1

The origin and evolution of Titan

G. TOBIE, J. I. LUNINE, J. MONTEUX, O. MOUSIS, AND F. NIMMO

1.1 Introduction

Although Titan is similar in terms of mass and sizeto Jupiter’s moons Ganymede and Callisto, it is dif-ferent in that it is the only one harboring a massiveatmosphere. Moreover, unlike the Jovian system, whichis populated with four large moons, Titan is the onlylarge moon around Saturn. The other Saturnian moonsare much smaller and have an average density at least25 percent less than Titan’s uncompressed density andmuch below the density expected for a solar composi-tion (Johnson and Lunine, 2005), although with a largevariation from satellite to satellite. Both Jupiter’s andSaturn’s moon systems are thought to have formed in adisk around the growing giant planet. However, the dif-ference in architecture between the two systems proba-bly reflects different disk characteristics and evolution(e.g., Sasaki et al., 2010), and, in the case of Saturn,possibly the catastrophic loss of one or more Titan-sized moons (Canup, 2010). Moreover, the presence ofa massive atmosphere on Titan, as well as the emis-sion of gases from Enceladus’ active south polar region(Waite et al., 2009), suggest that the primordial build-ing blocks that comprise the Saturnian system wereprobably more volatile-rich than those of Jupiter.

The composition of the present-day atmosphere,dominated by nitrogen, with a few percent methaneand lesser amounts of other species, probably doesnot directly reflect the composition of the primordialbuilding blocks and is rather the result of complex evo-lutionary processes involving internal chemistry andoutgassing, impact cratering, photochemistry, escape,crustal storage and recycling, and other processes. Thelow 36Ar/N2 ratio measured by the Gas Chromatograph

Mass Spectrometer (GCMS) on Huygens (Niemannet al., 2005, 2010) suggests that nitrogen was broughtto Titan not in the form of N2, but rather in the form ofNH3 (Owen, 1982). The argument goes as follows: Arand N2 have similar volatility and affinity with waterice. Thus, if the primary carrier of Titan’s N2 was molec-ular nitrogen itself, the Ar/N2 ratio on Titan should bewithin an order of magnitude of the solar composi-tion ratio of about 0.1 (Lunine and Stevenson, 1985),which is about 500,000 times larger than the observedratio. This indicates that nitrogen has been brought inthe form of easily condensable compounds, most likelyammonia, which then has been converted in situ byimpact-driven chemistry (McKay et al., 1988; Sekineet al., 2011; Ishimaru et al., 2011) or by photochemistry(Atreya et al., 1978). As will be further discussed in thischapter, both conversion processes required restrictiveconditions to occur, and therefore dictated specific evo-lutionary scenarios for Titan.

Photochemical processes are also known to lead toan irreversible destruction of methane, implying that asource of replenishment must exist to explain its fewpercent abundance (vertically averaged) in the atmo-sphere. In the absence of such a replenishment, theatmospheric methane should disappear over a timescaleof several tens of millions of years. Before the Cassini-Huygens mission, the atmospheric methane was pro-posed to be in equilibrium with a global surface liquidreservoir (e.g., Lunine et al., 1983). Even though liq-uid reservoirs have been identified on Titan’s surfaceby Cassini in the form of lakes and seas at high lat-itudes (Stofan et al., 2007), their estimated total vol-ume is too small to sustain methane in the atmosphereon geologic timescales (Lorenz et al., 2008a). This

29

30 The origin and evolution of Titan

suggests that a subsurface reservoir of methane existsand replenishes the atmospheric reservoir episodicallyor continuously. Different sources have been proposed,but despite a number of new constraints from theCassini-Huygens mission, we lack a strong discrim-inant among them. The 12C/13C ratio measured inmethane by the GCMS on Huygens (Niemann et al.,2010) is similar to the 12C/13C ratio measured in sev-eral solar system objects, and close to the solar value,possibly reflecting an atmospheric methane inventorythat has not been affected by atmospheric escape. Theanalysis of Mandt et al. (2009) indicates that if the cur-rent atmospheric methane is the remnant of methaneinjected into the atmosphere more than about two hun-dred million years ago, the 12C/13C should be decreasedby at least 10 percent compared with the initial value.The current atmospheric methane may then have beenrecently delivered to the surface, from whatever reser-voir. The detection of a significant amount of 40Ar(the decay product of 40K initially contained in sili-cate phases) further indicates that exchanges with therocky interior have occurred (Niemann et al., 2010).

In this chapter we argue that exchange betweenthe interior and atmosphere of Titan is likely to haveoccurred throughout Titan’s evolution, at least duringaccretion and in the following billion years, and pos-sibly until the present. A better understanding of theseexchange processes requires good knowledge of theinterior structure. Based on its mean density of 1881kg.m−3, Titan has long been known to be composedof a roughly 60–40 percent mixture of rock and ice bymass. Gravimetric and altimetric data collected by theCassini spacecraft permit us to constrain the distribu-tion of these components in the interior. However, as isfurther discussed in Section 1.2, the solution remainsnon-unique and the different possible structures of theinterior imply a range of formation and evolution sce-narios. Other data from the Cassini-Huygens missionsuggest the presence of a water ocean, possibly dopedwith ammonia or other agents depressing the oceancrystallization point, beneath the solid icy shell. Theexistence of such an internal water ocean at present hasimportant consequences for the evolution of Titan aswell as for its astrobiological potential.

In Section 1.2, we review the present-day constraintson the structure and dynamics of Titan’s interior andpresent different possible interior models. The present-

day composition and structure of the surface and atmo-sphere are already discussed in Chapters 2 through 5and will not be repeated here. A good knowledge ofthe present-day state of Titan is a necessary prelimi-nary to understanding its origin and evolution. Section1.3 presents some theoretical constraints on the forma-tion of the Saturnian system and describes the accretionprocess(es) of Titan. Coupled evolution of the interior,surface, and atmosphere, including internal differenti-ation and generation and recycling of the atmosphere,are then discussed in Section 1.4. Finally, we concludein Section 1.5 by listing a series of remaining questionsand how they may be addressed by future explorationmissions.

1.2 Present-day interior structure and dynamics

Models of Titan’s formation (Section 1.3) and subse-quent evolution (Section 1.4) are constrained by itspresent-day interior structure and dynamics. In this sec-tion, we focus on three questions that are of particularrelevance to Titan’s history.

1. To what extent is Titan differentiated? The degree of dif-ferentiation depends mainly on the thermal conditionsduring and just after accretion. Thus, Titan’s degree ofdifferentiation constrains the conditions under which itformed and evolved during the first billion years.

2. Is the ice shell convecting? The long-term thermal evolu-tion of Titan’s interior is determined by the rate at whichheat can be transferred across its icy outer layer(s). Thus,whether or not convection is taking place controls howrapidly Titan cools. If an ocean is present, the lifetime ofthe ocean (and thus its habitability) is controlled by thebehavior of the overlying ice shell. The ice shell also con-trols transport between the interior and the atmosphere,as well as potential tectonic and cryovolcanic activities.

3. Does Titan possess a subsurface ocean? If an oceanis present, Titan’s astrobiological potential is increased.Moreover, an ocean can have profound geophysical con-sequences; for instance, it will likely increase the rate oftidal dissipation in the overlying ice shell.

Prior to the Cassini-Huygens mission, there wereessentially no observational constraints on any of thesequestions. In the following sections we discuss the datasets relevant to the internal structure and dynamics ofTitan.

1.2 Present-day interior structure and dynamics 31

Table 1.1 Observed (Clm) un-normalized degree-twogravity coefficients

Clm × 10−6 Clm × 10−6 Chlm

SOL1 SOL2 m

l = 2, m = 0 −31.80 ± 0.80 −33.46 ± 1.26 −358l = 2, m = 2 9.98 ± 0.08 10.02 ± 0.14 63ratio −3.19 ± 0.08 −3.34 ± 0.13 −45.68

k f2 (C20) 0.965 1.015 –

k f2 (C22) 1.010 1.013 –

C/MT R2T (C20) 0.335 0.342 –

C/MT R2T (C22) 0.341 0.342 –

Observed values are from Iess et al. (2010), where SOL1and SOL2 refer to the multi-arc and global solutions, respec-tively, and errors are given as ±2σ . Un-normalized surfacetopography coefficients Ch

lm are derived from data detailedin Zebker et al. (2009b) and quoted from Nimmo and Bills(2010). Normalized moment of inertia C/MT R2

T was calcu-lated assuming hydrostatic equilibrium using equation 1.3.

1.2.1 Interior structure inferred fromthe gravity field data

Measurements of Titan’s gravity field are made byprecise radio tracking of the Cassini spacecraft as itpasses close to Titan. Unfortunately, because of the lackof a scan platform, gravity tracking precludes the acqui-sition of most other observations; therefore, the numberof fly-bys dedicated to these measurements is limited.Furthermore, Titan’s extended atmosphere can resultin nongravitational forces on the spacecraft, whichcomplicates the analysis and requires that the clos-est approach distance is farther from Titan than whatis desirable for an optimal resolution of the gravityfield. Despite these obstacles, Titan’s degree-two grav-ity coefficients have been determined with relativelyhigh confidence, as well as the degree-three coefficientsbut with much less confidence (Iess et al., 2010). Thedegree-two spherical harmonic coefficients, which arethe dominant coefficients for a body subjected to tidaland rotational forces, J2 = (−C20) and C22, are tabu-lated in Table 1.1. Because of the number of fly-bys,the coefficients were determined independently, with-out requiring a value of 10/3 for the ratio of J2 to C22.This situation differs from bodies where fewer fly-byshave been carried out, such as Callisto (Anderson et al.,2001) and Rhea (Mackenzie et al., 2008).

The degree-two coefficients provide very importantconstraints on the radial mass distribution within thebody. It can be shown mathematically that they aredirectly related to the principal moments of inertia(MoI), A, B, and C :

Cobs20 = (A + B)/2 − C

MT a2; Cobs

22 = B − A4MT a2

, (1.1)

with MT Titan’s mass and a the mean equatorial radius.Thus, measuring the C20 and C22 coefficients gives onlythe difference between the principal moments of iner-tia. To determine their absolute value, a third conditionis needed. For Mars and Earth, measurement of theprecession rate (which gives (C − A)/C , where C andA are the polar and averaged equatorial moments ofinertia) allows retrieval of the exact values of the threeprincipal moments of inertia. In the case of Titan, sucha measurement is not available. To infer the momentof inertia, one approach is to use the Radau-Darwinapproximation, which assumes that the body respondsto tidal and rotational forces as a fluid body (i.e., ithas no long-term strength, and therefore it is in hydro-static equilibrium). This approximation assumes thatthe body surface, as well as each interface betweeninternal layers, corresponds to an iso-potential surface.The potential at each interface is the sum of the tidalpotential, centrifugal potential, and gravitational poten-tial. If the body is in hydrostatic equilibrium, the degree-two gravitational potential is proportional to the tidaland centrifugal potential. The response of the body tothe tidal and centrifugal degree-two potential is clas-sically described with the fluid Love number k f

2 . Inhydrostatic equilibrium,

Cth20 = −

(α

2+ q

3

)k f

2 ; Cth22 = α

4k f

2 , (1.2)

where α = MS/MT × (RT /DST )3 is the tidal param-eter, and q = ω2 R3

T /G MT is the rotational parameter(with MS Saturn’s mass, RT Titan’s mean radius, DST

Saturn–Titan mean distance, ω Titan’s spin rate). In thecase of a synchronously rotating body in a circular orbit,α = q and therefore C20/C22 = −10/3 if the body ishydrostatic. For Titan, α = q = 3.9555 × 10−5.

In the Radau-Darwin approximation, the normalizedmoment of inertia about the spin axis, C/MT R2

T , canthen be calculated from the fluid Love number, inferredfrom the gravity coefficients assuming that they result

32 The origin and evolution of Titan

from a fluid response of the body and in the limit ofsmall density variations:

CMT R2

T= 2

3

⎡

⎣1 − 25

(5

(k f2 + 1)

− 1

)1/2⎤

⎦ . (1.3)

In the case of Titan, the ratio between the observedCobs

20 and Cobs22 is close to −10/3, suggesting that Titan

is relatively close to hydrostatic equilibrium. We cantherefore use the degree-two gravity coefficients to firstorder to estimate the fluid Love number and the corre-sponding moment of inertia. Depending on whether weuse the C20 or C22 coefficients as a reference to deter-mine the fluid Love number, the normalized momentof inertia varies between 0.335 and 0.342. This maybe compared with the value of 0.4 for a uniformsphere, 0.394 for the Moon, 0.33 for the Earth, 0.31for Ganymede, and 0.355 for Callisto (assuming hydro-static equilibrium [Anderson et al., 1996, 2001]). Themoment of inertia factor of Titan is therefore inter-mediate between Ganymede and Callisto, suggesting amodest increase of density toward the center.

As already mentioned, this interpretation, however,assumes that the Cobs

20 or Cobs22 coefficient is entirely

determined by the dynamical distortion of the satelliteunder the action of tidal and rotational forces. In reality,part of the gravity signal may be attributed to nonhydro-static contributions such as uncompensated relief, iceshell thickness variations, or lateral density variations.The fact that the Cobs

20 /Cobs22 ratio is not exactly −10/3

and that the degree-three gravity coefficients are notzero clearly indicates that there are some nonhydrostaticcontributions to the observed signals and that the grav-ity field of Titan is not attributed only to the dynamicaldistortion of a radially distributed mass interior. Thedegree-three coefficients indicate that non-negligiblelateral variations in the mass distribution exist on Titan.These should affect the estimation of the moment ofinertia. If we assume that the nonhydrostatic contri-butions to degree two are comparable to the inferreddegree-three terms, a few to 10 percent of the observeddegree-two signals may be attributed to the nonhydrostatic (Iess et al., 2010, Supplementary Informa-tion). According to the Radau-Darwin equation, a 5–10percent overestimate of the hydrostatic parts would shiftC/MR R2

T from 0.341 to 0.334–0.327.

In summary, the interpretation of the gravity dataindicates that the normalized moment of inertia of Titanis of the order of 0.33–0.34. This MoI value combinedwith the mean density of Titan indicates that the densityof the core is low, which may be explained by incom-plete separation of rock from ice (Iess et al., 2010)or the presence of highly hydrated silicate minerals(Castillo-Rogez and Lunine, 2010; Fortes, 2012). Basedon the present data, the presence of a small iron core(<500 km) cannot be ruled out (Castillo-Rogez andLunine, 2010; Fortes, 2012), but the low internal tem-peratures (<900–1000 K) required to maintain eitheran ice-rock mixture or highly hydrated silicate mineralsare incompatible with the segregation of iron and theformation of a iron core. The most likely present-daystructure is either a pure rock core, mostly anhydrous,surrounded by a mixture of ice and rock, surroundedby pure water ice (Fig. 1.5, top and middle), or a low-density rock core consisting mainly, but not entirely,of hydrated silicate minerals surrounded by H2O(Fig. 1.5, bottom). In both cases, a liquid water layermay be present between a high-pressure ice mantle andan outer ice shell. We will discuss the possible evo-lution scenarios leading to these interior structures inmore detail in Section 1.4.

1.2.2 Global shape: constraints on the structureand thermal state of the ice shell

Just as with gravity, the degree-two shape of a syn-chronous body can be used to determine its moment ofinertia – as long as the body is in hydrostatic equilib-rium. As with gravity, the degree-two shape coefficientsof such a body, Ch

20 and Ch22, will be in the ratio −10/3.

Thus, with sufficiently good topography, we can checkwhether or not the hydrostatic condition is satisfied.

Titan’s shape has been measured (Zebker et al.,2009b) by a combination of traditional radar altime-try (Zebker et al., 2009a) and a new technique thatuses overlaps in the Synthetic Aperture Radar (SAR)-imaging swaths to determine topography (Stiles et al.,2009). In contrast to the results obtained from gravity,Ch

20 and Ch22 were not in the expected −10/3 ratio (see

Table 1.1).This result is apparently paradoxical: How can

Titan have nearly hydrostatic gravity but nonhydrostatictopography? The hydrostatic gravity rules out scenarios

1.2 Present-day interior structure and dynamics 33

in which Titan’s shape is derived from an earlier epoch(a “fossil bulge” similar to that of the Moon). One pos-sibility is that Titan’s outer ice shell – presuming theexistence of an ocean – is of varying thickness or den-sity and is (Airy or Pratt) isostatically compensated.An isostatic shell would give rise to only very smallgravity anomalies (discussed later), whereas the shellthickness variations would result in a long-wavelengthnonhydrostatic shape.

Nimmo and Bills (2010) explored a version of theAiry scenario, in which shell thickness variations werecaused by spatial variations in tidal heating in the iceshell (cf. Ojakangas and Stevenson, 1989). They foundthat the observed topography could be fitted with amean ice shell of thickness D ≈100 km, heated frombelow at a rate of 4–5 mW m−2 (appropriate to a chon-dritic Titan) and from within at a rate about one-tenthas large (due to tidal dissipation). The tidal Love num-ber h2, which describes the response of the body to theeccentricity tides, was found to be ≈1.2, consistent withtheoretical predictions (Sohl et al., 2003). Given the rel-atively modest rate of tidal dissipation, the e-foldingtimescale to damp Titan’s eccentricity is quite long(2–3 Gyr). Alternatively, Choukroun and Sotin (2012)proposed that the flattened shape of Titan may resultfrom the accumulation of dense hydrocarbon clathratesat high latitudes owing to enhanced ethane precipita-tion. Their Pratt isostatic compensation model explainsthe ∼300 m difference between expected and measuredpolar radii if the ethane-rich clathrate crust is about3 km thick.

Although the variable-thickness shell is assumed iso-static, it will still have a small effect on the degree-twogravity coefficients. Correcting for this effect beforeusing equation (1.3) changes the resulting moment ofinertia by about 2 percent (the MoI factor is still between0.33 and 0.34). An isostatic shell is probably a goodassumption: a shell with an elastic thickness of 50 kmwould still be about 95 percent compensated at degreetwo (Turcotte et al., 1981), assuming a Young’s modulusof 9 GPa.

Perhaps the most surprising result of these studies isthat, if the variable-thickness model is correct, Titan’sice shell must be conductive, or at most weakly convect-ing, to support thickness variations in the long term. Inthe thickness-variation model, surface topography gen-erates pressure gradients at depth; if the ice shell were

convecting, the ice viscosity would be low enough toflow in response to these pressure gradients and pro-gressively remove the surface topography (Stevenson,2000). Whether or not an ice shell convects dependson whether the Rayleigh number Ra of the whole icelayer exceeds some critical value Racr (e.g., Mitri andShowman, 2008):

Ra = ρgα%T D3

κηb

≈ 3 × 108(

D100 km

)3(

1014 Pa sηb

)

≥ Racr . (1.4)

Here, ρ, α, κ , ηb are the density, thermal expan-sivity, thermal diffusivity, and basal viscosity of theice shell, respectively; %T is the temperature contrastacross the shell (%T ≃ 170–180 K between the surfaceand the ocean); and numerical values assumed are fromNimmo and Bills (2010). The value of Racr dependson how sensitive viscosity is to temperature changes,and the geometry of the shell (Solomatov, 1995; Barrand McKinnon, 2007), and for Titan is ≈4 × 107 in thestagnant lid regime. Equation (1.4) shows that whetheror not the shell convects is a sensitive function of theviscosity of the base of the ice shell. For pure ice, theviscosity in turn depends mainly on the temperatureof the ocean beneath, and secondarily on the ice grainsize. The ocean temperature is controlled by how muchammonia is present; a sufficiently ammonia-rich oceanand/or large ice grain sizes would prevent convectionbecause the ice would be too viscous (e.g., Tobie et al.,2005b).

Nevertheless, it is possible that the ice shell wasconvective in the past and that convection ceased asthe ocean crystallized and the ammonia concentrationincreased owing to reduced ocean volume (Tobie et al.,2005b; Mitri and Showman, 2008). The elevated orbitaleccentricity of Titan indicates, however, that this con-vective period should have lasted less than about 0.5 to1 billion years (Tobie et al., 2005b). In contrast to theGalilean moon where the Laplace resonance excites theeccentricity of Europa and Io, there is no major reso-nance in the Saturnian system to force Titan’s eccentric-ity. Unless some unknown process has recently forcedthe eccentricity, the latter should have monotonicallydecayed as a function of time since the formation ofthe system owing to tidal dissipation at the surface and

34 The origin and evolution of Titan

in the interior (Sears, 1995; Sohl et al., 1995; Tobieet al., 2005b). The dissipation rate is reduced either ifthe interior is totally rigid (with no internal ocean) (Sohlet al., 1995) or if the outer ice shell above the ocean isthin and/or cold (i.e., conductive) (Tobie et al., 2005b).Indeed, the volume wherein tidal energy is dissipated ismuch bigger in a convective ice shell (where about halfof the shell has a temperature higher than 240 K) thanin a conductive ice shell, where the warm temperatureis confined at the bottom of the ice shell.

If Titan has an internal ocean at present, it impliesthat the ice shell remained conductive and relativelythin (<30–40 km) during most of Titan’s evolution, toreduce the eccentricity damping (Tobie et al., 2005b).This is possible if the viscosity of the icy shell is largeenough due to large grain sizes or to low oceanic tem-peratures (owing to ammonia [Tobie et al., 2005b] ormethanol [Deschamps et al., 2010]) and/or if the icyshell contains a large fraction of low-conductivity mate-rials such as methane clathrate (Tobie et al., 2006). Fol-lowing this requirement, the icy shell has never beenconvecting or has been convecting only during a limitedperiod of time. A possible scenario may be that the iceshell started convecting during roughly the last billionyears and then stopped later because of ocean crystal-lization as the concentration in ammonia or methanolin the ocean became too high and the temperature toolow to allow convection in the overlying viscous iceshell (e.g., Mitri and Showman, 2008). Indeed, even foran initially moderate ammonia fraction (about 2–3%relative to the total mass of H2O), the concentrationincreases significantly as the ocean crystallizes and theocean temperature gets below 250 K when the ammo-nia concentration in the crystallizing ocean reaches10 wt% (Tobie et al., 2005b). At such a tempera-ture, the viscosity at the base of the ice shell exceeds5 × 1015–1016 Pa.s, which is too large to promote con-vection even for a 100 km-thick ice shell.

1.2.3 Geophysical evidence for an internalwater ocean

For the Galilean moons, the most convincing evidencefor subsurface oceans came from analysis of the mag-netic data acquired by the Galileo mission during closefly-bys. The tilt of Jupiter’s magnetic field relative to itsspin pole means that the Galilean satellites experience

a time-varying field and resulting induction currents.This effect has been used to probe the conductivitystructure of the satellites, and deduce the presence ofoceans for the three outer satellites (Khurana et al.,1998; Kivelson et al., 2002), and a partially moltenmantle for Io (Khurana et al., 2011). Unfortunately,Saturn’s spin and magnetic poles are almost exactlycoplanar, so the induction technique is very difficult touse at Titan (Saur et al., 2010). A recent study of Titan’smagnetic field did not see any induction effects and didnot detect any permanent internal magnetic field (Weiet al., 2010). Although it is possible to have a moltencore without generating a magnetic field (e.g., Venus,Io), this result is perhaps consistent with Titan’s interiorbeing cold and not fully differentiated.

On Titan, another technique based on the measure-ments of the electric field during the Huygens descentby the Permittivity, Wave, and Altimetry (PWA) instru-ment (Beghin et al., 2007) provides, however, someinformation on the subsurface electrical conductivity.The measurements of low-frequency waves and atmo-spheric conductivity by the PWA instrument revealedthe existence of a Schumann-like resonance trappedwithin Titan’s atmospheric cavity. The observed sig-nal may be triggered and sustained by strong electriccurrents induced in the ionosphere by Saturn’s mag-netospheric plasma flow (Simoes et al., 2007; Beghinet al., 2009, 2010). The characterisitics of this trapped-resonant mode suggest the presence of a conductivelayer at about 30 to 60 km below the surface (Beghinet al., 2010), which may correspond to the ice/oceaninterface provided that the ocean is sufficiently conduc-tive, possibly doped with small amounts of ammoniaand/or salts (Beghin et al., 2010).

As explained in Section 1.2.2, the observed topogra-phy may be explained by variations in ice shell thick-ness, resulting from inhomogeneous crystallization ofan underlying ocean. If this analysis is correct, it impliesthat an internal ocean is still present in Titan’s interiorand that it is currently slowly crystallizing. This consti-tutes additional indirect evidence for an internal oceaninside Titan.

Another piece of evidence is provided by Titan’s spinstate. A nonspherical body, such as a synchronous satel-lite, precesses about its rotation axis at a rate determinedby the ratio (C − A)/C . Because the gravity coefficientJ2 = (C − A)/MT R2

T , if J2 and the precession rate can

1.2 Present-day interior structure and dynamics 35

be measured, the moment of inertia C may be deter-mined directly without assuming hydrostatic equilib-rium. Although measuring a satellite’s precession ratedirectly is difficult, measuring the obliquity (the anglebetween spin pole and orbit pole) is relatively easy. If asatellite is dissipative, tidal torques drive the obliquityto a state in which the spin pole and orbit pole remaincoplanar as the latter precesses around a fixed (invari-able) pole (e.g., Peale, 1969). If this state, known as adamped Cassini state, exists, then measuring the obliq-uity is sufficient to determine the satellite moment ofinertia (e.g., Bills and Nimmo, 2011). In practice, othersmall effects (such as Saturn’s spin pole precession)have to be taken into account, but they do not alter theunderlying physical picture.

Titan’s spin pole location has been establishedthrough analysis of SAR images. It is very nearly copla-nar with the invariable and orbit poles (suggesting occu-pation of a Cassini state), and corresponds to an obliq-uity of 0.32◦ (Stiles et al., 2008, 2010). If Titan occupiesa damped Cassini state, the implied value of C/(MT R2

T )is 0.45 ± 0.02 (Bills and Nimmo, 2011). This result isclearly unreasonable, in that it suggests a body thatis under-dense in the interior and also disagrees withthe results derived from gravity. The solution is prob-ably a combination of two factors. First, the fact thatTitan’s spin pole is not quite coplanar with the orbitand invariable poles suggests some degree of present-day excitation of the obliquity. Two possible candidatesources of excitation are the atmosphere (Tokano et al.,2011) or fluid motion in a subsurface ocean, both ofwhich can transfer angular momentum to the ice shell.The second factor is that Titan is probably not acting asa rigid body (which the Cassini state analysis assumes),and that instead the ice shell is partially decoupled fromthe interior by an ocean. If the shell were totally decou-pled, its normalized moment of inertia would be 0.67.The derived value of 0.45 suggests some degree of cou-pling, perhaps due to pressure or gravitational torques(Baland et al., 2011). Fortunately, continued monitor-ing of the position of Titan’s spin pole should be ableto disentangle these two effects, but the current resultsare strongly indicative of a subsurface ocean.

Another potential source of evidence for a subsur-face ocean would be the existence of nonsynchronousrotation (NSR) of the ice shell (Ojakangas and Steven-son, 1989). It was initially thought that Titan’s shell was

exhibiting such behavior (Lorenz et al., 2008b), but thisobservation turned out to be erroneous and there is cur-rently no evidence for NSR (Stiles et al., 2010).

The most promising technique to confirm the exis-tence of a subsurface ocean would probably be to mea-sure the tidal fluctuation from accurate gravity or topog-raphy measurements. The Radio Science Experimenton Cassini has, in theory, the capability to measurethe tide-generated gravity signals (Rappaport et al.,1997, 2008). However, the atmospheric drag due to theextended atmosphere as well as the existence of rela-tively large signals in harmonics of degree three (Iesset al., 2010), suggesting significant signal at degreesfour and five (Sotin et al., 2009), make the analysisdifficult. In spite of these difficulties, a recent analysisof the Cassini gravity data obtained after six fly-byspermitted the estimation of the tidal Love number to0.6 ± 0.2 (Iess et al., 2012), providing further evidencefor a global ocean at depth on Titan (Iess et al., 2012). Afuture mission with dedicated multiple fly-bys or, evenbetter, in orbit around Titan will be required to confirmthis detection and to get more precise informations onthe ocean and ice shell thicknesses.

1.2.4 Summary

At the start of this section we posed three questions; wewill conclude by providing some provisional answers.

1. To what extent is Titan differentiated? Based on the grav-ity results (Section 1.2.1), the normalized moment of iner-tia of Titan is estimated between 0.33 and 0.34. Thisclearly indicates that Titan is not fully differentiated anddoes not have a substantial iron core, as does Ganymede.Earlier we suggested two possible internal structures –one in which rock and ice are incompletely separated, theother involving a rocky core that is partially hydrated –that satisfy the rather high moment of inertia. We discussthis further in Sections 1.3.4 and 1.4.1.1.

2. Is the ice shell convecting? A conductive ice shell isfavored by the model for Titan’s long-wavelength shape(Section 1.2.2) and is consistent with the requirementthat its eccentricity not damp too rapidly. To avoid con-vection, the ice shell must have a high viscosity, whichcan be achieved if the ocean beneath is ammonia-richand therefore cold. Convection may have occurred inthe past before the ammonia concentration increased, butthis convective period must have lasted less than 0.5 to

36 The origin and evolution of Titan

1 billion years to explain Titan’s present-day elevatedorbital eccentricity.

3. Does Titan possess a subsurface ocean? Electric signalsmeasured in Titan’s atmosphere by Huygens suggest thepresence of a conductive layer at depth, possibly a waterocean doped with small amounts of salts or ammoniato increase the electrical conductivity. Titan’s spin stateand tidal gravity response are also suggestive of a decou-pled shell (Section 1.2.3), whereas its long-wavelengthshape can be explained if a subsurface ocean is assumed(Section 1.2.2). Given the existence of subsurface oceanson the comparably sized Galilean satellites, Callisto andGanymede, it would be surprising if Titan did not pos-sess an ocean. However, one important difference is thatTitan’s ocean may be colder and more ammonia-rich(discussed later).

1.3 Origin of the Saturnian system andaccretion of Titan

Titan was formed as a regular satellite in a disk thatwas the outgrowth of the formation of Saturn itself.The conditions under which Saturn formed thereforehad a direct impact on the formation of Titan. Theydetermined its composition, and the architecture ofthe whole system, as well as the final assemblage thatled to the formation of Titan and of the other regularmoons. The present section is structured around threemain questions that are of particular relevance to Titan’sformation.

1. What was the composition of the primordial buildingblocks that formed Titan and the Saturnian system? Titanand the other Saturnian moons formed from a mixture ofrock and ice initially present in the vicinity of Saturn. Wepresent in this section different theoretical and observa-tional constraints that can be used to estimate the com-position of these ice-rock blocks, notably their volatilecontent.

2. What makes the Jupiter and Saturn systems so different?Contrary to the Jovian system, there is no gradual varia-tion of density, and hence rock/ice ratio, in the Saturnianmoons as a function of planet distance. Titan is the uniquebig moon, with the second largest moon, Rhea, having amass less than 2 percent of Titan’s mass. Another particu-lar characteristic of Saturn is its extended ring system. Inthis section, we describe the main phases of satellite for-mation in a disk around giant planets and try to identifywhich processes differed between Jupiter and Saturn.

3. Can Titan accrete undifferentiated? The Cassini grav-ity data suggest that Titan is moderately differentiatedat present (Section 1.2.1), suggesting that the differen-tiation process has been limited. One possibility is thatTitan accreted relatively cold, which would have delayedthe differentiation. To address that question, we reviewthe physics of accretion and determine under which con-ditions Titan might have remained cold during accretion.

1.3.1 Formation of ices in Saturn’s environment

The primitive building blocks that formed Titan and therest of the satellite system were presumably composedof a complex assemblage of various ices, hydrates, sili-cate minerals, and organics. An important issue associ-ated with the formation of giant planets and their satel-lite systems concerns the condensation of water ice andother ice compounds as the solar nebula cools down.The condensation sequence depends on the initial gasphase, the cooling rate of the disk, and the stability ofthe gas molecules in solid phases.

The composition of the initial gas phase of the diskcan be defined by assuming that the abundances of allelements, including oxygen, are protosolar (Asplundet al., 2009) and that O, C, and N exist only inthe forms of H2O, CO, CO2, CH3OH, CH4, N2, andNH3. The abundances of CO, CO2, CH3OH, CH4,N2, and NH3 are then determined from the adoptedCO/CO2/CH3OH/CH4 and N2/NH3 gas phase molecu-lar ratios. Once the abundances of these molecules arefixed, the remaining O gives the abundance of H2O.Concerning the distribution of elements in the mainvolatile molecules, the ratio of CO/CO2/CH3OH/CH4

can be set to 70/10/2/1 in the gas phase of the disk,values that are consistent with the interstellar medium(ISM) measurements considering the contributions ofboth gas and solid phases in the lines of sight (Mousiset al., 2009c). In addition, S is assumed to exist inthe form of H2S, with H2S/H2 = 0.5 × (S/H2)⊙, andother refractory sulfide components (Pasek et al., 2005).N2/NH3 is assumed equal to 10/1 in the nebula gasphase, a value compatible with thermochemical modelsof the solar nebula (Lewis and Prinn, 1980) and withobservations of cometary comae (e.g., Hersant et al.,2008).

To first order, the process by which volatilesare trapped in icy planetesimals can be calculated

1.3 Origin of the Saturnian system and accretion of Titan 37

assuming thermodynamic equilibrium. The sequenceof clathration1 and condensation can be determined byusing the equilibrium curves of stochiometric hydrates,clathrates, and pure condensates, and the thermody-namic path (hereafter cooling curve) detailing the evo-lution of temperature and pressure between roughly5 and 20 AU, a distance range largely encompass-ing the migration path followed by proto-Jupiter andproto-Saturn during their formation in the solar neb-ula (Alibert et al., 2005b). We refer the reader to theworks of Papaloizou and Terquem (1999) and Alibertet al. (2005a) for a full description of the turbulentmodel of accretion disk used here. For the pressurerange expected in the solar nebula, most of the gascompounds, except CO2, are enclathrated at tempera-tures higher than the temperatures they condense in theform of pure condensates. This has fundamental con-sequences for the trapping of volatiles in primordialbuilding blocks.

As illustrated in Figure 1.1, the cooling curve inter-cepts the equilibrium curves of the different ices at par-ticular temperatures and pressures. For each ice consid-ered, the domain of stability is the region located belowits corresponding equilibrium curve. The clathrationprocess stops when no more crystalline water ice isavailable to trap the volatile species. As CO2 is the onlyspecies that crystallizes at a higher temperature than itsassociated clathrate, solid CO2 is assumed to be the onlyexisting condensed form of CO2 in this environment. Inaddition, CH3OH is assumed to be in the form of pureice, as, to the best of our knowledge, no experimentaldata concerning the equilibrium curve of its associatedclathrate have been reported in the literature.

For simplicity, the clathration efficiency is assumedto be total, implying that guest molecules had the timeto diffuse through porous water-ice solids before theirgrowth into planetesimals and their accretion by proto-planets or proto-satellites. This statement remains plau-sible only if collisions between planetesimals haveexposed essentially all the ice to the gas over time scalesshorter than or equal to planetesimal lifetimes in thenebula (Lunine and Stevenson, 1985). In the present

1 Clathrate hydrates are crystalline water-based structures that form cagesin which gas molecules can be trapped (Lunine and Stevenson, 1985;Choukroun et al., 2013), and for this reason they play a crucial role in thevolatile trapping. They also play a key role in the evolution of the volatileinventory on Titan, as will be further discussed in Sections 1.4.1.3 and1.4.1.4.

Figure 1.1 Equilibrium curves of NH3-H2O hydrate,H2S, Xe, CH4, and CO clathrates (solid lines), CH3OH,CO2, Kr, CO, Ar, and N2 pure condensates (dotted lines),and thermodynamic path followed by the primordial neb-ula between 5 and 20 AU as a function of time, respec-tively, assuming a full efficiency of clathration. Speciesremain in the gas phase above the equilibrium curves.Below, they are trapped as clathrates or simply condense.The equilibrium curves of hydrates and clathrates derivefrom Lunine and Stevenson’s (1985) compilation of pub-lished experimental work, in which data are available atrelatively low temperatures and pressures. The equilib-rium curves of pure condensates used here derive from thecompilation of laboratory data provided by Lide (2004).

case, which is based on the protosolar abundances ofAsplund et al. (2009), NH3, H2S, Xe, CH4, and ∼31.5percent of CO form NH3-H2O hydrate and H2S, Xe,CH4 and CO clathrates with the available water in theouter nebula. The remaining CO, as well as N2, Kr, Ar,and Ne (not shown on the figure), whose clathration nor-mally occurs at lower temperatures, remain in the gasphase until the nebula cools enough to allow the forma-tion of pure condensates. As the gas phase compositionof the disk is assumed not to vary with heliocentric dis-tance, the calculated clathration conditions remain thesame in the 5–20 AU range of the nebula (see Marboeufet al. [2008] for details). Once crystallized, these iceswill agglomerate and form planetesimals large enoughto decouple from the nebular gas and will be accretedby the forming giant planets and satellites.

Many works published in the last decade and detail-ing the formation conditions of ices in the solar nebula(e.g., Gautier et al., 2001; Hersant et al., 2008; Mousisand Alibert, 2006; Alibert and Mousis, 2007) have used

38 The origin and evolution of Titan

Figure 1.2 Mole fraction of volatiles encaged in (a) H2S, (b) Xe, (c) CH4, and (d) CO clathrates. Gray and dark bars correspond to structureI and structure II clathrates, respectively, which are two possible clathrate structures with slightly different cage sizes (e.g., Sloan, 1998).

the equilibrium curves established experimentally forsingle guest molecules (similar to the curves shown onFigure 1.2). This simplified approach permits an under-standing to first order of the condensation sequencein the nebula, but in reality, clathrates forming in agas mixture always simultaneously trap several guestmolecules (e.g., Lunine and Stevenson, 1985; Sloan,1998). Here we consider the effect of simultaneous trap-ping of guest species on the predicted composition ofvolatiles trapped in the water ice.

The composition of the clathrate phases can be con-strained by using a statistical thermodynamic modelderived from the approach of van der Waals andPlatteeuw (1959) (see Lunine and Stevenson, 1985,and Mousis et al., 2010, for details). This thermody-namic approach is based on the use of intermolecu-lar potentials, which themselves depend on parametersdescribing the interaction between the guest molecule

and the cage made of water molecules, called Kiharaparameters. The aforementioned formation sequenceof the different ices occurring during the cooling ofthe solar nebula is considered (see Figure 1.1). Asa consequence, any volatile already trapped or con-densed at a higher temperature than the formationtemperature of the clathrate under consideration isexcluded from the coexisting gas phase composition.This implies that CO2, Xe, CH4, CO, Kr, Ar, andN2 are considered as possible guests in the case ofH2S clathrate. On the other hand, only N2, Ar, and Krcan become guests in CO clathrate. Figure 1.2 rep-resents the mole fraction f of volatiles encaged instructure I and structure II clathrates a priori domi-nated by H2S, Xe, CH4, and CO and formed in theprimordial nebula. This figure shows that H2S, Xe,CH4, and CO are the dominating guest species if struc-ture I clathrates form in the nebula. However, in the

1.3 Origin of the Saturnian system and accretion of Titan 39

case of structure II clathrate formation, CO becomes aminor compound in the clathrate that is expected to bedominated by this molecule. Indeed, Kr becomes moreabundant than CO in this clathrate, irrespective of theirinitial abundances in the gas phase of the nebula. On theother hand, structure I should be the most likely formof clathrate produced in the nebula because H2S, Xe,CH4, and CO clathrates are all of structure I (Lunine andStevenson, 1985). Figure 1.2 also shows that, depend-ing on the considered structure, the Kr/CH4 ratio inCH4–dominated clathrate is on the order of or greaterthan the value derived from the nebula gas phase com-position used in our model. This implies that Kr can beenclathrated at a temperature of ∼55 K instead of con-densing in the 20 to 30 K range in the solar nebula (seeFigure 1.1).

1.3.2 Formation of Saturn and the circumplanetarydisk: consequences for the delivery of volatiles

1.3.2.1 Saturn’s formation and envelope enrichment

Favored theories of giant planet formation centeraround two main paradigms, namely the core accretionmodel (Safronov, 1969; Goldreich and Ward, 1973; Pol-lack et al., 1996) and the gravitational instability model(Cameron, 1978; Boss, 1997). In the frame of the coreaccretion model, a solid core forms from the accretionof planetesimals and becomes massive enough (∼5–10M⊕) to initiate runaway gravitational infall of a largegaseous envelope in which gas-coupled solids continuetheir accretion (Alibert et al., 2005b; Hubickyj et al.,2005; Mordasini et al., 2009). This model provides thelarge amount of heavy elements necessary to explainthe supersolar metallicities observed in Jupiter and Sat-urn via the accretion of planetesimals in their envelopes(Gautier et al., 2001; Saumon and Guillot, 2004; Alib-ert et al., 2005b; Mousis et al., 2006, 2009c). In theframe of the gravitational instability model, gas giantprotoplanets form rapidly through a gravitational insta-bility of the gaseous portion of the disk and then con-tract to planetary densities more slowly (Boss, 1997,2005). In this scenario, owing to the limited efficiencyof planetesimals accretion during the planet formation,its metallicity should be almost equal to that of theparent star (Helled and Bodenheimer, 2010).

By using the sequences of ice formation describedin Section 1.3.1, it is possible to predict the abun-dances of major species (C, N, O, S, P) and noblegases (Ar, Kr, Xe) in Saturn’s atmosphere in the con-text of the core-accretion model (Hersant et al., 2008;Mousis et al., 2009c) and to compare to the observedabundances when available. For instance, adjusting themodel to the observed carbon abundance provides con-straints on the minimum mass range of ices requiredin the envelope (estimated between 14.3 and 15.6 M⊕[Mousis et al., 2009c]). Such an approach also providesinformation on the typical composition of the planetes-imals in the feeding zone of Saturn, which in turn canbe used to constrain the composition of the buildingblocks of Titan. The analysis of Hersant et al. (2008),based on a simplified description of volatile trappingby clathration and condensation, suggests that the plan-etesimals that reached Saturn’s envelope contained sig-nificant amounts of CO2, NH3, CH4, and H2S, andwere depleted in CO and N2. This also seems consis-tent with the composition of Enceladus’ plumes, whereCO2, NH3, CH4, and H2S have been detected by theCassini Ion and Neutral Mass Spectrometer (INMS)(Waite et al., 2009). If this is correct, it would implythat the solids that enrich Saturn’s envelope in heavy ele-ments have a composition close to the ones that formedthe regular satellites.

1.3.2.2 Structure and evolution of Saturn’s subnebula

At the beginning of Saturn’s formation process, the totalradius of the planet equals its Hill’s radius, and no sub-nebula can exist. At the end of the formation process,the cooling rate of the planet’s envelope is such that itsradius shrinks faster than the disk can supply additionalgas. Accretion of gas now proceeds through stream-ers that eventually collide and form a circumplanetarydisk (Lubow et al., 1999). When the radius of proto-Saturn becomes small enough, a subnebula emergesfrom the contracting atmosphere. This subnebula is fedby gas and solids accreted from Saturn’s feeding zone.The accretion rate of gas from the solar nebula to thesubnebula is therefore directly driven by the planet’saccretion rate of gas, which is derived from giant planetformation models (e.g., Alibert et al., 2005a).

Different models of subdisks have been proposed inthe literature (e.g., Canup and Ward, 2002; Mosqueiraand Estrada, 2003a; Estrada and Mosqueira, 2006;

40 The origin and evolution of Titan

Figure 1.3 Temperature inside the subnebula as a func-tion of the distance from Saturn (expressed in Saturn’sradii RSat ) and represented at different epochs of its evo-lution (figure adapted from Alibert and Mousis [2007] –see text).

Sasaki et al., 2010). Here we describe in more details theso-called gas-starved disk model initially introduced byCanup and Ward (2002) and then used by Alibert et al.(2005a) and Alibert and Mousis (2007). Alibert andMousis (2007) described a model of Saturn’s subneb-ula whose evolution proceeds in two phases. Duringthe first phase, when the solar nebula is still present,the subnebula is hot and dense, and fed through itsouter edge (fixed to 1/5 Hill’s radius – see Alibertand Mousis [2007] for details), at a rate controlled bySaturn’s formation. This justifies a posteriori the use ofan equilibrium model as an initial model of the sub-nebula. This phase lasts until the solar nebula has dis-appeared and the temporal evolution of the subnebulais therefore directly governed by the last phase of Sat-urn’s formation. The accretion rate as a function of thedistance to Saturn is no longer constant and the temper-ature and pressure of the subnebula decrease rapidly.

Figure 1.3 represents the temporal evolution ofthe temperature profile within the Saturn’s subnebuladepicted by Alibert and Mousis (2007). The solid linesare plotted, from top to bottom, at t = 0 Myr (corre-sponding to the time when Saturn has accreted 70% ofits final mass), 0.1 Myr, 0.2 Myr, 0.3 Myr, 0.35 Myr,0.4 Myr, 0.5 Myr, 0.6 Myr and 0.7 Myr. The switch fromphase 1 to phase 2 of the subnebula evolution occurs at

0.36 Myr. The dotted line gives the temperature 2000years after the beginning of phase 2, and illustrates thefast evolution of the subnebula during this stage.

1.3.3 Delivery of solids and conditionsof Titan’s formation

Regardless of the considered formation mechanism ofSaturn’s satellites within the subnebula (accretion attheir present orbit or migration and growth startingfrom the edge of the subnebula), the delivery of solidsto the disk is an essential step, because it determines theamount of volatiles that can be retained and ultimatelyincorporated in the interiors of these bodies. Differ-ent processes have been advocated (see, for instance,Estrada et al. [2009] and Canup and Ward [2009] for adetailed discussion): (1) direct transport of small par-ticles into the disk with the inflowing gas; (2) ablationand gas drag capture of planetesimals orbiting the Sunthrough the gas-rich circumplanetary disk; (3) breakup,dissolution, and recondensation of planetesimals in theextended envelope of the forming planet; and (4) colli-sional capture of planetesimals.

Mechanism (1) is valid for micron- to millimeter-sized particles that remain coupled to the subdisk’s gas,whereas mechanism (2) requires larger-sized particles(∼1 m) to allow their capture and inward drift due to gasdrag. Mechanism (3) is probably not valid for Saturn’ssubnebula because the D/H measurement performed atEnceladus by the Cassini spacecraft (Waite et al., 2009),which is close to the cometary value, suggests that thebuilding blocks of this satellite were formed in the solarnebula. In contrast, planetesimals formed via this mech-anism would have a protosolar D/H ratio. Mechanism(4) applies for gas-poor or no-gas environments andmay have taken place at epochs subsequent to the sub-nebula phase. In any case, the relative contribution ofeach of these mechanisms to the delivery of solids tothe disk is still under debate (see, for instance, Estradaet al. [2009] and Canup and Ward [2009]). In term ofvolatile delivery, each of these processes has very dif-ferent consequences. In mechanism (1), and to a lowerextent in mechanism (2), the particles may preserve thevolatile content they acquired in the solar nebula dur-ing the condensation sequence, whereas in mechanism(3) a large fraction of volatile would be likely releasedwhen incorporated into the disk. Even in mechanism

1.3 Origin of the Saturnian system and accretion of Titan 41

(4), a significant fraction may be lost during collisionsat high velocity.

As a consequence, the volatile content in the build-ing blocks of Titan may be significantly smaller than thecontent in the planetesimals initially formed in the solarnebula in the vicinity of Saturn. Even for small particlesentrained by the gas inflow, they could have been alteredonce embedded in the subdisk if they encountered gastemperature and pressure conditions high enough togenerate a loss of volatiles during their migration and ifthey remained relatively porous, which seems a reason-able assumption. Mousis et al. (2009b), for instance,favored this mechanism to explain the carbon monox-ide and argon deficiencies in the atmosphere of Titan.This scenario is based on the fact that CO-dominatedclathrate, CO, N2, and Ar pure condensates are notstable at temperatures greater than ∼50 K in the gasphase conditions of the nebula or the subnebula (seeFigure 1.1). This implies that if porous planetesimalsultimately accreted by Titan experienced intrinsic tem-peratures of ∼50 K during their migration and accre-tion in Saturn’s subnebula, they are expected to releasemost of their CO, Ar, Ne, and N2 while CH4 and NH3

are expected to remain trapped, or at least to partlyrecondense if released. Similarly, if planetesimals arecaptured by gas drag or collisional capture, the mostvolatile species, CO, N2, Ne, and Ar, should be releasedwhereas the most stable one should be preferentiallyretained.

Such a volatile loss scenario cannot, however,account for the deficiency of Titan’s atmosphere in Krand Xe because Kr is significantly trapped in CH4-dominated clathrate (see Figure 1.2) and Xe is alsoenclathrated at temperatures higher than ∼50 K. How-ever, as is further discussed in Section 1.4.2.2, bothXe and Kr may be sequestred in Titan’s icy shell andinterior owing to clathration process, partly dissolvedin the hydrocarbon seas of Titan (Cordier et al., 2010)or trapped in aerosols (Jacovi and Bar-Nun, 2008), thuspotentially explaining their deficiency while predictedto be present in significant amounts in the buildingblocks.

This scenario seems also problematic regarding thetentative detection of 22Ne by the Huygens GCMS(Niemann et al., 2010). According to this scenario,the building blocks should be strongly depleted in Newith 22Ne/36Ar ≪ 1, whereas the ratio estimated in the

atmosphere from the Huygens GCMS measurementsis close to 1 (Niemann et al., 2010). Note, however,that the GCMS detection of 22Ne is only tentative, andthe errors are still very large. In-situ measurements ofNe, Ar, Kr, and Xe abundances in Saturn’s atmosphereby a future probe, as well as in surface sampling onTitan, could constrain the origin of their unexpectedabundances in Titan’s atmosphere. This would help inunderstanding the process of noble gas trapping in theSaturnian system, thus providing key information onthe composition of Titan’s building blocks.

Another issue concerns the satellite growth from thesolids present in the disk. In a gas-rich disk, the growingsolid objects should migrate inward as a result of gasdrag for small objects and then of type I and/or II migra-tion for largest objects (∼1000 km) (e.g., Estrada et al.,2009; Canup and Ward, 2009). A growing satellite maysurvive if it grows large enough to open a gap in the disk,limiting the migration (Mosqueira and Estrada, 2003b),or if the gas disk dissipates on a timescale comparableto the satellite formation timescale (Canup and Ward,2002; Alibert et al., 2005a). In the case of Saturn, onlyone big moon has survived this migration process.

Nevertheless, several large moons may have existedoriginally and been lost before the disk dissipated(Canup and Ward, 2006; Sasaki et al., 2010). The differ-ence of satellite architecture between Jupiter and Saturnmay thus result from differences in disk evolution andsatellite migration (e.g., Sasaki et al., 2010). By consid-ering two different disk evolution/structure models forJupiter and Saturn, Sasaki et al. (2010) proposed thatin the case of Jupiter, the satellite migration may havestopped due to quick truncation of gas infall causedby Jupiter opening a gap in the solar nebula, whilethe subdisk was still relatively hot. In contrast, in thecase of Saturn, the gas infall would have never stoppedabruptly, and rather would have slowly decayed. Mostof the moons that formed collided with the planet, andonly one last big moon would have survived. This sur-viving satellite would have formed on long timescales(∼Myrs), owing to reduced mass infall. This is obvi-ously only one scenario among others, and several alter-native scenarios may be envisioned.

The accretion by Saturn of early formed satelliteswith sizes analogous to that of Titan may also providean explanation for the formation of the ring and theinner regular moons. Numerical simulations performed

42 The origin and evolution of Titan

by Canup (2010) suggest that the tidal removal of massfrom a differentiated, Titan-sized satellite as it migratesinward toward Saturn may be the origin of the giantplanet’s ring and inner moons. These simulations showthat planetary tidal forces preferentially strip materialfrom the satellite’s outer icy layers, whereas its rockycore remains intact and is lost by collision with theplanet. The result of these computations is a pure icering much more massive than Saturn’s current rings.As the ring evolves, its mass decreases and icy moonsare spawned from its outer edge with estimated massesconsistent with Saturn’s ice-rich moons interior to andincluding Tethys.

The latter scenario has been further investigated byCharnoz et al. (2011), who confirmed that all the moonsinterior to Titan’s orbit may have been formed by theseaccretion processes, if Saturn has a sufficiently lowdissipation factor Q (<2000). This latter condition isrequired to provide sufficient tidal expansion for thesatellites to move from the edge of a massive ring sys-tem to their present orbital position. They also showthat the accretion of the satellites from the outer edge ofsuch a ring system may explain the diversity of rock/iceratios among the mid-sized moons. Alternatively, thevariation in rock/ice fraction among the mid-sized satel-lites may be attributed to the stochastic nature of late-stage accretion, characterized by giant impacts, as sug-gested by Sekine and Genda (2012). Unfortunately, thedifferent proposed models lack strong observationaldiscriminants. Estimation of Saturn’s Q factor fromastrometric measurements is probably one of the mostconstraining observables. Preliminary analysis byLainey et al. (2011) indicates that the Q factor may belower than 2000. Future observations will be needed toconfirm this unexpectedly low dissipation factor, whichstill remains puzzling.

1.3.4 Accretion of Titan

There is some analogy between the formation of the ter-restrial planets and that of the icy moons. As describedearlier, icy moons probably accreted within the firstmillion years of the solar system history in circumplan-etary disks from aggregates of ice and rock (e.g., Canupand Ward, 2002; Estrada et al., 2009). Among the majoricy satellites, Ganymede and Callisto seem to exhibit astrong dichotomy in their degree of differentiation. This

has often been attributed to differences in the accretionprocesses. Titan may represent an intermediate case.However, the presence of its massive atmosphere sug-gests that Titan formed from different materials thandid its Galilean cousins, which may have significantconsequences for the final assemblage.

1.3.4.1 Assemblage sequence: role ofimpactor distribution

A number of models have been developed for accre-tional heating of icy moons (e.g., Schubert et al., 1981;Squyres et al., 1988; Coradini et al., 1989; Barr et al.,2010). Although some ingredients change from onemodel to another, all these models follow the sameapproach, which was initially proposed by Safronov(1978) and Kaula (1979) for the accretion of the Earth.The main differences between the different models con-cern the assumption for the impact energy partition-ing as well as for impactor size and velocity distribu-tions. The most recent study based on this approach wasdeveloped by Barr et al. (2010). In their approach, theyassumed that during the satellite growth the impactorsare small enough that the impact energy is deposited atshallow depth and is efficiently radiated away. In thiscondition, they showed that if Titan accreted slowlyfrom numerous cold and small embryos more than 5Myrs after CAI formation, the melting of the outerlayer may have been avoided and the undifferentiatedproto-core may extend up to the surface.

In reality, the size and velocity distribution of theimpactors have probably varied during the accretionphase. Two reservoirs of bodies may have contributedto the satellite growth: bodies in orbit around the planet,formed or captured within the circumplanetary disk,and bodies in orbit around the Sun, colliding directlywith the growing satellite (e.g. Squyres et al., 1988).The contribution of each of these two reservoirs prob-ably varied as the circumplanetary disk evolved. Whenthe accretion sequence started, centimeter-sized plane-tocentric particles were probably predominant, whereaskilometer-sized and larger bodies became more andmore frequent during the late-stage of accretion (seeFigure 1.4).

The temporal evolution of the size of impactorsplayed an important role in the differentiation of icysatellites, as it determined the accretion rate and the vol-ume of the heated region after an impact. The increase

1.3 Origin of the Saturnian system and accretion of Titan 43

Figure 1.4 Assemblage sequence and consequences forthermal evolution of the growing planet.

of temperature upon impact in the outer region of thesatellite is controlled mainly by the efficiency of post-impact cooling. The cooling is determined by radia-tion at the surface and subsurface thermal conduction,as well as local heat advection associated with melt-ing, excavation, and impact-induced convection (e.g.,Senshu et al., 2002). If the impact rate is too high orif the impact energy is buried at depths larger than afew kilometers, radiation and conduction are not effi-cient enough to cool the subsurface and the temperatureincreases.

As noted previously, the average size of impactorsis expected to increase as a function of time during theaccretion, and large impactors will be more and morefrequent, possibly leading to giant impacts during thelate stage of accretion (Sekine and Genda, 2012). Thishas major consequences for the efficiency of cooling ofthe satellite and the possible occurrence of melting dur-ing accretion. Heliocentric impactors with high veloc-ities may also become more and more frequent (Agnoret al., 1999), resulting in larger increases in temperature.High-velocity impacts may locally induce significantmelting and vaporization (Artemieva and Lunine, 2003,2005). As the volume of melted material scales withthe impact velocity and the impactor volume (Tonksand Melosh, 1992), giant impact events occurring atthe end of planetary accretion may generate oceans ofmelt (e.g. Sekine et al., 2011). Even for impact withmoderate velocities, their cumulative effect (especiallyfor large impactors) may lead to progressive increaseof temperature up to the melting point of water ice.

1.3.4.2 Energy budget during the early stage ofTitan’s evolution

During the early stages of Titan’s history, in addi-tion to impact heating, three main sources of energymay have contributed to the internal thermal budget:radiogenic heating by short-lived radioactive isotopes(mainly 26Al), tidal heating associated with despinning,and viscous heating due to ice-rock separation. In thissection, we provide some simple estimates of thesesources of energy to evaluate their potential impact onthe thermal budget of early Titan.

Impact heating results from the deposition ofimpactor kinetic energy during satellite accretion. Ittherefore depends on the velocity and the mass of theimpactor as well as the way kinetic energy is convertedinto heat. The impact creates a shock wave that com-presses the satellite beneath the impact site. Below theimpact site, the peak pressure is almost uniform in aquasi-spherical region, called the isobaric core (e.g.,Croft, 1982; Senshu et al., 2002). As shock compres-sion is an irreversible process, the entropy below theimpact site increases, leading to a temperature increase.Most existing shock simulations have been performedat high impact velocity (>10 km.s−1; Pierazzo et al.[1997]; Barr and Canup [2010]), so they cannot beused directly to constrain the shock heating in the low-velocity regime expected during the satellite growth.The heating and associated temperature increase can,however, be estimated from energy balance consider-ations. The temperature increase is determined fromthe peak pressure within the isobaric core, which isdirectly related to the impactor velocity, and from theintrinsic properties of the target materials (density, spe-cific heat, thermal expansion, etc.) (e.g., Senshu et al.,2002).

The impact velocity vimp is a contribution of theasymptotic velocity v∞ ∝

√e2 + i2, with e the eccen-

tricity and i the inclination of the planetary embryosand of the escape velocity vesc =

√8/3πGρR2 of the

growing icy satellite with R its radius, ρ its density, andG the gravitational constant:

v2imp = v2

esc + v2∞. (1.5)

The contribution of asymptotic and escape velocitiesto the impact velocity can be expressed under the form

44 The origin and evolution of Titan

of the gravitational focusing Fg (Kaula, 1979; Schubertet al., 1981; Barr et al., 2010):

Fg = 1 +(

vesc

v∞

)2

. (1.6)

During the early stages of planetary formation, dynam-ical friction reduces eccentricities and inclinations ofgrowing bodies. Hence, early impacts on Titan proba-bly occurred near the two-body escape velocity (i.e., v∞was negligible) (Kokubo and Ida, 1996; Agnor et al.,1999).

In this case and assuming that no significant meltingoccurs when the icy satellite is small, energy balancebetween kinetic energy and impact heating leads toa local temperature increase %T0 that is proportionalto the square of the radius of the growing icy moon(Monteux et al., 2007):

%T0 = 4π

9γ

f (m)ρG R2

C p, (1.7)

where γ is the percentage of kinetic energy retainedas heat deep below the surface, f (m) is the volumeeffectively heated over the volume of the isobaric core(typically equal to 2.5–3; Pierazzo et al. [1997]; Mon-teux et al. [2007]), and C p is the averaged heat capacity(typically varying between 1000 and 1500 J.kg−1.K−1

for ice-rock mixtures). As long as the radius of thegrowing satellite is below 1000 km, the local tempera-ture increase is less than 10 to 15 K after each impact.Above 2000 km, the increase of temperature is largerthan 50 K.

Radiogenic heating by short-lived radioisotopes(26Al, 60Fe, 53Mn) may also significantly contributeduring the accretion period. For simplicity, we considerhere only 26Al, which is the principal contributor (e.g.,Robuchon et al., 2010; Barr et al., 2010). The increaseof temperature due to the decay of 26Al at a radius rdepends on the time, t f , at which the layer at radius rhas been formed and it is given by (e.g., Barr et al.,2010):

%T26(r ) = mr q26(0)C p

τ26

ln 2exp

(− ln 2

t f (r )τ26

), (1.8)

where mr is the rock mass fraction in the accreting mate-rial (around 0.5 in the case of Titan), τ26 is the half-lifetime of 26Al (τ26 = 0.716 Myr), and q26(0) is the initial

26Al heating rate at the time of the calcium-aluminium-rich inclusions (CAIs). t f denotes the time after the for-mation of CAIs. The initial heating rate is determinedby the specific heat production of 26Al H26(0) = 0.341W.kg−1 (Robuchon et al., 2010), the Al content in rocks,[Al], and the initial 26Al/27Al. Assuming a rock com-position similar to CI chondrites, [Al] = 0.865 wt%and initial 26Al/27Al of 5.85 × 10−5 (Thrane et al.,2006; Barr et al., 2010) leads to q26(0) = 1.725 × 10−7

W.kg−1. This is about 20 million times larger than theradiogenic heating rate owing to the decay of long-livedisotopes expected at present within Titan. However, thishuge heat production rapidly decays. For an accretiontime of 2 Myr after CAI formation, the temperatureincrease due to 26Al decay would be of the order of100 K, but it becomes less than 10 K for accretiontimes above 5 Myr. Radiogenic heating by short-livedradioisotopes may have been of importance for the firstmaterials to accrete; that is, for the materials that formedthe inner part of the satellite. By contrast, impact heat-ing becomes significant for layers above a radius ofabout 1000 km.

An additional source of heating corresponds to dissi-pation due to enhanced tidal friction before the satellitereaches synchronous rotation. It is indeed likely thatTitan had an initial spin period much shorter that itspresent-day period of 15.95 days. Initial spin periodsas low as 6 to 10 hours are possible (Lissauer andSafronov, 1991; Robuchon et al., 2010). In this con-dition, the tidal bulge raised by Saturn rotates relativeto the satellite body frame. Owing to the anelastic prop-erties of the interior, the fast-rotating satellite does notrespond perfectly to the tidal forcing, resulting in amisalignment of the tidal bulge with respect to Saturn’sdirection. This acts as a torque that tends to slow downthe satellite spin, until the satellite is locked in spin-orbit resonance. The gradual decrease of the spin rateω as a function of time t is determined from the tidaltorque acting on the satellite and is given by

dω

dt= 3

2k2G M2

S R5

D6ST QC

, (1.9)

with k2 tidal Love number, MS Saturn’s mass, DST meandistance between Titan and Saturn, Q dissipation factorof Titan, and C its polar moment of inertia (C ≃ 0.4 ×MT R2 at the time of accretion). As the despinning rate

1.3 Origin of the Saturnian system and accretion of Titan 45

scales with R5, it increases rapidly when the satellitereaches its final size. For k2 ≃ 0.07, the expected valuefor an homogeneous rigid interior (Sohl et al., 1995) andQ =100–500 and initial rotation period of 10 hours, thetime needed to reach the 16-day period range between2 and 10 Myr, suggesting that despinning is achievedshortly after the accretion. Note that once large-scalemelting occurs, a global liquid layer may form, thusincreasing the Love number and the despinning rate,so that the current period is reached even faster. Thetotal amount of energy dissipated during the despinningstage is given by the difference of body rotation energybetween the initial and final periods and leads to thefollowing temperature increase:

%Tdesp = 12

C(ω20 − ω2

c )C p MT

, (1.10)

with ω0 and ωc the initial and current angular rota-tion rates, respectively, and MT Titan’s mass. For ini-tial periods varying between 10 and 6 hours, the aver-aged temperature increase would range between 25 and75 K, respectively. The localization of the dissipation isvery sensitive to the internal structure (Tobie et al.,2005a). For a homogeneous interior, the dissipationwill be rather uniform, with a maximum of dissipationat about mid-radius. For a differentiated structure, thedissipation will be located mainly in the outer regions(liquid water and ice I) and would be more efficientlytransported to the surface.

Ice-rock separation, which may be triggered alreadyduring the accretion phase, may also contribute sig-nificantly to the energy budget. The total increase ofinternal temperature due to ice-rock differentiation,%Tdiff, can be estimated from the change of gravita-tional energy, %Eg, between a homogeneous interiorand a differentiated interior consisting of a rock coreand a thick H2O mantle (e.g. Friedson and Stevenson,1983):

%Eg ≃ 35

G M2T

RT

[

1 −(

ρr

ρ

)2

X5c +

(ρi

ρ

)2 (1 − X5

c

)

+ 52

ρi

ρ

(ρr

ρ− ρi

ρ

)X3

c

(1 − X2

c

)], (1.11)

with ρr the density of rock, ρi the averaged den-sity of ices and liquid water, Xc the ratio betweenthe rock core radius, Rc and Titan’s radius, RT . For

ρr = 2700 kg.m−3 and ρi = 1250 kg.m−3, Xc = 0.75and %Tdiff = %Eg/cp M ≃ 100–150 K (for cp = 1000–1500 kg.m−3). If the differentiation process is slow (>1Gyr), the convective heat transfer should be able totransport this additional energy. If the ice-rock sepa-ration is fast enough (<0.5 Gyr), the dissipation ofpotential energy may induce runaway melting and thusmay create a catastrophic differentiation (Friedson andStevenson, 1983; Kirk and Stevenson, 1987). In partic-ular, giant impact events occurring at the end of plane-tary accretion (e.g., Sekine et al., 2011) may generate alocally large volume of melts, leading to the formationof rock diapir migrating toward the center (Figure 1.4).The viscous heating associated with the migration oflarge rock diapirs may generate significant melting inthe surrounding ice, potentially triggering differentia-tion runaway in a way similar to what is expected interrestrial planets (e.g., Monteux et al., 2009; Ricardet al., 2009).

1.3.4.3 Possible post-accretional interior structure

Depending on the conditions under which Titanaccreted, different post-accretional interior structuresare possible just after accretion, as illustrated in Fig-ure 1.5. If the accretion is very slow (>1 Myr) and theimpactors remain small all along the accretion period(<100 m), the surface should able to cool down betweensuccessive impacts, and differentiation may be avoided.A fully undifferentiated structure also requires that theaccretion starts about 3 to 4 Myr after the CAI forma-tion, in order to limit the increase of internal temper-ature due to short-lived radiogenic elements, and thatthe initial spin period of Titan was larger than 8 to10 hours. In such a cold accretion case, no significantatmosphere would be generated, and the formation ofa significant atmosphere should occur later during thesatellite evolution (for instance, during the Late HeavyBombardment [LHB] event).

If the growing satellite is hit by larger impactorsat the end of accretion, a surface water ocean may begenerated (R > 2300–2400 km) (Figure 1.5, middle).Even if only 10 to 20 percent of the total volume of icemelts, release of gases (mainly CO2 and CH4; Tobieet al. [2012]) trapped in the ice phase may generatean atmosphere up to 5 to 10 bars. Even if a globalocean is generated, most of the mass of Titan would

46 The origin and evolution of Titan

Figure 1.5 Possible evolution scenarios for the interior of Titan for different initial states. Depending mostly on the efficiency of heattransfer in the interior, different bifurcations in the evolutionary path may have occurred.

1.4 Coupled evolution of the atmosphere, surface, and interior of Titan 47

still be undifferentiated. Finally, if the accretion is fasterand/or if the impactors are large even during the earli-est stage of accretion, ice melting occurs rapidly dur-ing the satellite growth. In this case, a very massiveocean in equilibrium with a very massive atmosphere(P > 50 bars) would be generated. Rocky blocks result-ing from the melting of icy impactors would sedimentat the base of the ocean, forming a thick layer ofunconsolidated rock (Kirk and Stevenson, 1987; Lunineand Stevenson, 1987). The subsequent evolution ofthese different possible initial structure is discussed inSection 1.4.1.

1.3.5 Summary

At the start of this section we posed three questions; wewill conclude by providing some provisional answers.

1. What was the composition of the primordial buildingblocks that built Titan and the Saturnian system? Mostof the solids that formed the Saturnian system probablyoriginate from the surrounding solar nebula. They arepredicted to be composed of rock and ice and to containsignificant amounts of volatiles (NH3, CH4, CO2, H2S)trapped in the form of hydrate and pure condensate. Theapparent deficit in primordial CO, N2, and Ar in Titanrelative to proto-solar composition may be explained bypartial devolatilization of building blocks during theirmigration in Saturn’s subnebula.

2. What make the Jupiter and Saturn systems so different?The difference in satellite architectures may be explainedby different evolutions of the circumplanetary disk. In thecase of Jupiter, the feeding of the disk may have stoppedabruptly, whereas in the case of Saturn, the satellite accre-tion in the disk may have lasted longer owing to reducedmass infall. During the accretion process, several Titan-size moons may have been formed and lost by migrationinto proto-Saturn. Titan would be the only surviving bigmoon of Saturn, whereas Jupiter was able to retain fourbig ones.

3. Can Titan accrete undifferentiated? It seems extremelydifficult to avoid ice melting during the accretion phase,owing to combined effects of impact heating, radio-genic heating by short-lived isotopes, and tidal despin-ning. Melting may be prevented only if Titan accretedslowly (>1 Myr) from a swarm of small building blocks(<100 m). Even for faster accretion with larger impactors,large-scale melting is reached only toward the end of the

accretional growth, so the inner part of the interior struc-ture ( <1500 km) remains undifferentiated. In any case,a totally undifferentiated structure after accretion is notneeded to explain the present-day moderately differenti-ated structure.

1.4 Coupled evolution of the atmosphere, surface,and interior of Titan

During and after the accretion, exchanges among theinterior, surface, and atmosphere are likely to haveoccurred. The icy shell has probably played a crucialrole all along Titan’s evolution by storing and transfer-ring volatile compounds from the internal ocean to thesurface. These exchanges have probably influenced thecomposition of the atmosphere during most of Titan’sevolution, and they may still be operating at present. Inthis section, we discuss the coupled evolution betweenthe different envelopes of Titan, and try to address thethree questions listed below.

1. To what extent did the interior of Titan influence the evo-lution of its atmosphere? Evidence for exchange withthe interior is provided by the detection of a significantamount of 40Ar in the atmosphere. However, the degreeand timing of internal outgassing still remain uncon-strained. In this section, we review the different internalprocesses that may have occurred all along Titan’s evo-lution and discuss their implications for the evolution ofthe atmosphere.

2. Did Titan always possess a massive N2-CH4 atmosphere?The composition and the mass of Titan’s atmosphere haveprobably varied since the satellite accretion. We discussthe different evolution scenarios that can be envisioneddepending on the initial state acquired during the accre-tion and on the efficiency of internal outgassing.

3. Where is the missing hydrocarbon inventory? Hydrocar-bon lakes have been identified at the surface of Titan(Stofan et al., 2007). However, the estimated volume ismuch smaller than what is expected from continuous pho-tolysis of methane since the formation of Titan (Lorenzet al., 2008a). This suggests that either a large fraction ofhydrocarbon products has been removed from the surfaceor that the photochemical production has not been con-stant through time. We discuss the implications of thesetwo different hypotheses for the long-term evolution ofTitan.

48 The origin and evolution of Titan

1.4.1 Differentiation and long-term evolutionof the interior

1.4.1.1 Ice-rock segregation and formationof a rocky core

As already discussed in Section 1.2, the moment ofinertia (C/MT R2

T = 0.33–0.34) inferred from Cassinigravity measurements suggests that Titan’s differentia-tion was not complete. Titan probably does not have aniron core, and a layer consisting of a rock-ice mixturemay still exist between a rocky core and a high-pressureice layer. This possible uncompleted ice-rock separationhas been proposed to be the result of a cold accretionprocess (Barr et al., 2010). However, limited differenti-ation not only requires slow accretion but also impliesthat the convective heat transfer was efficient enoughto remove heat from the interior (Friedson and Steven-son, 1983; Mueller and McKinnon, 1988; Nagel et al.,2004). This is possible if the viscosity of the ice-rockmixtures is sufficiently low. Nagel et al. (2004) showedthat, in the case of Callisto, if the average size of rockswas of the order of meters to tens of meters, the interiormay have experienced a gradual, but still incomplete,unmixing of the two components. Significant lateralheterogeneities among the ice-rock mixtures existingafter accretion could also have enhanced the differen-tiation processes (O’Rourke and Stevenson, 2011). Inthese conditions, convective motions associated withice-rock unmixing processes may be efficient enoughto avoid melting. Following this scenario, the interior ofTitan would be relatively similar to Callisto, but with athinner layer of ice-rock mixed materials, as suggestedby its smaller MoI factor.

Nevertheless, incomplete ice-rock separation is notnecessarily needed, as the MoI factor may also be con-sistent with a highly hydrated silicate interior, possi-bly dominated by antigorite serpentine (Castillo-Rogezand Lunine, 2010; Fortes, 2012). The formation ofa highly hydrated core may result from a rapid dif-ferentiation process, implying large-scale melting andefficient water-rock interactions (e.g., Fortes et al.,2007; Castillo-Rogez and Lunine, 2010; Fortes, 2012).The combined effect of accretional heating, radiogenicdecay by short-lived isotopes, tidal heating associatedwith despinning, and viscous heating associated withnegative rock diapirs may increase the internal temper-ature above the melting point of ice. The presence of

freezing point depressing agents such as ammonia ormethanol (Lunine and Stevenson, 1987; Grasset andSotin, 1996; Grasset and Pargamin, 2005; Deschampset al., 2010) would further favor the generation of melts,as they strongly decrease the melting point of ice. Theaqueous melt, which is less dense than ice-rock mix-tures, at least at pressures above 100 MPa, will per-colate upward, thus interacting with surrounding rockparticles. During their ascent, aqueous solutions inproto-Titan may have been able to interact chemicallywith silicate minerals (Engel and Lunine, 1994) andwith other substances in solution, such as sulfates(Fortes et al., 2007). Interactions of aqueous fluidswith initially anhydrous silicate minerals at low tem-perature may result in the formation of hydrated sil-icate minerals, called serpentine (Fortes et al., 2007;Castillo-Rogez and Lunine, 2010; Fortes, 2012). Thehydrated minerals may remain stable as long as therocky core temperature remains below approximately900 K (Grindrod et al., 2008; Castillo-Rogez andLunine, 2010). Above this temperature, dehydrationwould be expected, releasing warm supercritical waterfluids that may interact again with overlying rocks dur-ing their ascent (Castillo-Rogez and Lunine, 2010).

As pointed out by Iess et al. (2010), the main chal-lenge with this hypothesis is to keep most of the corehydrated (i.e., <900 K until present). Thermal convec-tion is unlikely to start in a cold core dominated byhydrated silicates (Castillo-Rogez and Lunine, 2010).As thermal diffusion is not efficient enough to extractheat, the core temperature should rise up to the dehydra-tion temperature. The temperature increase can, how-ever, be limited if the core is depleted in radiogenicelements. Castillo-Rogez and Lunine (2010) showedthat if at least 30 percent of potassium is leached intothe ocean from the silicate minerals during the earlierhydration event, the internal temperature would slowlyincrease and the outer part of the core would remainhydrated. As shown by Castillo-Rogez and Lunine(2010), a rock core consisting of only antigorite serpen-tine is not required to explain the present-day momentof inertia. An inner anhydrous core with radius up to1300 km overlaid by a hydrated mantle is compati-ble with the MoI factor of 0.33–0.34 (see Figure 1.5,bottom right). The scenario of a fully hydrated coreinitially investigated by Fortes et al. (2007) must beconsidered as a very extreme case, maybe not fully

1.4 Coupled evolution of the atmosphere, surface, and interior of Titan 49

realistic. Nevertheless, the idea that a large fraction ofthe core remains hydrated during most of Titan’s evolu-tion should be considered as a possible scenario. Morecomplete chemical and thermal evolution models areprobably needed in the future to test this hypothesis inmore detail.

As illustrated in Figure 1.5, different evolutionarypaths are possible, some of them combining both solid-state unmixing and ice melting processes. For instance,for an initially warm and gravitationally unstable inte-rior (Figure 1.5, bottom), the rock core may form rapidlyowing to runaway melting of the undifferentiated proto-core, resulting in a catastrophic core overturn (Kirk andStevenson, 1987), or more slowly with progressive ice-rock separation and limited ice melting (e.g., Estradaand Mosqueira, 2011). By contrast, an initially coldundifferentiated interior may slowly differentiate withno melting and no formation of ocean during the entireevolution. Any intermediate evolution is also possible.

1.4.1.2 Water-rock interactions in Titan’s conditions

On Titan, the crystallization of the internal liquid waterresults in the formation of a low-pressure ice layer atits top and a high-pressure ice layer at its base. Onsmaller icy moons, such a high-pressure ice layer doesnot exist due to reduced pressure (high-pressure phasesexist only above about 200 MPa). On Europa and Ence-ladus, for instance, the internal water ocean is perma-nently in contact with the rocky core, thus favoringwater-rock interactions all along the satellite evolution.By contrast, on Titan, a thick layer of high-pressureice rapidly forms during the interior evolution (Figure1.5). The water ocean is expected to be in direct contactwith underlying rocks only during a brief period of time(<100 Myr) just after accretion, and only in the case ofrelatively warm accretion (Figure 1.5).

However, even if there is no direct contact betweenthe water ocean and the rocky core, water-rock inter-actions may have occurred at different periods duringTitan’s evolution, especially during the differentiation.As discussed in Section 1.3.1, non-water ice compo-nents such as CH3OH, NH3, CO2, and H2S are likelyto have been incorporated inside Titan. In this condi-tion, eutectic solutions at temperatures well below 273K would have formed (e.g., Sohl et al., 2010). Thesehighly concentrated eutectic solutions should facilitatedissolution of mineral phases and chemical reactions.

In the case of ammonia-rich fluids, Engel and Lunine(1994) indicated that NH+

4 ions contained in the aque-ous solution may efficiently replace ions such as K+,Na+, Rb+, and Ca2+ contained in the mineral phase.Ammonia-rich liquids are also expected to react withmagnesium sulfate, contained in the chondritic phase,and with carbon dioxide brought in the ice phase to formammonium sulfate and ammonium carbonate, respec-tively (Kargel, 1992; Fortes et al., 2007).

During the migration of the enriched fluids towardthe surface, their composition should vary as new ele-ments are extracted from primary minerals and otherprecipitate in secondary phases (e.g., Zolotov, 2010). Inturn, the circulation of aqueous solutions leads to oxy-dation and hydration reactions, analogous to aqueousprocesses on chondrite parent bodies (e.g., Sohl et al.,2010). These reactions are well documented at low-pressure conditions and have been used, for instance, toconstrain the possible composition of Europa and Ence-ladus’ ocean (e.g., Zolotov, 2007; Zolotov and Kargel,2009), but they are still poorly known at the high pres-sure range of Titan’s proto-core (>1 GPa). Only onestudy by Scott et al. (2002) tried to investigate the aque-ous alteration of chondritic materials at high pressure.However, Scott et al. (2002) did not take into accountthe probable presence of ammonia, methanol, or othercompounds, which may significantly affect the chem-ical equilibrium. Future experimental and thermody-namical investigations will be needed to better under-stand the complex chemistry that probably occurredduring the differentiation.

After the complete separation of rock and ice, water-rock interactions are limited to the ice-rock interface atthe base of the high-pressure layer. Even though wateris mostly in the form of high-pressure ice phase at thisinterface owing to the elevated pressure (>0.8–1 GPa),liquid water may exist locally and potentially circulatein the upper part of the rock core. To efficiently trans-fer the heat coming from the rocky core through thehigh-pressure ice layer, the temperature profile in thehigh-pressure ice is probably close to the melting curveof water ice (e.g., Sotin and Tobie, 2004). In this con-text, the formation of a transient liquid reservoir at therock-ice interface is likely. These liquids may poten-tially interact with the underlying rocks before perco-lating upward to the ocean. Warm liquid water mayalso be released from the core when the dehydration

50 The origin and evolution of Titan

temperature is reached (Castillo-Rogez and Lunine,2010). This hydration event may have occurred rela-tively late during Titan’s evolution (>3–3.5 Gyr afterformation), provided sufficient amounts of radiogenicpotassium were leached from the core and that thermaldiffusivity in the core was not too low (Castillo-Rogezand Lunine, 2010) (see Section 1.4.1.1). In summary, inwarm evolution scenarios, intensive water-rock interac-tions were likely during two main sequences of Titan’sevolution: before 1 Gyr during the differentiation andafter 3 to 3.5 Gyr during core dehydration. Moreover,reduced water-rock interactions may also exist all alongTitan’s evolution owing to circulation of high-pressureice melts at the rock-ice interface.

It has been proposed that water-rock interactions maylead to efficient production of methane from carbondioxide, carbon grains, or organic matters by serpen-tinization (Atreya et al., 2006). However, this modeof production seems incompatible with the deuterium/hydrogen (D/H) ratio observed at present in the atmo-spheric methane, which is about three times greaterthan the value expected in hydrothermally producedmethane (Mousis et al., 2009a). It has also been sug-gested that the oxidation of primordial NH3 to N2 mayalso have occurred in hydrothermal systems on earlyTitan (Glein et al., 2009), as well as on early Ence-ladus (Matson et al., 2007; Glein et al., 2008). Theendogenic production of N2 from NH3 is possible onlyif the hydrothermal systems were sufficiently oxidizedwith C existing predominately in the form of CO2, notCH4. It also implies that primordial NH3 was enrichedin 15N, to explain the elevated 15N/14N observed inatmospheric N2. Although water-rock interactions andassociated chemical reactions are thermally plausible,there are essentially no constraints available to assesstheir likelihood. For instance, constraints on the possi-ble primordial value of 15N/14N in NH3 or D/H in waterare still missing. Future measurements in comets byRosetta, for instance, will bring new bearings on theseideas.

Circumstantial evidence for water-rock interactionsis, however, provided by the detection of 40Ar in theatmosphere (Niemann et al., 2005; Waite et al., 2005;Niemann et al., 2010). 40Ar is one of the decay prod-ucts of 40K initially contained in silicate minerals (e.g.,Engel and Lunine, 1994). The mass of 40Ar in the atmo-sphere is estimated at between 4.0 and 4.6 × 1014 kg

(Niemann et al., 2010), which corresponds to about 7.5to 9 percent of the total mass that may be contained inthe rocky core, assuming a CI chondritic composition(Tobie et al., 2012). As a comparison, on Earth, Mars,and Venus, the fraction of 40Ar outgassed from the sil-icate mantle is estimated to about 50 percent on Earth(Allegre et al., 1996) and about 25 percent Mars andVenus (Hutchins and Jakosky, 1996; Kaula, 1999). Theoutgassing process is probably very different on Titan.On Earth and terrestrial planets, 40Ar is progressivelyoutgassed from the silicate mantle through volcanism(e.g., Xie and Tackley, 2004). 40K decays into 40Ar inthe silicate minerals and the latter is released when theminerals are brought near the surface by convectivemotions and melt. On Titan, as suggested by the grav-ity data, Titan’s rocky core is probably relatively coldduring most of its evolution (<1000 K; Castillo-Rogezand Lunine [2010]); therefore, it has probably neverreached magmatic temperatures (at least on a globalscale). Moreover, even if silicate melting has occurred,outgassing of erupting silicate magmas should be lim-ited due to the high-pressure conditions expected atthe surface of the rocky core (1–1.5 GPa) (McKinnon,2010).

On Titan, the extraction of 40Ar is more likelythe result of a three-step mechanism. 40K may havebeen extracted first from the silicate minerals by waterinteractions (e.g., Engel and Lunine, 1994). Then dis-solved 40K+ ions in the ocean would have decayed into40Ar. Finally, the 40Ar contained in the ocean wouldbe extracted through direct outgassing of the oceanor progressive incorporation in the outer shell in theform of clathrate hydrate and then destabilization nearthe surface by cryovolcanism (e.g., Tobie et al., 2006)or impact (e.g., Barr et al., 2010). Assuming that 30percent of potassium has been leached into the oceanimplies that about 25 to 30 percent of the produced 40Arhas been outgassed through the outer icy shell, whichseems a reasonable outgassing efficiency.

The presence of significant amounts of 22Ne in theatmosphere, if confirmed, may also be indicative ofwater-rock interactions and efficient outgassing pro-cesses (Tobie et al., 2012). Indeed, a chondritic core maycontain 22Ne in sufficient amounts to explain the tenta-tively observed abundance in the atmosphere (Niemannet al., 2010). As further discussed in Section 1.4.1.4,neon is the least stable gas compound in the form

1.4 Coupled evolution of the atmosphere, surface, and interior of Titan 51

of clathrate and the least soluble in liquid water, andtherefore it is expected to be very efficiently extractedduring methane-driven outgassing events. The accuratedetermination of the abundance of the different isotopesof Ne (20Ne, 21Ne, and 22Ne) by a future mission willpermit confirmation of the tentative detection of neonby Huygens, and it will provide pertinent tests on theorigin of neon and on its extraction mechanism fromthe interior.

1.4.1.3 Volatile storage in the H2O mantle

A major issue on Titan concerns the storage and trans-port of gas compounds in the H2O mantle, which con-stitute about 60 to 70 percent of the satellite volume.The amount of gas compounds that can be stored in theocean and in the ice layers is determined by their solu-bility in water and their stability relative to the clathratephase. NH3 and CH3OH, for instance, are highly solu-ble in liquid water and do not form clathrate structures.At low temperature (T < 177 K), ammonia combineswith water molecules to form ammonia hydrate (notclathrate) and methanol solidifies at relatively similartemperatures. Such low temperature conditions can bereached only near the cold surface, and both ammo-nia hydrates and methanol ice become unstable belowa depth of 5 to 10 km (depending on the subsurfacethermal gradient). As a consequence, ammonia andmethanol should be nearly exclusively stored dissolvedin the internal ocean.

Although the solubility of most of the other gases –for instance, methane – in liquid water, is low, largequantities of gas can be stored dissolved in the liquidwater, as the ocean is very massive on Titan and it isconfined at relatively elevated pressures (>25 MPa).For example, more than 2000 times the present-daymass of atmospheric methane may be stored dissolvedin the ocean (Tobie et al., 2012).

Depending on their stability relative to the clathratephase, a fraction of the gases dissolved in liquid watercan, however, be incorporated in the form of clathratehydrate (Choukroun et al., 2010; Tobie et al., 2012;Choukroun et al., 2013). The most stable speciesexpected in the ocean, such as H2S and Xe, shouldbe preferentially incorporated in the form of clathratehydrate, whereas the least stable species, such as COand Ne, remain mainly in the liquid phase (Tobie et al.,2012). The fraction incorporated in the clathrate phase

should also increase as the ocean cools down and crys-tallizes. Clathrate hydrates that have a compositiondominated by H2S have a density larger that of liq-uid water, so they are expected to sink to the base of theocean. As an example, clathrate structures containing75 percent of H2S, 20 percent of CO2, and 5 percent ofCH4 have a density of about 1050 kg.m−3. These H2S-rich sinking clathrates will also preferentially incor-porate Xe owing to its very high stability in clathratephase. As a consequence, H2S and Xe are expected tobe stored mostly in the form of clathrate hydrate withinthe high-pressure ice layer.

During the clathration process, the dissolved fractionof H2S progressively reduces, and CH4 and CO2 even-tually become the dominant clathrate former. WhenCH4 becomes the dominant clathrate former gas, theclathrate density gets lower than the density of liquidwater, and clathrate hydrates accumulate at the top ofthe ocean. Differentiation of Titan’s interior during theearly stage of its evolution may also have led to massiverelease of methane-rich clathrate hydrates, and henceaccumulation in the outer shell above the ocean (Tobieet al., 2006). Both early differentiation and late oceancrystallization may have resulted in the accumulationof methane-rich clathrate hydrates in the outer shell.However, even in these conditions, the total thicknessof methane-rich clathrates incorporated in the outer iceshell should not exceed 5 to 10 km, corresponding toabout 20 percent of the total mass of methane that ispotentially contained in the interior (Figure 1.6; Tobieet al., 2012). A fraction of CO2, CO, and Ar shouldalso be incorporated in the outer shell in the form ofclathrate, but to a lesser extent than methane owing todifferences in solubility and clathrate stability.

1.4.1.4 Volatile release from the H2O mantle

Even though the interior may potentially contain largequantities of volatiles, their release into the atmosphereis limited owing to their solubility in the internal oceanand owing to clathration processes within the H2Omantle. Outgassing from the interior may occurthrough two different mechanisms: (1) destabilizationof clathrates stored in the outer ice shell by thermalconvective plumes or by impacts, and (2) mobilizationof dissolved gas molecules contained in the water oceanand subsequent bubble formation if the ocean gets sat-urated (Tobie et al., 2006, 2009). Both processes are

52 The origin and evolution of Titan

Figure 1.6 Possible present-day structure of Titan’s interior and estimated volatile distribution within the thick H2O mantle. Theradius, temperature, and pressure of each interface are only indicative.

controlled by methane, as it is the dominant gas in theform of clathrate hydrate in the outer shell and becauseit is the only gas compound that may reach the sat-uration point in the ocean during periods of reducedice shell thicknesses. The other dominant compounds,CO2, NH3, or H2S, are more soluble in liquid water, sothey will never reach saturation. For noble gases, suchas Ne or Ar, although they have a low solubility, theirconcentration is very low and always remains belowthe saturation point. However, when methane reachesthe saturation point, Ne and Ar, owing to their low sol-ubility in water, may be preferentially incorporated inCH4-rich bubbles formed at the ocean-ice interface andtransported upward with the percolating fluids. Follow-ing this extraction mechanism, the detection of 40Ar(and potentially 22Ne) suggests that the atmosphericmass of methane has been renewed at least 20 to 30times during Titan’s evolution (Tobie et al., 2012).

1.4.2 Formation and evolution of the atmosphere

1.4.2.1 Origin of the massive N2-CH4 atmosphere

The formation of Titan’s atmosphere is certainly anevent that occurred relatively early in its history;

otherwise, we should see a number of impact craters toosmall to have been formed under the dense atmosphereof today. However, whether it is an accretionally gen-erated atmosphere (e.g., Lunine and Stevenson, 1987;Kuramoto and Matsui, 1994) or a late veneer gener-ated during the subsequent Late Heavy Bombardment(LHB) (Zahnle et al., 1992; Sekine et al., 2011) is anissue that cuts to the heart of the question of whetherTitan’s formation environment was chemically and ther-mally distinct from that of Jupiter (discussed earlier). Ifthe disk around Saturn was colder than that at Jupiter,rich in volatiles such as methane and ammonia, thenthe atmosphere most likely was primordial. The likelypresence of an interior water ocean, outgassing throughtime of methane needed to replace that destroyed byphotolysis, and the elevated nitrogen isotopic ratio inTitan’s atmosphere (Niemann et al., 2010) all hint atthe likelihood that Titan acquired its volatile inventorythroughout accretion, the present atmosphere being anoutgassed remnant of a volatile-rich Titan.

A volatile-rich Titan formed under all but the cold-est accretional scenarios would have had an initiallywarm ammonia-rich atmosphere, if indeed the domi-nant carrier of nitrogen was ammonia. This seems to

1.4 Coupled evolution of the atmosphere, surface, and interior of Titan 53

be required by the measurement of 36Ar to nitrogen(Niemann et al., 2010), which can be interpreted asrequiring that molecular nitrogen was initially in theform of ammonia (Owen, 1982). Arguments to the con-trary – that somehow the argon is being sequestered bya condensed or crustal reservoir – do not seem to makesense in light of the presence of radiogenic argon ata two orders of magnitude larger mixing ratio. Thuswe can assume that ammonia was in abundance onearly Titan, initially in the form of an atmosphere andsurface ammonia-water ocean supported by accretionalheating and a strong greenhouse effect (Lunine et al.,2009). Conversion of ammonia to molecular nitrogenby photolysis (Atreya et al., 1978) and shock heating(Jones and Lewis, 1987; McKay et al., 1988) all pro-vide a plausible story in which the ammonia is rapidlyconsumed, the atmosphere cools and condenses out, andnitrogen takes its place. Ishimaru et al. (2011), however,demonstrated that efficient conversion of NH3 into N2

by impact shock requires an oxidizing CO2-rich proto-atmosphere on Titan. This is consistent with CO2 asthe dominant gas compound in the building blocks thatformed Titan, as predicted by Hersant et al. (2008);Mousis et al. (2009c); and Tobie et al. (2012).

Sekine et al. (2011) also proposed, based on laser-gun experiments, that ammonia ice contained in theouter icy shell may have been efficiently converted toN2 by high-velocity impacts during the Late HeavyBombardment. The generation of the N2-rich atmo-sphere by this process requires the presence of sev-eral percent of ammonia in the outer shell, whichimplies that Titan’s outer layer remained undifferen-tiated and unmelted until the LHB era. As alreadydiscussed, it is very difficult to avoid melting andpartial differentiation during the accretion, so it isunlikely that Titan’s icy shell contained ammonia insufficient amounts at the time of the LHB. Althoughthe LHB event has probably affected the evolution ofthe atmosphere, notably by inducing a significant atmo-spheric erosion (Korycansky and Zahnle, 2011; Ishi-maru et al., 2011), it seems unlikely that it has been themain contributor for the atmosphere generation; a mas-sive atmosphere probably already existed prior to thatevent.

Outgassing of methane during the accretion andpossibly during the LHB would have been vigorous.However, the present-day methane in the atmosphere

is unlikely to be the remnant of this primordial out-gassing. If that was the case, the C13/C12 in CH4

should be much larger than the observed value, whichis close to the solar value (Mandt et al., 2009; Niemannet al., 2010). This nearly unmodified ratio suggests thatmethane released during the early stage of evolutionshould have been totally lost or rapidly retrapped inthe subsurface without being affected by atmosphericfractionation processes. Today’s methane was proba-bly injected much later during the evolution, possiblydue to the destabilization of crustal reservoir of CH4-rich clathrate (Tobie et al., 2006). The only record ofthe accretional time lies in the nitrogen isotopic ratio,modestly heavier than the terrestrial value (Mandt et al.,2009), suggesting a source distinct from the Earth, orloss of nitrogen to space, or both. However, a singleisotope ratio is a poor constraint on the rate and volumeof the loss because isotopic enrichment is a functionof the rate and mechanism of escape. Further extensivediscussion of the early Titan atmosphere may be foundelsewhere (Lunine et al., 2009).

1.4.2.2 Interactions with the icy shell

All along Titan’s evolution, interactions with the solidicy shell have likely influenced the composition of theatmosphere. During the post-accretional cooling lead-ing to the crystallization of the primordial water ocean,part of the atmospheric gases, notably CO2 and CH4,may have been retrapped first by clathration and thenby condensation (Choukroun et al., 2010; Lunine et al.,2009). Large quantities of CO2 and CH4 were proba-bly incorporated in the primitive icy shell. AdditionalCH4-rich clathrates released during the differentiationof the deep interior may have also accumulated at thebase of this primitive outer shell, leading to the forma-tion of a thick (>10 km) CH4-rich clathrate layer (Tobieet al., 2006). This shell would then thin in response tothe release of heat associated with the core evolution(Tobie et al., 2006), resulting in the release of largeamounts of methane that would be mainly stored dis-solved in the ocean; only the gas in excess may havebeen outgassed (Tobie et al., 2009).

According to this scenario, several kilometers ofCH4-rich clathrates may still be present in the outer icyshell at present, and recent thermal destabilizations ofthis reservoir may explain the present-day atmosphericmethane (Tobie et al., 2006). CH4-rich clathrates are

54 The origin and evolution of Titan

probably not the only crustal reservoir. Large volumesof liquid hydrocarbons may also potentially be storedin a porous crust (Kossacki and Lorenz, 1996; Vosset al., 2007). The morphology and distribution of thehydrocarbon lakes on Titan suggest that they are con-nected to a subsurface reservoir (Hayes et al., 2008)(see also Chapter 2). However, no constraints exist onthe extension of this subsurface reservoir as well as onthe methane content. The lakes seems dominated byethane (Mitri et al., 2007; Brown et al., 2008; Cordieret al., 2009), which suggests that the methane fractionin a subsurface liquid reservoir is limited at present.Based on the available observations, it is impossible toconstrain the contribution of this potential reservoir tothe methane inventory. Additional observations, suchas subsurface sounding by a future exploration mis-sion, will be needed to better constrain the nature of thesubsurface reservoir.

The mysterious absence or sparsity of the primor-dial noble gases in Titan’s atmosphere may also beindicative of interactions of the atmosphere with thesolid icy shell. It has been shown recently that theapparent deficit in noble gases could be essentiallydue to the formation of multiple guest clathrates at thesurface-atmosphere interface of the satellite (Mousiset al., 2011). These clathrates would preferentially storesome species, including all heavy noble gases, overothers. These results have been obtained from the useof a statistical-thermodynamic model based on recentdeterminations of intermolecular potentials and super-sede the previous ones that suggested a limited trap-ping of argon (Osegovic and Max, 2005; Thomas et al.,2008). In fact, the clean water ice needed for the for-mation of these clathrates may be delivered by suc-cessive episodes of cryovolcanic deposits. In this sce-nario, the formation of clathrates in the porous lavasand their propensity for trapping Ar, Kr, and Xe wouldprogressively remove these species from the atmo-sphere of Titan over its history. A global clathrate shellwith an average thickness not exceeding a few metersmay be sufficient on Titan for a complete removalof the heavy noble gases from the atmosphere. How-ever, there has been no convincing and unambigu-ous evidence for cryovolcanic activities on Titan sofar (see, for instance, Moore and Pappalardo, 2011),so the proposed trapping mechanism still remainshypothetical.

If clathrate formation is an efficient process, ashypothesized for the noble gases, then it is possible thatclathrates are also a sink for the prodigious amounts ofethane, and lesser amounts of propane, expected to havebeen produced over the history of Titan by photochem-ical processing of methane. Episodes of crustal vol-canism could encourage formation of ethane clathrate(Mousis and Schmitt, 2008), as could impacts (Lunineet al., 2010), or simple percolation through a heavilyfractured crust might progressively allow the substitu-tion of ethane for methane in crustal clathrate hydrateand reconversion of water ice that once was clathrateprior to losing methane by outgassing. Regardless of theprocess, whereas methane clathrate hydrate is slightlyless dense than the underlying water-ammonia ocean,ethane clathrate is not. Hence any ethane clathrate willhave a tendency to move downward in a predominantlymethane clathrate shell, and once at the bottom will sinkin the ocean to exist stably at the bottom. Timescalesmight be an issue here, leaving most of the ethane inthe outer shell but sequestered away from the surface-atmosphere system. If this view is correct, Titan’s outershell is laden with ethane clathrate from billions ofyears of methane photolysis, the lakes and seas beingonly the most recently produced ethane.

1.4.2.3 Long-term climate stability ofTitan’s atmosphere

Titan’s atmosphere has been unstable on timescales ofbetween tens and hundreds of millions of years unlesslarge areas of methane seas and oceans, an order ofmagnitude or more larger than the present-day lakes andseas, have existed on the surface throughout geologictime. In the absence of such sustaining bodies of liq-uid methane, or an implausibly finely tuned outgassingrate that keeps the atmospheric methane content withina factor of several of the current abundance for bil-lions of years, the methane in Titan’s atmosphere musthave been periodically depleted. For the first half ofTitan’s history, when the Sun’s luminosity was less than80 to 85 percent of its current value, loss of methanemay have caused the nitrogen to rain out onto Titan’ssurface (McKay et al., 1993; Lorenz et al., 1997a), pro-ducing a global sea of liquid nitrogen under a thinnerclear atmosphere.

This scenario, bizarre in its contrast with the present-day state, raises all sorts of issues: (1) Is there geologic

1.5 Conclusion and questions for future missions 55

evidence for a long-standing ancient global sea of eitherliquid nitrogen or (to avoid this catastrophe) a standingancient global sea of methane and ethane? (2) Howcould a thicker atmosphere be restored, even with mas-sive outgassing of methane (why wouldn’t the methanesimply freeze out at the lower global temperatures forwhich liquid nitrogen is stable)? (3) What would pre-vent, under a thin atmosphere, the formation of massiveice caps of nitrogen at cold poles, possibly rendering,as in point (2), a permanently frozen Titan? It is pos-sible that freeze-out did not occur even under a fainteryoung Sun, and indeed for values of the solar luminos-ity within 15 to 20 percent of the present-day values,it seems that a nitrogen atmosphere not much differentfrom that at present could have been sustained in theabsence of methane (Lorenz et al., 1997a).

As crustal and deeper sources of methane continueto be depleted by photolysis, it seems inevitable thata state will be reached in which Titan will lose itssurface methane and revert to a nearly pure nitrogenatmosphere, largely clear of aerosols, but with a sur-face temperature not much different from today or evenlarger as the Sun continues to grow in luminosity. Even-tually, when the Sun becomes a red giant, Titan willhave lost much of its atmosphere thanks to a stronglyenhanced solar wind, but will for a time support a sur-face water ocean until the red giant’s luminosity beginsto fade (Lorenz et al., 1997b). Perhaps any life extant inthe deep water interior might then adapt to the surface,develop a form of photosynthesis, and so on, if time isnot too brief. Such speculations would be in the realmof science fiction were it not for the likelihood that Titanis representative of a class of planets that would stablyexist within 1 AU of M to K dwarfs, with a spectrum ofenvironments from the present-day state to an ancientnitrogen sea state, to the ultimate warmer version dis-cussed here, simply by virtue of adjusting the stellarspectral type from late M to K (Lunine, 2010).

1.4.3 Summary

At the start of this section we posed three questions; wewill conclude by providing some provisional answers.

1. To what extent did the interior of Titan influence the evo-lution of its atmosphere? The probable existence of an

internal ocean at present suggests that the differentia-tion of Titan’s interior has been accompanied by signifi-cant water-rock interactions and massive releases of gascompounds initially incorporated in clathrate structure.Although most of the gas species should be stored dis-solved in the massive internal ocean and in the form ofclathrate in both the outer ice layer and the high-pressureice layer, a significant fraction of CH4 was likely out-gassed through time.

2. Did Titan always possess a massive N2-CH4 atmosphere?Titan likely acquired its atmosphere during its first hun-dred millions of years, at the latest during the LHB. Theatmospheric composition and mass have probably var-ied along Titan’s evolution. The early atmosphere mayhave been more massive and dominated by CO2. Follow-ing the post-accretional cooling phase, nitrogen proba-bly rapidly became the dominant gas. It is still unclearhow nitrogen has been produced from ammonia and howatmospheric escape has influenced the nitrogen inventory.The methane abundance has probably varied significantlythrough time, depending on the balance between replen-ishing and destruction. Long periods without atmosphericmethane may even be envisioned.

3. Where is the missing hydrocarbon inventory? Large vol-umes of hydrocarbon may be stored in the form of liq-uid or clathrate hydrate in the solid icy shell. However,there are still no direct constraints on the volume of theseputative reservoirs. Alternatively, the lack of extensivehydrocarbon reservoirs may indicate that the productionof hydrocarbons from methane photolysis has been lessefficient than previously anticipated. This may suggestthat methane has been present in the atmosphere duringlimited periods of time.

1.5 Conclusion and questions forfuture missions

Since the arrival of the Cassini-Huygens mission tothe Saturn system, a large amount of data has beencollected on Titan and various models have been devel-oped to interpret them in an evolution context. In spiteof considerable progress in our understanding of Titan’sorigin and evolution, many questions remain, and theywill probably not be answered by the end of the Cassini-Huygens mission. A series of investigations in situ andfrom the orbit will be needed in the future to addressall the mysteries of Titan that Cassini-Huygens has juststarted to unveil. Specifically, some of the remainingproblems include the following:

56 The origin and evolution of Titan

1. What was the composition of the building blocks thatformed Titan? Theorotical models of planetary formationsuggest that the building blocks that formed Titan con-tained a few percent of various gas compounds, mainlyCO2, NH3, and CH4, as well as sulfur compounds,methanol, and traces of noble gases. Most of the gascompounds initially expected have not been detected sofar. Xe and Kr, for instance, appear to be depleted relativeto the predicted abundance. By contrast, Ne, which wasnot expected, has been tentatively detected. CO2, which ispredicted to be the dominant carbon-bearing species, hasbeen identified only at Huygens’ landing site (Niemannet al., 2010) and tentatively in some other specific areas(McCord et al., 2008). Constraining the CO2 fraction onTitan as well as the noble gas content will provide keyconstraints on the conditions under which the buildingblocks formed. Atmospheric sampling by a future explo-ration mission will permit us to determine more accuratelythe noble gas content of Titan’s atmosphere, in particularto confirm the detection of neon. Determination of noblegas abundance in Saturn’s atmosphere by an entry probewill also be extremely useful to determine the efficiencyof noble gas trapping in planetesimals embedded in thefeeding zone of Saturn. Moreover, in situ analysis in dif-ferent materials collected at the surface will also providepertinent information on the possible CO2 content in theouter shell as well as on the possible sequestration ofnoble gases in the subsurface.

2. How did Titan’s N2-dominating atmosphere form andevolve? It is likely that Titan’s atmosphere was acquiredrelatively early during the satellite evolution. However, itis still unclear how the N2-dominating atmosphere wasproduced and how its composition has evolved throughtime. The elevated 15N/14N ratio observed in Titan’s nitro-gen is commonly attributed to enhancement due to mas-sive atmospheric escape. However, we are still missingconstraints on the initial value in Titan. Determina-tion of the 15N/14N ratio in NH3 in comets, in partic-ular by the Rosetta mission on Comet 67P/Churyumov-Gerasimenko, will provide the first direct constraint. Suchmeasurements, together with determination of isotopicratios in nonradiogenic isotopes of argon (38Ar/36Ar) andneon (22Ne/20Ne) by in situ sampling, will allow us todetermine how much mass Titan’s atmosphere have lostsince its formation, thus providing constraints on the earlystate of Titan’s atmosphere.

3. How efficient was the interaction between rock and H2O inthe interior? The moment of inertia inferred from Cassinigravity data, as well as the detection of 40Ar in the atmo-sphere, is indicative of prolonged interactions between

silicate minerals and H2O (liquid water or water ice).Such an interaction has key implications for the chemicalevolution of Titan’s interior, in particular for the astrobi-ological potential of its internal ocean. Its efficiency maybe constrained by future exploration missions (1) by con-straining the degree of differentiation and ocean charac-teristics (size and composition) from tidal and magneticmonitoring, as well as from seismic detection using afixed lander network at the surface; (2) by detecting saltsand ammonia hydrates at the surface from spectroscopyand surface samplings; and (3) by measuring the ratiobetween radiogenic and nonradiogenic isotopes in noblegases other than Ar from in situ mass spectrometry in theatmosphere and/or in hydrocarbon lakes.

4. Are Titan’s ice shell and interior currently active? In spiteof the Earth-like landscape of Titan, there is still somedoubt regarding the internal activity of Titan. Althoughthere is some evidence of interactions with the interior inthe past, it is uncertain whether Titan’s interior still has aninfluence on the evolution of the atmosphere and surface.Precise determination of the topography and gravity fieldby a future mission may provide key information on thestructure of the icy shell, and thus on its internal dynam-ics. Identification of recent (unambiguous) cryovolcanicflows using high-resolution images and IR spectra wouldhint at the likelihood of recent activities. The detection ofhelium, which has a very short residence time in the atmo-sphere, would be indicative of contemporaneous internaloutgassing. Finally, direct observations of eruption wouldbe the most convincing evidence.

5. How long has the methane cycle been active on Titan?At present, Titan has a very active methane cycle, char-acterized by photochemical production of more complexhydrocarbon species, formation of a haze layer, cloudformation and rainfall, and accumulation of hydrocarbonliquids in near-polar lakes and seas. However, it is stillunclear whether this cycle has existed during most ofTitan’s evolution or whether it was triggered only rela-tively recently. A more comprehensive discussion of thephysical processes that control the methane cycle on Titanis presented in Chapter 6. Determining the possible vol-ume of hydrocarbon stored under the surface by subsur-face sounding with ground-penetrating radar by a futuremission will provide pertinent constraints on the dura-tion of CH4-driven chemistry on Titan. Identification ofclathrates of methane, ethane, and others at the surface byin situ sampling at the surface will also permit us to eval-uate the role of clathration in the storage of hydrocarbonin the subsurface. Determination of 13C/12C in variouscarbon-bearing materials collected at the surface will also

References 57

provide key information on the atmospheric fractionationprocess, and potentially on the duration of active methanecycle. High-resolution mapping of the surface may alsoprovide information on different climatic conditions inthe past, such as evidence for more extended methaneseas associated with wetter climate or for N2 ice capsassociated with a dryer and colder climate.

Acknowledgments

Part of the research leading to these results receiv-ed funding from the European Research Councilunder the European Community’s Seventh Frame-work Programme (FP7/2007-2013 Grant Agreementno. 259285). J. M. is funded by the Agence Nationalede Recherche (project “Accretis”, no. ANR-10-PDOC-001-01). Support from the Cassini Project, NASA-PGG, and NASA CDAP in the preparation of this chap-ter is also gratefully acknowledged.

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