 K) or methane (and -.- 10 km deep at 101 K), or anything between. Dissolved nitrogen (2-6%) can be incorporated. If the ocean is warm (that is if the surface temperature is closer to the upper limit), it could be an efficient CO reservoir. Finally, for most of the constituents, saturation is possible in any type of ocean. Fig. 1 shows a nominal and the associated extreme thermal profiles for the equatorial region, as proposed by Coustenis et al. (1989a). 2.1.2

The temperature

profile in t h e t h e r m o s p h e r e

In a paper by Lellouch et al. (1990), the thermal balance of Titan's atmosphere was reevaluated. The need to do so followed the discovery of some errors in the solar heating profile and in the expressions of the collisional relaxation rates in the work of Friedson and Yung. Indeed, the vertical temperature profiles proposed for the whole 0-1300 km altitude range, which were based on the study of Friedson and Yung, incorporated high temperatures in the stratosphere to compensate for the cold mesopause. Calculations (Coustenis and Gautier,


A. Coustenis

1987) showed that these profiles generated excessive fluxes in the v4 methane band and were therefore incompatible with observations. The new thermal profiles derived from the re-evaluation were on the one hand tested against the emission observed in the v, methane band, according to the procedure described above, which yielded information in the 150-400 km zone, and on the other hand constrained by the UV temperature-density point obtained at 1265 km (see Introduction).

Fig. 9 of

Lellouch et al. (1990) shows a temperature profile which satisfies these two conditions. It is composed of a "warm" stratospheric region (T -,~ 175 K) up to 500 km, to compensate for a cold mesopause (135 K) at 800 km. The associated methane mole fraction is 1.7%. The results of the Lellouch ct al. (1989) paper on the lower atmosphere impose an additional constraint at 200 km (T=173 + 4 K) and provide at the same time the continuity of the thermal profile down to the surface. The model thus obtained shows that the thermal structure in Titan's thermosphere is extremely dependant on the solar heating efficiencies and on the relaxation rates of the cooling agents (mainly acetylene). Given our little information on these parameters, one may, in theory, obtain several solutions. The only firm conclusion of this exercise seems to be the existence of a mesopause, due to the very efficient cooling resulting from the radiative emission in the Vs band of C~H2. Precise and complete knowledge on Titan's thermal structure is hindered at the present time by lack of information on the atmospheric zone above ,~ 400 km. However, in the future some improvement could be gained from ground-based observations during occultations (see 2.3) or from millimeter observations in the CO and HCN bands from which we can derive temperature values in the 400-600 km range (Paubert et al., 1984). Also, the Cassini mission will certainly complete our knowledge of Titan's thermal structure in 2003 through in situ measurements at the equator and from data acquired by the experiments on board the orbiter.

These conclusions ~hould be compared to the a~umptions and first ~tudie~ deacribed in the previou~ chapter.


N e w analyses of Voyager 1 data

In the past years, beside the information on the thermal structure, improvements were acquired in other aspects, such as spectroscopic data or cloud model of Titan's atmosphere (Samuelson and Mayo, 1990). The author and eolleagues have undertaken the task to reinvestigate the temperature and chemical composition of Titan's atmosphere at all latitudes. The method used, fuRRy described in Coustenis et aL (1989a), consists of using a radiative

Titan: Recent Developments



8 ,







O aO

\ \

0 \




'l f~



? 0




(Joq~) 3~lr~SS]~ld



A. Coustenis






























: ~ / "~ ~o~

¢J m


f .I~° "~'~

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Titan: Recent Developments


transfer program, in which the temperature, clouds and spectroscopic parameters are incorporated, to derive the mean mole fractions of the gases with good precision. Assuming a methane mole fraction of 1.8%, nominal temperature profiles are proposed for different latitudes in the 0-400 km altitude range (as the one shown in Fig.

1 for the

equator). Outside the equator, there are no radio-occultation data available. We have thus used only the v4 CH4 band to infer information on the temperature in the 150-350 km altitude range, assuming that methane is uniformly mixed in the atmosphere. When the thermal profile is known, the radiative transfer equation can be used to derive the mixing ratios of the absorbers. The iterative procedure we use can be simply defined as follows: a Line-by-line radiative transfer program, incorporating the temperature profile and using the abundances of the absorbers as parameters, generates synthetic spectra in the 200-1200 cm -I spectral range. These spectra are compared to the observations until the best agreement is obtained by trial and error. We then extract the associated mixing ratios of the molecules presenting emission bands in the spectra at a given latitude. The atmospheric opacity in our model includes contributions from molecular sources and from clouds. The radiative transfer calculations include the colLision-induced absorption by N=-N=, N2-H2, N2CH4, CH4-CH4, and molecular bands from CH4, CH3D, C2H2, HCN, CsH4, HCsN, C4H=, C2H4, Calls, C2Hs, C2N2 and C02. The continuum opacity is simulated by a cloud model similar to that of Samuelson (1981). The synthetic spectra, generated by the radiative transfer program, are convolved with the Fourier transform of the Hamming function at the spectral resolution of Voyager: 4.3 cm -I, prior to comparison with the IRIS spectra. Fig. 2 shows such a synthetic spectrum superimposed with an observational spectral average of Titan's equator. Spectroscopic parameters in our calculations include some very recent analyses and originate from different sources (as in Goustenis et al. , 1989a). With the exception of a special north pole sequence - f r o m which vertical distributions were derived, as will be explained below- at all the other locations mean mole fractions of the absorbers were thus derived in the 200-1200 cm -I spectral region (Coustenis et a l . , 1989a, b; Coustenis and B6zard, 1990). The 200-800 cm -1 spectral region yielded mean abundances for C2Hz, C~H4, C2Hs, CsH4, Calls, C4H2, HCN, CzN2, HCaN mad C02 while emission bands around 1150 cm -1 allowed us to check the propane abundance and then to evaluate the CHaD mole fraction at the equator. The latter allows us to obtain a value of D/H, a ratio most important in cosmological models. Our value (1.5 x10-4), in good agreement with other measurements from groundbased observations (de Bergh et a l . , 1988), indicates an enrichment of deuterium in Titan's atmosphere, with respect to the protosolar value. The reasons for this axe investigated in


A. Coustenis

Coustenis et al. (1989b). This method has been applied to 8 different data sets up to now. From the 300 usable IRIS spectra (Coustenis et al. , 1989a), eight first selections were made and the spectra in each one were co-added. This allows us to obtain an average spectrum with improved signalto-noise ratio. Each selection contains -,. 30 spectra and corresponds to locations on Titan's disk from 60 ° S to 70 ° N. There are three selections near the equator at a mean latitude of 7° N, and four near +50 °, +70 °, - 3 2 ° and - 5 3 °, respectively. The observational sequence over Titan's north pole (RTSAT) contains 9 spectra taken at grazing incidence and divided into 3 averages. The latter are associated with altitudes 269, 232 and 121 km above the surface in the atmosphere. Due to its particular geometry, this sequence includes information on the height-dependence of the gas mixing ratios. 2.2.1


From a study of the temperatures obtained at 0.4 and 1 mbar as a function of latitude for seven different selections, one remarks that there is a larger temperature decrease as one moves from the equator to the north pole than when moving from the equator to the south pole. The south pole appears to be 2-3 K colder than the equator, while the north pole exhibits a drop in temperature by up to 20 K with respect to the equator (see Fig. 3). This is in excellent agreement with the results of Flasax and Conrath (1990), who retrieved the latitudinal variations of temperature from inversions of IRIS spectral measurements in the v4 methane band for various selections over the disk. The Voyager encounter took place at the time of the northern spring equinox. Flasar and Conrath suggest that the observed equatorto-north pole temperature gradient may represent the response of Titan's stratosphere to seasonal forcing, which lags behind the radiative equilibrium state because of strong dynamical inertia.

Another possibility is that the observed cooling in the north polar region is

the result of radiative equilibrium in the presence of an enhanced concentration of infrared emitters (gases and possibly aerosols).

Titan: Recent Developments




TITAN temperature latitudinal variations "






o.¢ mbar

170 1.0 r n b a r v

i,i n,-

16o rv i,i O. LAJ




















Fig. 3 Temperature latitudinal variations from thermal profiles obtained at 52° S, 32° S, 7° N, 50° N, and 70° N. The temperatures are given at two distinct pressure leveb (0.4 and 1 mhar). The 7°N pointl correspond to the three spectral averages near the equator, which were fitted with synthetic spectra using the same atmospheric model. The north polar region shows a dramatic decrease in temperature (by 15-20 K) with respect to the equator in the stratosphere, whereas the south pole is only colder by 2-3 K. This equator-to-pole asymmetry can be explained by seasonal and dynamical effects (see text). Error bars are indicated. Figure from Coustenis et aL (1989c).



In terms of composition, Table 2 (adapted from Owen and Gautier, 1989) shows the updated abundances measured by Coustenis et al. (.1.989a) at the equator and by Coustenis et al. (1990) at the north polar region. On a planetary scale, the following results were obtained. Ethane, acetylene and propane are the three most abundant hydrocarbons after methane.


molecules were found to be deprived of any noticeable compositional variations in latitude. They seem to be homogeneously mixed in Titan's atmosphere from pole to pole.


A. Coustenis

T~tble II Stratospheric composition of Titan Gas

Mole fraction

I Comments-References

Ma~or components Nitrogen Argon Methane Hydrogen

N2 Ar CH4 H~

0.757-0.99 0-0.21 0.005-0.034 0.002-0.006


Hydrocarbons Acetylene C2H2 Ethylene C~H4 Ethane C2Hs Methylacetylene CsH4 Propane CsH9 Diacetylene C+H2 Nitriles Hydrogen cyanide HCN Cyanoacetylene HCsN Cyanogen C2N2 Oxygen compounds Carbon dioxide C02 Carbon monoxide CO

(a) (b) (c) (d) (e)

Inferred indirectly (a) Inferred indirectly (a) In the stratosphere (a)


North pole

xlO - s xlO -s xlO -s xlO -9

2.2 xlO -s 3 xlO -s 8.8 xlO -s 1.7 xlO -s

7.0 xlO -z 1.4 xlO -9

5.0 xlO -7 2.2 xlO -s

(c) (c) (c) (c) (c) (c)

1.6 x 10-7 _<1.5x10 -9 < 1 . 5 x 1 0 -9

6 x 10-7 7.0 xlO -s 4.5 ×10 -9

(c) (e) (c)

2.2 9.0 1.3 4.4

1.4 x l 0 -s _<7 x l 0 -9 6 x l 0 -s _< 4 x10 -6

(c) In the troposphere (d) In the stratosphere (e)

Lellouch et 02. (1989). Samuelson et 02. (1981), Toon et 02. (1988). Coustenis et 02. (1989a), Coustenis eta/. (1990). Lutz et 02. (1983). Marten et 02. (1988).

Ethylene (and methylacetylene), tend to increase (by a factor of 3-4) near the north polar region, while their mixing ratios remain constant from mid-latitudes to the south pole. For the C4H2 mixing ratio, there is practically no difference between the equator and the south pole, whereas a •ctor of --- 20 increase is observed near the north pole. Two of the nitriles (HC3N and C~N2) are absent in the spectra corresponding to latitudes < 60 °. However, upper limits at a 3~ level obtained from these selections indicate an increase of a factor of at least -,- 50 in HC3N and of > 4 in C2N2 in the north polar region. The other nitrile observed as a gas in the IRIS spectra is HCN, whose enfission spectral

Titan: Recent Developments








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o~ ~

,, !!,






-~.~ . ~



A. Coustenis


" ~ Z r~



I °~,,~

o o


0 o

(.mctm) ~anssoad


Titan: Recent Developments


feature is found at all latitudes. HCN shows a steady increase in abundance from pole-to-pole by a factor of -~ 12 (factor of -,~ 4 from the equator to the north pole). CO=, on the other hand, is the only gas up to now found to decrease near the north pole. Its constant abundance of 1.4 × 10 -8 from the south pole to the equator is reduced by a factor of at least 2 at the north pole. In the spectra of high northern latitudes the v2 band of carbon dioxide is totally absent. Fig.

4 illustrates the above findings for five molecules: C=H=, HCN, CO2, C4H2 and

HCsN. It should be noted at this point, that Titan's composition may he more complex than current evidence suggests. Indeed, laboratory simulations (see for instance Cerceau etal., 1985) imply that we may not actually have access to the maximum order of chemical complexity in the satellite's atmosphere. Macromolecules, more complex hydrocarbons and nitriles than those observed from earth or by Voyager could be present in fair amounts on Titan. However, due to weak bands or to lower quantities than have those already detected, we do not have evidence for these species yet.


Vertical variations

Coustenis etal. (1990) investigated the special observational sequence (RTSAT), consisting of ,,-30 spectra recorded over Titan's north polax region. Nine of the spectra were taken at grazing incidence over the sateUite'ssurface,in other words they correspond to lines of sight which do not intercept the surface but follow opticalpaths entirelysituated in the atmosphere. The advantage of these observations is that they can be associated to certain altitude levels and therefore yield height-dependent information on temperature and composition. Also, as we move deeper into the atmosphere, at lower altitudes, the optical paths become longer and so the emission features in the spectra axe enhanced. This latter aspect has allowed the detection of H C s N and C2N2 in the R T S A T spectra. Also, two emission bands of C4H2 and Call4 at 220 and 328 c m -I respectively,absent in the equatorial spectra, were observed in the north polar ones. The vertical concentration profilesof all the gases present in the R T S A T

spectra were

derived. Five of them (CsH4, HCsN, HCN, C2H2, C~Hs) are plotted in Fig. 5. The general tendency is an increase of the mixing ratio with altitude, confirming the prediction of the photochemical models that these species form in the upper atmosphere and then diffuse downwards in the stratosphere. Below the condensation levelof each gas, the distributionsare assumed to decrease following the respective vapor saturation law. However, the gradients in


A. Coustenis

abundance represented by these profiles are not always in accordance with theoretical models (by Yung et al., 1984 or Yung, 1987). The latitudinal variations of the gases have not been fully explained up to date. The nitrile enhancement observed near the north pole was tentatively explained by Yung (1987) who suggested that the nitriles accumulate in that region during the winter when the polar region is in the shadow (see Fig. 4 of Yung, 1987). At the time of the Voyager 1 flyby, the north polar region was just coming out of winter, having maximum nitrile abundances. Although the Voyager encounter occurred -.~ 8 months after spring equinox, chemical time constants may have been large enough to prevent the UV flux from immediately depleting the strong nitrile concentrations observed. This could also explain the depletion in CO2, if this molecule is mainly produced from the H20 photolysis (Yung et al., 1984). Its formation may have then been inhibited in the north pole because of the low level of UV radiation. On the whole, the results show that at the time of the Voyager encounter with the satellite, latitudinal variations in temperature and in composition existed, notably from the equator to the north pole. Current models are unable to secure an explanation for the observations made by IRIS which are localized and seasonaLly-dependent (e.g. for the enrichment of some hydrocarbons (C4H2, C2H4 and CsH4). The vertical mixing ratios derived over the north pole can only be associated to models including seasonal and dynamical effects. A complete pole-to-pole model of Titan's atmosphere in terms of temperature and composition is a goal for the future. Based on such a model, other physical parameters in Titan's atmosphere will be constrained.




Information on the thermal structure and composition of Titan's atmosphere was secured not only through space missions but also from ground-based observations using different techniques: spectroscopy, radar, stellar occultations, etc.. 2.3.1

Deuterium: CHsD and D/H ratio

As mentioned in 2.2.2, observations of the 3v2 CHsD band at 1.6 pm by de Bergh et al. (1988) confirm the value of 1.5 × 10 -4 found by Coustenis et al. (1989b) from Voyager data analyses of the v6 CHsD band at 8.fi ~m and support the evidence for a deuterium enrichment in Titan's atmosphere with respect to the protosolar value and the one found in the giant planets (D/H --~ 2 - 3.4 × 10-s). However, the enrichment factor is not known with enough precision to allow

Titan: Recent Developments


firm determination of its origin. Several hypotheses have been suggested. Owen et al. (1986) propose a primitive nebula model in which two distinct deuterium reservoirs co-existed before the formation of the solar system: the main one in gaseous hydrogen (HD) and the second smaller one in deuterated isotopes of CNO compounds -such as CH4, NHs, and H~O- trapped in ices or clathrates (Lunine and Stevenson, 1985) or adsorbed on grains. Hubbard and MacFarlane (1980) were the first to suggest that Titan's atmosphere could result from volatile degassing of grains originating from the second reservoir and having accreted to form the satellite.

The fractionation mechanisms acting during or after the

formation of the satellite could have contributed to the deuterium enrichment observed on Titan. Pinto et al. (1986) have suggested several processes including exchanges of deuterium between methane gas and cloud particles or the putative ethane-methane ocean, or even the hypothetical icy crust underneath the ocean. An additional possibility could involve isotopic exchange catalyzed in the presence of metallic grains in the saturnian nebula. Better precision in measurements is needed to distinguish between the different theories which could also combine to explain the deuterium enrichment observed on Titan and elsewhere in the solar system (see Fig. 6).


Other constituents

Marten et al. (1988) estimated the CO/N~ mixing ratio in the stratosphere of Titan from millimeter heterodyne observations at 115.3 GHz and found it to be ~ 2 x l 0 -e, in conflict with an estimate by Lutz et al. (1983) of 6 x l 0 -5 for the tropospheric value, from infrared observations, unless one accepts that the mixing ratio of CO strongly varies with altitude. Also from millimeter observations, very recently, the stratospheric profile of HCN on Titan was derived by Tanguy et aL (1990). This vertical distribution, obtained from an analysis of the (1-0) line of HCN at 88.6 GHz, is found to increase with altitude. The mean mixing ratio scale height is ~ 47 km in the 100-300 km region and the HCN abundance around the 170-kin level is -.- 3.3 x l 0 -T. These results are consistent with the analyses of Voyager infrared measurements (Coustenis et al., 1990). Conversely, the inferred vertical concentration gradient is much steeper and the abundance in the lower stratosphere smaller than predicted by photochemical models (Yung et aL , 1984).


A. Coustenis







0 0

? o

Protosolar range URANUS

ISM range

Fig. @ The D/H ratio observed in different bodies of our solar system. The protosolar value comes from recent revisions by Anders and Grevesse (1988) of the generally adopted estimate of Geiss and Bochsler (1981). This new value takes into account the re-evaluation of the 3He/4He ratio in the solar wind and the most recent estimations of the He/H protosolar ratio from solar evolution models. Value (a) for Titan is the result of Voyager infrared data analysis Coustenis et al. (1989b) and the measurements from earth by de Bergh et al. (1988). lue (b) allows for the fractionation which could have taken place in Titan's atmosphere since its formation (Pinto et al., 1986). The values in the giant planets and in comet Halley are taken from Gautier and O w e n (1988).


Titan: Recent Developments







~'.~,~ •

® ~N ~.~ ® ~ ® •

. ~ ~.~.~ ~.g ~ ~ ~'~ •





,mqw " : J U N S S ] ~ I


~ 9~ ~

~ E'-, .,-~







A. Coustenis



The July 3, 1989 stellar occultation of Titan allowed Sicardy et al. (1990) to infer some information on Titan's stratosphere. The occultation of star 28 Sgr by Titan was observed from Paris Observatory (Meudon), Pic du Midi Observatory and Catania Observatory. This rare event provided information in the 250-500 km altitude range.

Sicardy et al. (1990)

derived a mean scale height of 48 km at 450 km altitude (-,~ 3 p b a r level). This allowed them to constrain the mean temperature between 149 and 178 K at that level. Titan's shadow center passed within ~ 20 km south of Meudon. The central flash observed there was a unique opportunity to constrain the apparent oblateness of Titan at the 0.25 mbar level (~ 250 km of altitude), giving a value which may be as high as 0.014. The same event was studied by Hubbard et al. (1990). They find mean temperatures around 180 K in agreement with models of Titan's mesosphere by Lellouch et al. (1990).

3 3.1

P r o g r e s s in o t h e r a s p e c t s Surface and clouds

Because of the dense haze-cloud deck surrounding Titan, the properties of the surface, aerosols and clouds remain unknown to this date. However, based on earth or spatial measurements, models have been developed very recently to complement and improve previous studies (see 1.3.3 and 1.3.4). An important parameter in haze models is the geometric albedo, which for Titan is around 0.2. From earth-based observations Neff et al. (1984) obtained data on the geometric albedo of Titan's atmosphere over the wavelength range 0.3 to 1.0 p m and extended them to 2.5 p m (Neff et al., 1985). Lunine and colleagues have been suggesting, since 1983, the presence on Titan's surface of an ethane-methane ocean (see 1.3.2). A schematic model of the ocean-atmosphere evolution on Titan is shown in Fig. T. The fundamental driving force in the long-term evolution of Titan's atmosphere is the photolysis of methane in the stratosphere to form higher hydrocarbons and aerosols. The current rate of photolysis, together with the inferred undersaturation of methane in the lower troposphere, suggests a reservoir of methane mixed with higher hydrocarbons. This "ethane-methane ocean" serves as both the source and the sink of photolysis, and contains a mass of nitrogen comparable to the atmospheric abundance. In the absence of outgassing or external resupply, the ocean composition evolves toward more ethane rich as

Titan: Recent Developments


methane is photolyzed over geologic time. The atmosphere responds to the change in ocean composition with a corresponding change in gueous composition and spatially averaged cloud composition (Lunine et al., 1989). Recent investigations of Titan's surface modeling include also studies by Lellouch et al. (1989) and Dubouloz et al. (1990), who, based on the assumption of an ethane-methane

surface have elaborated models to constrain the surface temperature and composition (see 2.1.1). The latter authors have derived the solubilities and thicknesses of precipitate for several minor oceanic components. Considering the extreme temperature values given by Lellouch et al., they derive the following main ocean features: T,~, = 92.5 K, CH4(lower trop.) =1.55%, Ar=0 and ocean composed of _~ 90% ethane, 7.3% methane and 1.8% nitrogen; T~,, = 101 K, CH4(lower trop.) =21.1%, Ar=17% and ocean composed of _~ 5% ethane, 83.4% methane, 5.6% argon and 6% nitrogen. The ocean can be an important CO reservoir. Conversely to these studies, radar soundings of Titan by Muhleman et al. (1990) indicate that the surface may be non-uniform: mostly solid with small lakes. The echoes obtained by these authors Show that Titan is not covered with a deep global ocean of ethane. The mean radar cross section obtained (0.35 ±0.08) is very high, most consistent with that of the Galilean sateRites. The reflexion data showed no evidence of the putative liquid ethane. The surface could have a component of solid hydrocarbons, but liquid hydrocarbons could exist on Titan's surface in the form of small lakes. Concerning methane condensation in Titan's atmosphere, Toon et al. (1988) created a model investigating the possible properties of methane clouds and based on the Voyager infrared data. Their model suggests the presence of patchy methane clouds limited below 30 km of altitude and containing > 0.5-#m-size particles which can be classified as rain. Compared with terrestrial clouds, Titan's methane clouds have larger but fewer particles and more closely resemble terrestrial rain than terrestrial clouds. McKay et al. (1989) developed a radiative-convective model of the thermal structure of Titan's atmosphere, as an extension to the work of Samuelson (1983). Their model is founded on improved determinations of coliision-induced gas absorption coefficients and constraints of the Voyager IRIS data. A review of their results -which in general confirm those of Toon et al. (1988) as concerns the clouds- is given, among other, by Lunine et al. (1989). The

calculations are based on the surface temperature found by Voyager, which corresponds to a balance between the amount of solar visible and atmospheric thermal radiation emitted (Samuelson, 1983). They find that gas and haze opacity alone suggest radiative convective balance everywhere in the atmosphere. Among other, McKay et al. have examined the effects on the greenhouse warming of pressure-induced gas opacity, in the infrared absorption of


A. Coustenis

sunlight by gases and particles in the stratosphere and methane clouds in Titan's atmosphere. The first two factors were found to make important contributions to the greenhouse warming. In agreement with Toon et al. (1988), the latter factor is found to have at best a neutral effect or, at worst, a cooling effect. McKay et al. (1989) find however a better agreement with the IRIS data when some cloud infrared opacity is added. These clouds could be limited in a certain atmospheric zone, formed of very large particles or broken in areal coverage. However, additional caiculations, coupled with laboratory scattering measurements, are still needed before a particle size and shape is demonstrated to be able to reproduce all the Titan observations. Improved cloud physics caiculations, using the temperature-pressure structure of Titan's atmosphere determined by Voyager with appropriate aerosol mass generation rates and reasonable variations of particle size with depth, could give additional constaints on the total thickness of aerosol haze required for thermal balance in Titan's lower atmosphere. 3.2




The three possible scenarios that have been proposed to explain the presence of a massive atmosphere of nitrogen on T i t a n - - ( a ) direct capture of N2 from the solar or protoSaturnian nebula, (b) devolatilization of nitrogen-containing ices condensed from the nebula (the clathrate hydnate mode/), and (c) photolysis of ammonia in the early phases of accretion-are briefly but efficiently explained in their current understanding by S. K. Atreya (1986) in his book "Atmospheres and Ionospheres of the Outer Planets and their Satellites". Considering all possible mechanisms, Titan's atmosphere is probably "secondary" and not captured as suggested by (a). A combination of (b) and (c) on primordial Titan appears to be the most plausible scenario for the evolution of its present atmosphere. Studies which have recently led to interesting models of the various physical processes in Titan's atmosphere, include, among others, the foUowing. The photochemistry active in the satellite's atmosphere was modeled by Yung et al. (1984) and upadated by Yung (1987). Lunine and Stevenson investigated recently the evolution of Titan's coupled ocean-atmosphere system and interaction of ocean with bedrock (1985: in Ices in the Solar System) and the origin of methane on Titan (1987). The present state and chemical evolution of the atmospheres of Titan, Triton and Pluto is exposed in a chapter of Origin and Evolution o[

Planetarll and Satellite Atmospheres by Lunine, Atreya and Pollack (1989) and references therein. Some of the questions related to the issue of Titan's origin and evolution, which still remain unanswered, are asked by these authors. A selected list runs as follows:

Titan: Recent Developments


1. What is the origin of nitrogen in Titan's atmosphere? Is it derived from a gas containing molecular nitrogen or was it produced by chemical processing of ammonia? 2. How much carbon monoxide was introduced into Titan and its early atmosphere? 3. Was methane brought into Titan as frozen condensate or entrapped in clathrate? 4. Why does Titan have an atmosphere, while the Galilean satellites do not? 5. Could most of Titan's atmosphere have been added by impact of volatile-rich comets after accretion? 6. What would be the effect of an ocean in contact with the atmosphere, as opposed to solid surface, on the evolution of Titan? Several attempts have been made to answer these questions and modeling is still under development. Lunine

e t al.

(1989) address these questions and others and, in particular, they

thoroughly investigate surface models involving an ethane-methane ocean and its implications on the coupled evolution of Titan's ocean-atmosphere system (see above). Below are summarized the tentative explanations proposed for some of the other questions in the above list. Precise knowledge of the abundances of the noble gases, such as argon and deuterium, in Titan's atmosphere may be diagnostic tests of the competing hypotheses on the origin of nitrogen and methane on Titan. The outstanding problem of the origin of atmospheric nitrogen on Titan (was it introduced as molecular nitrogen or NHa?), could find an answer in measurements of the Ar/N2 ratio in the present atmosphere: a large argon abundance suggests N2, while a small one favors NH3. Unfortunately, argon has not yet been detected on Titan. A deuterium enrichment could originate from different procedures acting during or after the satellite's formation: fractionation processes involving deuterium exchange with cloud particles, ocean or crust; isotopic exchange catalyzed by metal grains in the saturnian nebula; outgassing of already D-enriched grains of interstellar origin, etc. A precise measurement of the D/H value on Titan should tell us something about the ways in which hydrogen-containing compounds were incorporated in planets and satellites. The atmospheric CO and C02 observed in Titan require the existence of an exterior oxygen-bearing component in the atmosphere, such as water, which has not yet been detected. If present, H20 would then quickly dissociate into OH which combines with CH4 photolysis products (such as CH2, CHs, etc..) and produces CO and CO2 (Samuelson

Jt~VA 3 4 s l / . ~ - D

e t ai. ,

1983). The


A. Coustenis

source for water on Titan can be found in the tings of Saturn, the meteorites or comets. CO could otherwise be directly captured in a clathrate-hydrate directly from the solar nebula. However, all of these mechanisms involve numerous uncertainties which have not yet been quantified and do not allow for a definite preference between the models. There are several reasons which could explain why Titan is the only satellite with a large N2 atmosphere in contrast with the Galilean moons of Jupiter and Triton: (i) accumulation of an atmosphere on Titan could have proceeded unperturbed by losses from energetic charged particles, because Titan remains mostly outside Saturn's magnetospheric environment; (ii) Titan formed in the cooler part of the nebula from which it was accreted, which permits a more favorable endowment of methane and ammonia ices (in addition to water ice), for the development of an atmosphere (Hunten et al., 1984); (iii) Titan became massive enough to have undergone some internal differentiation which concentrated ices towards the surface leading to substantial outgassing, but still had enough gravity to retain all but the lightest gases; (iv) finally, the surface temperature on Titan (as opposed to that on Triton) is high enough to maintain significant amounts of CI'I4 and NI,I3 gas in equilibrium with their corresponding surface ices. Titan's origin and evolution is another aspect which is still treated mainly in form of questions and models. Using the physical processes operating in the present-day environment (such as radiation-driven chemistry, atmospheric escape, photochemistry, surface-atmosphere coupling etc.) to "run the dock back" would help model the origin and evolution of Titan.



The fabulous scientific adventure of the Voyager missions, which ended recently after a 12year Odyssey in the solar system, is at the present time being pursued by a mission to the Jovian system (Galileo) and will be extended in the next decade by a mission to Saturn and Titan (Cassini), currently under development. Planned to be launched in 1996 toward the saturnian system, the Cassini spacecraft will be composed of an orbiter and a descent probe. The instruments on board this mission will acquire improved (spectrally and spatially) time-dependent measurements of Titan which will enable us to complete our knowledge of the satellite. For example, CIRS, the infrared spectrometer will acquire a large number of data during the ~ 30 fly-by planned for the orbiter, with a spectral resolution (0.5 cm -1) an order of magnitude better than IRIS, a better

Titan: Recent Developments


sensitivity (in the 600-1400 cm -1 region) and a vertical resolution improved by a factor of 2-3 over the Voyager nadir soundings. Also, the closest approach of Cassini with respect to Titan is planned to be at a cUstance of ~ 1000 km from the surface: 6 times closer than Voyager. Other instruments will concentrate on the temperature-pressure-density profiles, secure with precision the atmospheric composition as a function of height, collect samples, explore the surface conditions, take images, etc. As concerns the composition, measurements of Titan's spectrum with higher resolution than the one actually available, might allow the detection of more complex species in Titan's atmosphere, the presence of which is implied from laboratory experiments. Cassini will no doubt bring a second revelation of Titan to us. In the meantime, ground-based measurements as well as observations made by satellites (such as ISO), will uncover, Little by little, some of the satellite's remaining mysteries.

Acknowledgments The author wish~s to thank Bruno B6zard and Th6r6se Encrenaz for critically reading the manuscript and for offering valuable remarks. Also, some of the figures were elaborated by Josiane Tualy.



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