Abstract et al., 2009) suggest that the primordial building blocks that comprise the Saturnian system were probably more Saturn’s moon Titan is the second largest natural satellite volatile-rich than Jupiter’s. in the solar system, and the only one that possesses a sub- stantial atmosphere. The origin of the massive nitrogen at- The composition of the present-day atmosphere, domi- mosphere and the source of the present-day methane, which nated by nitrogen, with a few percent methane and lesser is continuously destroyed by UV-driven photochemistry, has amounts of other species, probably does not directly reﬂect long been debated. Data collected by the Cassini-Huygens the composition of the primordial building blocks and is mission since its arrival to Saturn in 2004 now provide key rather the result of complex evolutionary processes involv- constraints on Titan’s atmosphere composition, surface mor- ing internal chemistry and outgassing, impact cratering, phology and interior structure, which restrict the possible photochemistry, escape, crustal storage and recycling, and 36 scenarios of formation and evolution. In the present chap- other processes. The low Ar/N2 ratio measured by the ter, after reviewing our present knowledge about the interior GCMS on Huygens (Niemann et al., 2005, 2010) suggest structure and composition based on the Cassini-Huygens that nitrogen was not brought to Titan in the form of N2, data, we present the diﬀerent physical and chemical pro- but rather in the form of NH3 (Owen, 1982). The argument cesses that have potentially aﬀected the origin and evolution goes as follows: Ar and N2 have similar volatility and aﬃn- of Titan. In particular, we argue that all along Titan’s evolu- ity with water ice. Thus, if the primary carrier of Titan’s tion, from accretion to present, interactions between the icy N2 were molecular nitrogen itself, the Ar/N2 ratio on Titan shell, the internal water ocean and the deep rock-dominated should be within an order of magnitude of the solar com- interior may have aﬀected the evolution of the atmosphere. position ratio of about 0.1 (Lunine and Stevenson, 1985), In spite of considerable progresses in our understanding of which is about 500,000 times larger than the observed ratio. Titan thanks to Cassini-Huygens,manyquestionsstillre- This indicates that nitrogen has been brought in the form of main, and we conclude on how they may be addressed by easily condensable compounds, most likely ammonia, which future exploration missions. then has been converted in situ by impact-driven chemistry (McKay et al., 1988; Sekine et al., 2011; Ishimaru et al., 2011) or by photochemistry (Atreya et al., 1978). As will be further discussed in the chapter, both conversion processes 2.1 INTRODUCTION required restrictive conditions to occur, and therefore spe- ciﬁc evolutionary scenarios for Titan. Although Titan is similar in terms of mass and size to Jupiter’s moons, Ganymede and Callisto, it is the only one Photochemical processes are also known to lead to an harboring a massive atmosphere. Moreover, unlike the Jo- irreversible destruction of methane, implying that a source vian system populated with four large moons, Titan is the of replenishment must exist to explain its few percent only large moon around Saturn. The other Saturnian moons abundance (vertically averaged) in the atmosphere. In the are much smaller and have an average density at least 25% absence of such a replenishment, the atmospheric methane less that Titan’s uncompressed density and much below should disappear over a timescale of several tens of millions the density expected for a Solar composition (Johnson and of years. Before the Cassini-Huygens mission, the atmo- Lunine, 2005), although with a large variation from satellite spheric methane was proposed to be in equilibrium with to satellite. Both Jupiter’s and Saturn’s moon systems are aglobalsurfaceliquidreservoir(e.g.Lunineetal.,1983). thought to have formed in a disk around the growing giant Even though liquid reservoirs have been identiﬁed on Ti- planet. However, the diﬀerence in architecture between the tan’s surface by Cassini in the form of lakes and seas at two systems probably reﬂects diﬀerent disk characteristics high latitudes (Stofan et al., 2007), their estimated total and evolution (e.g. Sasaki et al., 2010), and in the case volume is too small to sustain methane in the atmosphere of Saturn, possibly the catastrophic loss of one or more on geologic timescales (Lorenz et al., 2008a). This suggests Titan-sized moons (Canup, 2010). Moreover, the presence that a subsurface reservoir of methane exists and replen- of a massive atmosphere on Titan as well as the emission ishes the atmospheric reservoir episodically or continuously. of gases from Enceladus’ active south polar region (Waite Diﬀerent sources have been proposed, but despite a number
12 The Origin and Evolution of Titan 13 of new constraints from the Cassini-Huygens mission, we 2.2 PRESENT-DAY INTERIOR lack a strong discriminant among them. The 12C/ 13Cratio STRUCTURE AND DYNAMICS measured in methane by the GCMS on Huygens (Niemann 12 13 et al., 2010) is similar to the C/ Cratiomeasuredin Models of Titan’s formation (Section 2.3) and subsequent several solar system objects, and close to the solar value, evolution (Section 2.4) are constrained by its present-day possibly reﬂecting an atmospheric methane inventory that interior structure and dynamics. In this Section, we focus on has not been aﬀected by atmospheric escape. The anal- three questions which are of particular relevance to Titan’s ysis of Mandt et al. (2009) indicates that if the current history. atmospheric methane is the remnant of methane injected into the atmosphere more than about two hundred million 1. To what extent is Titan diﬀerentiated? The degree of 12 13 years ago the C/ Cshouldbedecreasedbyatleast diﬀerentiation depends mainly on the thermal condi- 10% compared to the initial value. The current atmospheric tions during and just after accretion. Thus, Titan’s methane may then have been recently delivered to the sur- degree of diﬀerentiation constrains the conditions under face, from whatever reservoir. The detection of a signiﬁcant which it formed and evolved during the ﬁrst billion years. amount of 40Ar (the decay product of 40Kinitiallycon- tained in silicate phases) further indicates that exchanges 2. Is the ice shell convecting? The long-term thermal evolu- with the rocky interior have occurred (Niemann et al., 2010). tion of Titan’s interior is determined by the rate at which heat can be transferred across its icy outer layer(s). Thus, In this chapter we argue that exchange between the in- whether or not convection is taking place controls how terior and atmosphere of Titan is likely to have occurred rapidly Titan cools. If an ocean is present, the lifetime throughout Titan’s evolution, at least during accretion and of the ocean (and thus its habitability) is controlled by in the following billion years, and possibly until the present. the behaviour of the overlying ice shell. The ice shell Abetterunderstandingoftheseexchangeprocessesrequires also controls transport between the interior and the at- good knowledge of the interior structure. Based on its mean mosphere, as well as potential tectonic and cryovolcanic −3 density of 1881 kg.m ,Titanhaslongbeenknowntobe activities. composed of a roughly 60%-40% mixture of rock and ice by mass. Gravimetric and altimetric data collected by the 3. Does Titan possess a sub-surface ocean? If an ocean is Cassini spacecraft permit us to constrain the distribution present, Titan’s astrobiological potential is increased. of these components in the interior. However, as is further Moreover, an ocean can have profound geophysical con- discussed in Section 2, the solution remains non-unique sequences; for instance, it will likely increase the rate of and the diﬀerent possible structures of the interior imply tidal dissipation in the overlying ice shell. arangeofformationandevolutionscenarios.Otherdata from the Cassini-Huygens mission suggest the presence of a Prior to the Cassini-Huygens mission, there were essen- water ocean, possibly doped with ammonia or other agents tially no observational constraints on any of these questions. depressing the ocean crystallization point, beneath the solid Below we discuss the data sets relevant to the internal struc- icy shell. The existence of such an internal water ocean at ture and dynamics of Titan. present has important consequences for the evolution of Titan as well as for its astrobiological potential. 2.2.1 Interior structure inferred from the gravity In Section 2.2, we review the present-day constraints on ﬁeld data the structure and dynamics of Titan’s interior and present diﬀerent possible interior models. The present-day com- Measurements of Titan’s gravity ﬁeld are made by pre- position and structure of the atmosphere and surface are cise radio tracking of the Cassini spacecraft as it passes already discussed in other chapters and will not be repeated close to Titan. Unfortunately, because of the lack of a scan here. A good knowledge of the present-day state of Titan platform, gravity tracking precludes the acquisition of most is a necessary preliminary to understanding its origin and other observations, therefore the number of ﬂybys dedicated evolution. Section 2.3 presents some theoretical constraints to these measurements is limited. Furthermore, Titan’s ex- on the formation of the Saturnian system and describes tended atmosphere can result in non-gravitational forces on the accretion process(es) of Titan. Coupled evolution of the the spacecraft, which complicates the analysis and requires interior, surface and atmosphere, including internal diﬀer- that the closest approach distance is farther from Titan entiation and generation and recycling of the atmosphere, than what is desirable for an optimal resolution of the are then discussed in Section 2.4. Finally, we conclude in gravity ﬁeld. Despite these obstacles, Titan’s degree-two Section 2.5 by listing a series of remaining questions and gravity coeﬃcients have been determined with relatively how they may be addressed by future exploration missions. high conﬁdence, as well as the degree-three coeﬃcients but with much less conﬁdence (Iess et al., 2010). The degree-two spherical harmonic coeﬃcients, which are the dominant co- eﬃcients for a body subjected to tidal and rotational forces, J =( C )andC are tabulated in Table 2.1. Because 2 − 20 22 of the number of ﬂybys, the coeﬃcients were determined in- 14 Tobie, Lunine, Monteux, Mousis & Nimmo dependently, without requiring a value of 10/3 for the ratio Table 2.1. Observed (Clm)un-normalizeddegree-twograv- of J2 to C22.Thissituationdiﬀersfrombodieswherefewer ity coeﬃcients. Observed values are from Iess et al. (2010), ﬂybys have been carried out, such as Callisto (Anderson where SOL1 and SOL2 refer to the multi-arc and global et al., 2001) and Rhea (Mackenzie et al., 2008). solutions, respectively, and errors are given as 2σ.Un- h ± normalized surface topography coeﬃcients Clm are derived The degree-two coeﬃcients provide very important con- from data detailed in Zebker et al. (2009b) and quoted from straints on the radial mass distribution within the body. It Nimmo and Bills (2010). Normalized moment of inertia 2 can be shown mathematically that they are directly related C/MT RT was calculated assuming hydrostatic equilibrium to the principal moments of inertia (MoI), A, B and C: using equation 2.3. −6 −6 h (A + B)/2 C B A Clm 10 Clm 10 C obs obs × × lm C20 = 2 − ; C22 = − 2 (2.1) MT a¯ 4MT a¯ SOL1 SOL2 m l =2,m=0 -31.80 0.80 -33.46 1.26 -358 with MT Titan’s mass anda ¯ the mean equatorial radius. ± ± l =2,m=2 9.98 0.08 10.02 0.14 63 Thus, measuring the C20 and C22 coeﬃcients only gives the ± ± ratio -3.19 0.08 -3.34 0.13 -5.68 diﬀerence between the principal moments of inertia. In or- ± ± f der to determine their absolute value, a third condition is k2 (C20) 0.965 1.015 - needed. For Mars and the Earth, measurement of the pre- f k2 (C22) 1.010 1.013 - cession rate (which gives (C A)/C,whereC and A are the C/M R2 (C ) 0.335 0.342 - − T T 20 polar and averaged equatorial moments of inertia) allows 2 C/MT RT (C22) 0.341 0.342 - retrieval of the exact values of the three principal moments of inertia. In the case of Titan, such a measurement is not available. To infer the moment of inertia, one approach is to use the Radau-Darwin approximation, which assumes that degree two gravity coeﬃcients to ﬁrst order to estimate the body responds to tidal and rotational forces as a ﬂuid the ﬂuid Love number and the corresponding moment of body (i.e. it has no long-term strength, and therefore it is in inertia. Depending on whether we use the C20 or C22 coef- hydrostatic equilibrium). This approximation assumes that ﬁcients as a reference to determine the ﬂuid Love number, the body surface as well as each interface between internal the normalized moment of inertia varies between 0.335 and layers correspond to an iso-potential surface. The potential 0.342. This may be compared with the value of 0.4 for a at each interface is the sum of the tidal potential, centrifugal uniform sphere, 0.394 for the Moon, 0.33 for the Earth, 0.31 potential and gravitational potential. If the body is in hy- for Ganymede and 0.355 for Callisto (assuming hydrostatic drostatic equilibrium, the degree-two gravitational potential equilibrium (Anderson et al., 1996, 2001)). The moment is proportional to the tidal and centrifugal potential. The re- of inertia factor of Titan is therefore intermediate between sponse of the body to the tidal and centrifugal degree-two Ganymede and Callisto, suggesting a modest increase of potential is classically described with the ﬂuid Love number f density toward the center. k2 .Inhydrostaticequilibrium, α q α Cth = + kf ; Cth = kf (2.2) As already mentioned, this interpretation however as- 20 − 2 3 2 22 4 2 obs obs sumes that the C20 or C22 coeﬃcient is entirely deter- “ ” 3 where α = MS/MT (RT /DST ) ,isthetidalparameter, mined by the dynamical distortion of the satellite under 2 3 × and q = ω RT /GMT is the rotational parameter (with MS the action of tidal and rotational forces. In reality, part Saturn’s mass, RT Titan’s mean radius, DST Saturn-Titan of the gravity signal may be attributed to non-hydrostatic mean distance, ω Titan’s spin rate). In the case of a syn- contributions such as uncompensated relief, ice shell thick- chronously rotating body in a circular orbit, α = q and ness variations or lateral density variations. The fact that obs obs therefore C /C = 10/3ifthebodyishydrostatic.For the C20 /C22 ratio is not exactly -10/3 and that the 20 22 − Titan, α = q =3.9555 10−5. degree three gravity coeﬃcients are not zero clearly indi- × cates that there are some non-hydrostatic contributions to In the Radau-Darwin approximation, the normalized mo- the observed signals and that the gravity ﬁeld of Titan 2 ment of inertia about the spin axis, C/MT RT can then be is not only attributed to the dynamical distortion of a calculated from the ﬂuid Love number, inferred from the radially-distributed mass interior. The degree-three coeﬃ- gravity coeﬃcients assuming that they result from a ﬂuid cients indicate that non negligible lateral variations in the response of the body and in the limit of small density vari- mass distribution exist on Titan. These should aﬀect the ations: estimation of the moment of inertia. If we assume that the 1/2 non-hydrostatic contributions to degree two are comparable C 2 2 5 = 1 1 (2.3) to the inferred degree 3 terms, a few to 10% of the observed M R2 3 2 − 5 f − 3 T T k2 +1 ! degree two signals may be attributed to the non-hydrostatic 4 5 (Iess et al., 2010, Supplementary Information). According to the Radau-Darwin equation, a 5%-10% overestimate of obs 2 In the case of Titan, the ratio between the observed C20 the hydrostatic parts would shift C/MRRT from 0.341 to obs and C22 is close to -10/3, suggesting that Titan is relatively 0.334-0.327. close to hydrostatic equilibrium. We can therefore use the The Origin and Evolution of Titan 15
In summary, the interpretation of the gravity data indi- of thickness D 100 km, heated from below at a rate of ≈ cates that the normalized moment of inertia of Titan is of 4-5 mW m−2 (appropriate to a chondritic Titan) and from the order of 0.33-0.34. This MoI value combined with the within at a rate about one-tenth as large (due to tidal mean density of Titan indicates that the density of the core dissipation). The tidal Love number h2,whichdescribesthe is low, which may be explained by incomplete separation of response of the body to the eccentricity tides, was found to rock from ice (Iess et al., 2010) or the presence of highly hy- be 1.2, consistent with theoretical predictions (Sohl et al., ≈ drated silicate minerals (Castillo-Rogez and Lunine, 2010; 2003). Given the relatively modest rate of tidal dissipation, Fortes, 2011). Based on the present data, the presence of a the e-folding timescale to damp Titan’s eccentricity is quite small iron core (< 500 km) cannot be ruled out (Castillo- long (2-3 Gyr). Alternatively, Choukroun and Sotin (2012) Rogez and Lunine, 2010; Fortes, 2011), but the low internal proposed that the ﬂattened shape of Titan may result from temperatures ( < 900 1000K) required to maintain either the accumulation of dense hydrocarbon clathrates at high − an ice-rock mixture or highly hydrated silicate minerals are latitudes owing to enhanced ethane precipitation. Their incompatible with the segregation of iron and the formation Pratt isostatic compensation model may explain the 300 ∼ of a iron core. The most likely present-day structures is ei- mdiﬀerencebetweenexpectedandmeasuredpolarradiiif ther a pure rock core, mostly anhydrous, surrounded by a the ethane-rich clathrate crust is about 3-km thick. mixture of ice and rock, surrounded by pure water ice (Fig. 2.5, top and middle), or a low density rock core consisting Although the variable-thickness or variable-density shell mainly, but not entirely, of hydrated silicate minerals sur- is assumed isostatic, it will still have a small eﬀect on the rounded by H2O(Fig.2.5,bottom).Inbothcases,aliquid degree-2 gravity coeﬃcients. Correcting for this eﬀect before water layer may be present between a high-pressure ice man- using equation (2.3) changes the resulting moment of inertia tle and an outer ice shell. We will discuss the possible evo- by about 2% (the MoI factor is still between 0.33 and 0.34). lution scenarios leading to these interior structures in more An isostatic shell is probably a good assumption: a shell details in Section 2.4. with an elastic thickness of 50 km would still be about 95% compensated at degree-2 (Turcotte et al., 1981), assuming aYoung’smodulusof9GPa. 2.2.2 Global shape: Constraints on the structure and thermal state of the ice shell Perhaps the most surprising result of these studies is that, Just as with gravity, the degree-two shape of a synchronous if it is correct, Titan’s ice shell must be conductive, or body can be used to determine its moment of inertia - as at most weakly convecting, in order to support thickness long as the body is in hydrostatic equilibrium. As with or density variations on the long term. In the thickness- gravity, the degree-two shape coeﬃcients of such a body, variation model, surface topography generates pressure gra- h h C20 and C22,willbeintheratio-10/3.Thus,withsuﬃ- dients at depth; if the ice shell were convecting, the ice vis- ciently good topography we can check whether or not the cosity would be low enough to ﬂow in response to these hydrostatic condition is satisﬁed. pressure gradients and progressively remove the surface to- pography (Stevenson, 2000). Whether or not an ice shell Titan’s shape has been measured (Zebker et al., 2009b) by convects depends on whether its Rayleigh number Ra of the acombinationoftraditionalradaraltimetry(Zebkeretal., whole ice layer exceeds some critical value Racr (e.g. Mitri 2009a), and a new technique which uses overlaps in the and Showman, 2008): SAR-imaging swaths to determine topography (Stiles et al., h 2009). In contrast to the results obtained from gravity, C20 h 3 3 14 and C22 were not in the expected -10/3 ratio (see Table 2.1). ρgα∆TD 8 D 10 Pas Ra = 3 10 Racr κη ≈ × 100km η ≥ b „ « „ b « This result is apparently paradoxical: how can Titan have (2.4) nearly hydrostatic gravity but non-hydrostatic topography? Here ρ,α,κ,ηb are the density, thermal expansivity, ther- The hydrostatic gravity rules out scenarios in which Titan’s mal diﬀusivity and basal viscosity of the ice shell, respec- shape is derived from an earlier epoch (a ”fossil bulge” tively, ∆T is the temperature contrast across the shell similar to that of the Moon). One possibility is that Titan’s (∆T 170 180 K between the surface and the ocean) and & − outer ice shell - presuming the existence of an ocean - is numerical values assumed are from Nimmo and Bills (2010). of varying thickness or density and is (Airy or Pratt) iso- The value of Racr depends on how sensitive viscosity is to statically compensated. An isostatic shell would give rise temperature changes, and the geometry of the shell (Solo- to only very small gravity anomalies (see below), while the matov, 1995; Barr and McKinnon, 2007), and for Titan is shell thickness variations would result in a long-wavelength 4 107 in the stagnant lid regime. Equation (2.4) shows ≈ × non-hydrostatic shape. that whether or not the shell convects is a sensitive function of the viscosity of the base of the ice shell. For pure ice, the Nimmo and Bills (2010) explored a version of the Airy viscosity in turn depends mainly on the temperature of the scenario, in which shell thickness variations were caused ocean beneath, and secondarily on the ice grain size. The by spatial variations in tidal heating in the ice shell (cf. ocean temperature is controlled by how much ammonia is Ojakangas and Stevenson, 1989). They found that the ob- present; a suﬃciently ammonia-rich ocean and/or large ice served topography could be ﬁtted with a mean ice shell grain sizes would prevent convection because the ice would 16 Tobie, Lunine, Monteux, Mousis & Nimmo be too viscous (e.g. Tobie et al., 2005b). 2.2.3 Geophysical evidence for an internal water ocean Nevertheless, it is possible that the ice shell was convec- For the Galilean moons, the most convincing evidence for tive in the past and that convection ceased as the ocean subsurface oceans came from analysis of the magnetic data crystallized and the ammonia concentration increased ow- acquired by the Galileo mission during close ﬂybys. The tilt ing to reduced ocean volume (Tobie et al., 2005b; Mitri of Jupiter’s magnetic ﬁeld relative to its spin pole means and Showman, 2008). The elevated orbital eccentricity of that the Galilean satellites experience a time-varying ﬁeld Titan indicates however that this convective period should and resulting induction currents. This eﬀect has been used have lasted less than about 0.5-1 billion of years (Tobie to probe the conductivity structure of the satellites, and et al., 2005b). In contrast to the Galilean moons where the deduce the presence of oceans for the three outer satel- Laplace resonance excites the eccentricity of Europa and lites (Khurana et al., 1998; Kivelson et al., 2002), and a Io, there is no major resonance in the Saturnian system partially-molten mantle for Io (Khurana et al., 2011). Un- to force Titan’s eccentricity. Unless some unknown pro- fortunately, Saturn’s spin and magnetic poles are almost cess has recently forced the eccentricity, the latter should exactly coplanar, so the induction technique is very diﬃcult have monotonically decayed as a function of time since the to use at Titan (Saur et al., 2010). A recent study of Titan’s formation of the system owing to tidal dissipation at the magnetic ﬁeld did not see any induction eﬀects, and did not surface and in the interior (Sears, 1995; Sohl et al., 1995; detect any permanent internal magnetic ﬁeld (Wei et al. Tobie et al., 2005b). The dissipation rate is reduced either 2010). Although it is possible to have a molten core without if the interior is totally rigid (with no internal ocean) (Sohl generating a magnetic ﬁeld (e.g. Venus, Io), this result is et al., 1995) or if the outer ice shell above the ocean is thin perhaps consistent with Titan’s interior being cold and not and/or cold (i.e. conductive) (Tobie et al., 2005b). Indeed, fully diﬀerentiated. the volume wherein tidal energy is dissipated is much bigger in a convective ice shell (where about half of the shell has On Titan, another technique based on the measurements atemperaturelargerthan240K)thaninaconductiveice of the electric ﬁeld during the Huygens descent by the Per- shell where the warm temperature is conﬁned at the bottom mittivity, Wave and Altimetry (PWA) instrument (B´eghin of the ice shell. et al., 2007) provides however some information on the subsurface electrical conductivity. The measurements of If Titan has an internal ocean at present, it implies low frequency waves and atmospheric conductivity by the that the ice shell remained conductive and relatively thin PWA instrument revealed the existence of a Schumann-like (< 30 40 km) during most of Titan’s evolution, in order − resonance trapped within Titan’s atmospheric cavity. The to reduce the eccentricity damping (Tobie et al., 2005b). observed signal may be triggered and sustained by strong This is possible if the viscosity of the icy shell is large electric currents induced in the ionosphere by Saturn’s mag- enough due to large grain sizes or to low oceanic tempera- netospheric plasma ﬂow (Sim˜oes et al., 2007; B´eghin et al., tures (owing to ammonia Tobie et al. (2005b) or methanol 2009, 2010). The characterisitics of this trapped-resonant Deschamps et al. (2010)) and/or if the icy shell contains mode imply the presence of a conductive layer at about large fraction of low conductivity materials such as methane 30-60 km below the surface (B´eghin et al., 2010), which clathrate (Tobie et al., 2006). Following this requirement, may correspond to the ice/ocean interface provided that the icy shell has never been convecting or has been convect- the ocean is suﬃciently conductive, possibly doped with ing only during a limited period of time. A possible scenario small amounts of ammonia and/or salts (B´eghin et al., may be that the ice shell started convecting during the 2010). last roughly billion years and then stopped later because of ocean crystallization as the concentration in ammonia or As explained in Section 2.2.2, the observed topography methanol in the ocean became too high and the temper- may be explained by variations in ice shell thickness, re- ature too low to allow convection in the overlying viscous sulting from inhomogeneous crystallization of an underlying ice shell (e.g. Mitri and Showman, 2008). Indeed, even for ocean. If this analysis is correct, it implies that an inter- initially moderate ammonia fraction (about 2-3% relative to nal ocean is still present in Titan’s interior and that it is the total mass of H O), the concentration signiﬁcantly in- 2 currently slowly crystallizing. This constitutes additional creases as the ocean crystallizes and the ocean temperature indirect evidence for an internal ocean inside Titan. gets below 250 K when the ammonia concentration in the crystallizing ocean reaches 10 wt% (Tobie et al., 2005b). At Another piece of evidence is provided by Titan’s spin such a temperature, the viscosity at the base of the ice shell state. A non-spherical body, such as a synchronous satel- exceeds 5.1015 1016 Pa.s, which is too large to promote − lite, precesses about its rotation axis at a rate determined convection even for a 100 km-thick ice shell. by the ratio, (C A)/C.Sincethegravitycoeﬃcient − J =(C A)/M R2 ,ifJ and the precession rate can be 2 − T T 2 measured, the moment of inertia C may be determined di- rectly without assuming hydrostatic equilibrium. Although measuring a satellite’s precession rate directly is diﬃcult, measuring the obliquity (the angle between spin pole and The Origin and Evolution of Titan 17 orbit pole) is relatively easy. If a satellite is dissipative, tained during six close ﬂybys indicates that Titan responds tidal torques drive the obliquity to a state in which the spin to the variable tidal ﬁeld exerted by Saturn at a detectable pole and orbit pole remain coplanar as the latter precesses level, providing further evidence for a global ocean at depth around a ﬁxed (invariable) pole (e.g. Peale, 1969). If this on Titan (Iess et al., 2012). A future mission with dedicated state, known as a damped Cassini state, exists, then mea- multiple ﬂybys or even better in orbit around Titan will be suring the obliquity is suﬃcient to determine the satellite required to conﬁrm this detection and to get more precise moment of inertia (e.g. Bills and Nimmo, 2011). In practice, informations on the ocean and ice shell thicknesses. other small eﬀects (such as Saturn’s spin pole precession) have to be taken into account, but they do not alter the underlying physical picture.
Titan’s spin pole location has been established through analysis of SAR images. It is very nearly coplanar with the invariable and orbit poles (suggesting occupation of a Cassini state), and corresponds to an obliquity of 0.32◦ (Stiles et al., 2008, 2010). If Titan occupies a damped Cassini state, the implied value of C/(M R2 )is0.45 0.02 T T ± 2.2.4 Summary (Bills and Nimmo, 2011). This result is clearly unreason- able, in that it suggests a body which is under-dense in At the start of this Section we posed three questions; we the interior, and also disagrees with the results derived will conclude by providing some provisional answers. from gravity. The solution is probably a combination of two factors. First, the fact that Titan’s spin pole is not 1. To what extent is Titan diﬀerentiated? Based on the quite coplanar with the orbit and invariable poles suggests gravity results (Section 2.2.1), the normalized moment of some degree of present-day excitation of the obliquity. Two inertia of Titan is estimated between 0.33 and 0.34. This possible candidate sources of excitation are the atmosphere clearly indicates that Titan is not fully diﬀerentiated (Tokano et al., 2011) or ﬂuid motion in a subsurface ocean, and does not have an iron core like Ganymede. Earlier both of which can transfer angular momentum to the ice we suggested two possible internal structures – one in shell. The second factor is that Titan is probably not acting which rock and ice are incompletely separated, the other as a rigid body (which the Cassini state analysis assumes), involving a rocky core that is partially hydrated – that and that instead the ice shell is partially decoupled from the satisfy the rather high moment of inertia. We discuss interior by an ocean. If the shell were totally decoupled, its this further in Sections 2.3.4 and 22.214.171.124. normalized moment of inertia would be 0.67. The derived value of 0.45 suggests some degree of coupling, perhaps due 2. Is the ice shell convecting? to pressure or gravitational torques (Baland et al., 2011). Aconductiveiceshellisfavouredbythemodelfor Fortunately, continued monitoring of the position of Titan’s Titan’s long-wavelength shape (Section 2.2.2) and is spin pole should be able to disentangle these two eﬀects, but consistent with the requirement that its eccentricity not the current results are strongly indicative of a subsurface damp too rapidly. To avoid convection, the ice shell must ocean. have a high viscosity, which can be achieved if the ocean beneath is ammonia-rich and therefore cold. Convection Another potential source of evidence for a subsurface may have occurred in the past before the ammonia con- ocean would be the existence of non-synchronous rotation centration increased, but this convective period must (NSR) of the ice shell (Ojakangas and Stevenson, 1989). have lasted less than 0.5-1 billion years in order to ex- It was initially thought that Titan’s shell was exhibiting plain Titan’s present-day elevated orbital eccentricity. such behaviour (Lorenz et al., 2008b), but this observation turned out to be erroneous and there is currently no evi- 3. Does Titan possess a sub-surface ocean? Electric signals dence for NSR (Stiles et al., 2010). measured in Titan’s atmosphere by Huygens suggest the presence of a conductive layer at depth, possibly a water Probably, the most promising technique to conﬁrm the ocean doped with small amounts of salts or ammonia existence of a subsurface ocean would be to measure the to increase the electrical conductivity. Titan’s spin state tidal ﬂuctuation from accurate gravity or topography mea- and tidal gravity response are also suggestive of a de- surements. The Radio Science Experiment on Cassini has coupled shell (Section 2.2.3) while its long-wavelength in theory the capability to measure the tide-generated grav- shape can be explained if a subsurface ocean is assumed ity signals (Rappaport et al., 1997, 2008). However, the (Section 2.2.2). Given the existence of subsurface oceans atmospheric drag due to the extended atmosphere as well on the comparably-sized Galilean satellites, Callisto and as the existence of relatively large signals in harmonics of Ganymede, it would be surprising if Titan did not pos- degree three (Iess et al., 2010), suggesting signiﬁcant signal sess an ocean. However, one important diﬀerence is that at degrees 4 and 5 (Sotin et al., 2010), make the analysis Titan’s ocean may be colder and more ammonia-rich (see more diﬃcult than initially anticipated. In spite of these below). diﬃculties, a recent analysis of the Cassini gravity data ob- 18 Tobie, Lunine, Monteux, Mousis & Nimmo
2.3 ORIGIN OF THE SATURNIAN SYSTEM that O, C, and N exist only in the form of H2O, CO, CO2, AND ACCRETION OF TITAN CH3OH, CH4,N2,andNH3.TheabundancesofCO,CO2, CH3OH, CH4,N2 and NH3 are then determined from the Titan was formed as a regular satellite in a disk that was the adopted CO/CO2/CH3OH/CH4 and N2/NH3 gas phase outgrowth of the formation of Saturn itself. The conditions molecular ratios. Once the abundances of these molecules under which Saturn formed therefore had a direct impact are ﬁxed, the remaining O gives the abundance of H2O. on the formation of Titan. They determined its composi- Concerning the distribution of elements in the main volatile tion, the architecture of the whole system as well as the molecules, the ratio of CO/CO2/CH3OH/CH4 can be set ﬁnal assemblage that led to the formation of Titan and of to 70/10/2/1 in the gas phase of the disk, values that are the other regular moons. The present section is structured consistent with the ISM measurements considering the con- around three main questions which are of particular rele- tributions of both gas and solid phases in the lines of sight vance to Titan’s formation. (Mousis et al., 2009c). In addition, S is assumed to exist in the form of H S, with H S/H =0.5 (S/H )#,andother 2 2 2 × 2 1. What was the composition of the primordial building refractory sulﬁde components (Pasek et al., 2005). N2/NH3 blocks that formed Titan and the Saturnian system? is assumed equal to 10/1 in the nebula gas-phase, a value Titan and the other Saturnian moons formed from a compatible with thermochemical models of the solar nebula mixture of rock and ice initially present in the vicinity (Lewis and Prinn, 1980) and with observations of cometary of Saturn. We will present in this section diﬀerent the- comae (e.g. Hersant et al., 2008). oretical and observational constraints that can be used to estimate the composition of these ice-rock blocks, To ﬁrst order, the process by which volatiles are trapped notably their volatile content. in icy planetesimals can be calculated assuming thermo- dynamic equilibrium. The sequence of clathration1 and 2. What makes the Jupiter and Saturn systems so diﬀerent? condensation can be determined by using the equilibrium Contrary to the Jovian system, there is no gradual varia- curves of stochiometric hydrates, clathrates and pure con- tion of density and hence rock/ice ratio in the Saturnian densates, and the thermodynamic path (hereafter cooling moons as a function of planet distance. Titan is the curve) detailing the evolution of temperature and pressure unique big moon, the second largest moon Rhea having between roughly 5 and 20 AU, a distance range largely amasslessthan2%ofTitan’smass.Anotherparticular encompassing the migration path followed by proto-Jupiter characteristic of Saturn is its extended ring system. In and proto-Saturn during their formation in the solar nebula this Section, we will describe the main phases of satellite (Alibert et al., 2005b). We refer the reader to the works of formation in a disk around giant planets and try to iden- Papaloizou and Terquem (1999) and Alibert et al. (2005a) tify which processes diﬀered between Jupiter and Saturn. for a full description of the turbulent model of accretion disk used here. For the pressure range expected in the 3. Can Titan accrete undiﬀerentiated? The Cassini gravity solar nebula, most of the gas compounds, except CO2,are data suggests that Titan is moderately diﬀerentiated at enclathrated at temperatures higher than the temperatures present (section 2.2.1), suggesting that the diﬀerentiation they condense in the form of pure condensates. This has process has been limited. One possibility is that Titan fundamental consequences for the trapping of volatiles in accreted relatively cold, which would have delayed the primordial building blocks. diﬀerentiation. To address that question, we will review the physics of accretion and determine under which con- As illustrated in Figure 2.1, the cooling curve intercepts ditions Titan might have remained cold during accretion. the equilibrium curves of the diﬀerent ices at particular tem- peratures and pressures. For each ice considered, the domain of stability is the region located below its corresponding equilibrium curve. The clathration process stops when no 2.3.1 Formation of ices in Saturn’s environment more crystalline water ice is available to trap the volatile The primitive building blocks that formed Titan and the species. As CO2 is the only species that crystallizes at a rest of the satellite system were presumably composed of higher temperature than its associated clathrate, solid CO2 acomplexassemblageofvariousices,hydrates,silicate is assumed to be the only existing condensed form of CO2 minerals and organics. An important issue associated with in this environment. In addition, CH3OH is assumed to be the formation of giant planets and their satellite systems in the form of pure ice, since, to the best of our knowledge, concerns the condensation of water ice and other ice com- no experimental data concerning the equilibrium curve of pounds as the solar nebula cools down. The condensation its associated clathrate have been reported in the literature. sequence depends on the initial gas phase, the cooling rate of the disk and the stability of the gas molecules in solid phases. 1 Clathrate hydrates are crystalline water-based structures that form cages in which gas molecules can be trapped (Lunine and Stevenson, 1985; Choukroun et al., 2011), and for this reason The composition of the initial gas phase of the disk can they play a crucial role in the volatile trapping. They also play a be deﬁned by assuming that the abundances of all elements, key role in the evolution of the volatile inventory on Titan asit including oxygen, are protosolar (Asplund et al., 2009) and will be further discussed in Sections 126.96.36.199 and 188.8.131.52. The Origin and Evolution of Titan 19
Figure 2.2 Mole fraction of volatiles encaged in H2S(a),Xe (b), CH4 (c) and CO (d) clathrates. Grey and dark bars correspond to structure I and structure II clathrates, Figure 2.1 Equilibrium curves of NH3-H2Ohydrate,H2S, Xe, respectively, which are two possible clathrate structures with CH4 and CO clathrates (solid lines), CH3OH, CO2,Kr,CO,Ar slightly diﬀerent cage sizes (e.g. Sloan, 1998). and N2 pure condensates (dotted lines), and thermodynamic path followed by the primordial nebula between 5 and 20 AU as afunctionoftime,respectively,assumingafulleﬃciencyof curves established experimentally for single guest molecules clathration. Species remain in the gas phase above the (similar to the curves shown on Fig. 2.2). This simpliﬁed equilibrium curves. Below, they are trapped as clathrates or simply condense. The equilibrium curves of hydrates and approach permits an understanding to ﬁrst order of the con- clathrates derive from Lunine and Stevenson (1985)’s densation sequence in the nebula, but in reality, clathrates compilation of published experimental work, in which data are forming in a gas mixture always simultaneously trap several available at relatively low temperatures and pressures. The guest molecules (e.g. Lunine and Stevenson, 1985; Sloan, equilibrium curves of pure condensates used here derive from 1998). Here we consider the eﬀect of simultaneous trapping the compilation of laboratory data provided by Lide (2004). of guest species on the predicted composition of volatiles trapped in the water ice.
For simplicity, the clathration eﬃciency is assumed to be The composition of the clathrate phases can be con- total, implying that guest molecules had the time to diﬀuse strained by using a statistical thermodynamic model de- through porous water-ice solids before their growth into rived from the approach of van der Waals and Platteeuw planetesimals and their accretion by proto-planets or proto- (1959) (see Lunine and Stevenson, 1985; Mousis et al., satellites. This statement remains plausible only if collisions 2010, for details). This thermodynamic approach is based between planetesimals have exposed essentially all the ice on the use of intermolecular potentials, which themselves to the gas over time scales shorter or equal to planetesimal depend on parameters describing the interaction between lifetimes in the nebula (Lunine and Stevenson, 1985). In the the guest molecule and the cage made of water molecules, present case, which is based on the protosolar abundances called “Kihara parameters”. The afore-mentioned forma- of Asplund et al. (2009), NH ,H S, Xe, CH and 31.5% tion sequence of the diﬀerent ices occurring during the 3 2 4 ∼ of CO form NH3-H2OhydrateandH2S, Xe, CH4 and CO cooling of the solar nebula is considered (see Fig. 2.1). As clathrates with the available water in the outer nebula. The aconsequence,anyvolatilealreadytrappedorcondensed remaining CO, as well as N2,Kr,ArandNe(notshown at a higher temperature than the formation temperature on the ﬁgure), whose clathration normally occurs at lower of the clathrate under consideration is excluded from the temperatures, remain in the gas phase until the nebula coexisting gas phase composition. This implies that CO2, cools enough to allow the formation of pure condensates. Xe, CH4,CO,Kr,ArandN2 are considered as possible As the gas phase composition of the disk is assumed not to guests in the case of H2Sclathrate.Ontheotherhand, vary with heliocentric distance, the calculated clathration only N2,Ar,KrcanbecomeguestsinCOclathrate.Figure conditions remain the same in the 5–20 AU range of the 2.2 represents the mole fraction f of volatiles encaged in nebula (see Marboeuf et al. (2008) for details). Once crys- structure I and structure II clathrates aprioridominated tallized, these ices will agglomerate and form planetesimals by H2S, Xe, CH4 and CO and formed in the primordial large enough to decouple from the nebular gas and will be nebula. This ﬁgure shows that H2S, Xe, CH4 and CO are accreted by the forming giant planets and satellites. the dominating guest species if structure I clathrates form in the nebula. However, in the case of structure II clathrate Many works published in the last decade and detailing formation, CO becomes a minor compound in the clathrate the formation conditions of ices in the solar nebula (Gautier that is expected to be dominated by this molecule. Indeed, et al., 2001; Hersant et al., 2008; Mousis and Alibert, 2006; Kr becomes more abundant than CO in this clathrate, Alibert and Mousis, 2007, e.g.) have used the equilibrium irrespective of their initial abundances in the gas phase 20 Tobie, Lunine, Monteux, Mousis & Nimmo of the nebula. On the other hand, structure I should be also seems consistent with the composition of Enceladus’ the most likely form of clathrate produced in the nebula plumes, where CO2,NH3,CH4 and H2Shavebeendetected since H2S, Xe, CH4 and CO clathrates are all of structure by the Cassini Ion and Neutral Mass Spectrometer (INMS) I(LunineandStevenson,1985).Figure2.2alsoshowsthat, (Waite et al., 2009). If this is correct, it would imply that depending on the considered structure, the Kr/CH4 ratio the solids that enrich Saturn’s envelope in heavy elements in CH4–dominated clathrate is order of or greater than the have a composition close to the ones that formed the regular value derived from the nebula gas phase composition used satellites. in our model. This implies that Kr can be enclathrated at atemperatureof 55 K instead of condensing in the 20–30 ∼ Krangeinthesolarnebula(seeFig.2.1). 184.108.40.206 Structure and evolution of Saturn’s subnebula At the beginning of Saturn’s formation process, the total ra- dius of the planet equals its Hill’s radius, and no subnebula 2.3.2 Formation of Saturn and the circumplanetary can exist. At the end of the formation process, the cooling disk: Consequences for the delivery of volatiles rate of the planet’s envelope is such that its radius shrinks faster than the disk can supply additional gas. Accretion of 220.127.116.11 Saturn’s formation and envelope enrichment gas now proceeds through streamers which eventually col- Favored theories of giant planet formation center around lide and form a circumplanetary disk (Lubow et al. 1999). two main paradigms, namely the core accretion model When the radius of proto-Saturn becomes small enough, a (Safronov, V. S., 1969; Goldreich and Ward, 1973; Pol- subnebula emerges from the contracting atmosphere. This lack et al., 1996) and the gravitational instability model subnebula is fed by gas and solids accreted from Saturn’s (Cameron, 1978; Boss, 1997). In the frame of the core feeding zone. The accretion rate of gas from the solar nebula accretion model, a solid core forms from the accretion of to the subnebula is therefore directly driven by the planet’s planetesimals and becomes massive enough ( 5–10 M⊕) accretion rate of gas, which is derived from giant planet ∼ to initiate runaway gravitational infall of a large gaseous formation models (e.g. Alibert et al., 2005a). envelope in which gas-coupled solids continue their accre- tion (Alibert et al., 2005b; Hubickyj et al., 2005; Mordasini Diﬀerent models of subdisks have been proposed in the et al., 2009). This model provides the large amount of heavy literature (e.g. Canup and Ward, 2002; Mosqueira and elements necessary to explain the supersolar metallicities Estrada, 2003a; Estrada and Mosqueira, 2006; Sasaki et al., observed in Jupiter and Saturn via the accretion of plan- 2010). Here we describe in more details the so-called gas- etesimals in their envelopes (Gautier et al., 2001; Saumon starved disk model initially introduced by Canup and Ward and Guillot, 2004; Alibert et al., 2005b; Mousis et al., 2006, (2002) and then used by Alibert et al. (2005a) and Alibert 2009c). In the frame of the gravitational instability model, and Mousis (2007). Alibert and Mousis (2007) described a gas giant protoplanets form rapidly through a gravitational model of Saturn’s subnebula whose evolution proceeds in instability of the gaseous portion of the disk and then more two phases. During the ﬁrst phase, when the solar nebula slowly contract to planetary densities (Boss, 1997, 2005). In is still present, the subnebula is hot and dense, and fed this scenario, due to the limited eﬃciency of planetesimals through its outer edge (ﬁxed to 1/5 Hill’s radius – see accretion during the planet formation, its metallicity should Alibert and Mousis (2007) for details), at a rate controlled be almost equal to that of the parent star (Helled and by Saturn’s formation. This justiﬁes a posteriori the use of Bodenheimer, 2010). an equilibrium model as an initial model of the subnebula. This phase lasts until the solar nebula has disappeared By using the sequences of ice formation described in and the temporal evolution of the subnebula is therefore Section 2.3.1, it is possible to predict the abundances of directly governed by the last phase of Saturn’s formation. major species (C,N, O, S, P) and noble gases (Ar, Kr, Xe) The accretion rate as a function of the distance to Saturn in Saturn’s atmosphere in the context of the core-accretion is no longer constant and the temperature and pressure model (Hersant et al., 2008; Mousis et al., 2009c), and to conditions of the subnebula decrease rapidly. compare to the observed abundances when available. For instance, adjusting the model to the observed carbon abun- Figure 2.3 represents the temporal evolution of the tem- dance provides constraints on the minimum mass range perature proﬁle within the Saturn’s subnebula depicted by of ices required in the envelope (estimated between 14.3 Alibert and Mousis (2007). The solid lines are plotted, from and 15.6 M⊕,Mousisetal.(2009c).Suchanapproachalso top to bottom, at t =0Myr(correspondingtothetime provides information on the typical composition of the plan- when Saturn has accreted 70% of its ﬁnal mass), 0.1 Myr, etesimals in the feeding zone of Saturn, which in turn can 0.2 Myr, 0.3 Myr, 0.35 Myr, 0.4 Myr, 0.5 Myr, 0.6 Myr be used to constrain the composition of the building blocks and 0.7 Myr. The switch from phase 1 to phase 2 of the of Titan. The analysis of Hersant et al. (2008), based on a subnebula evolution occurs at 0.36 Myr. The dotted line simpliﬁed description of volatile trapping by clathration and gives the temperature 2000 years after the beginning of condensation, suggests that the planetesimals that reached phase 2, and illustrates the fast evolution of the subnebula Saturn’s envelope contained signiﬁcant amounts of CO2, during this stage. NH3,CH4 and H2S, and were depleted in CO and N2.This The Origin and Evolution of Titan 21
of solids to the disk is still under debate (see for instance Estrada et al. (2009) and Canup and Ward (2009)). In term of volatile delivery, each of these processes has very diﬀerent consequences. In mechanism (1), and to a lower extent in mechanism (2), the particles may preserve the volatile content they acquired in the solar nebula during the condensation sequence, while in mechanism (3) a large frac- tion of volatile would be likely released when incorporated into the disk. Even in mechanism (4), a signiﬁcant fraction may be lost during collisions at high velocity.
As a consequence, the volatile content in the building blocks of Titan may be signiﬁcantly smaller than their content in the planetesimals initially formed in the solar nebula in the vicinity of Saturn. Even for small particles entrained by the gas inﬂow, they could have been altered once embedded in the subdisk, if they encountered gas tem- perature and pressure conditions high enough to generate a loss of volatiles during their migration, and if they remained relatively porous, which seems a reasonable assumption. Mousis et al. (2009b), for instance, favored this mechanism Figure 2.3 Temperature inside the subnebula, as a function of to explain the carbon monoxide and argon deﬁciencies in the distance from Saturn (expressed in Saturn’s radii RSat)and the atmosphere of Titan. This scenario is based on the represented at diﬀerent epochs of its evolution (ﬁgure adapted fact that CO-dominated clathrate, CO, N2 and Ar pure from Alibert and Mousis (2007) – see text). condensates are not stable at temperatures greater than 50 K in the gas phase conditions of the nebula or the ∼ subnebula (see Fig. 2.1). This implies that if porous plan- 2.3.3 Delivery of solids and conditions of Titan’s etesimals ultimately accreted by Titan experienced intrinsic formation temperatures of 50 K during their migration and accretion ∼ Regardless the considered formation mechanism of Saturn’s in Saturn’s subnebula, they are expected to release most of satellites within the subnebula (accretion at their present their CO, Ar, Ne and N2 while CH4 and NH3 are expected orbit or migration and growth starting from the edge of the to remain trapped, or at least to partly re-condense if re- subnebula), the delivery of solids to the disk is an essential leased. Similarly, if planetesimals are captured by gas drag step, because it determines the amount of volatiles that can or collisional capture, the most volatile species, CO, N2, be retained and ultimately incorporated in the interiors of Ne, Ar should be released while the most stable one should these bodies. Diﬀerent processes have been advocated (see be preferentially retained. for instance Estrada et al. (2009) and Canup and Ward (2009) for a detailed discussion): (1) direct transport of Such a volatile loss scenario cannot however account for small particles into the disk with the inﬂowing gas, (2) the deﬁciency of Titan’s atmosphere in Kr and Xe because ablation and gas drag capture of planetesimals orbiting the Kr is signiﬁcantly trapped in CH4-dominated clathrate (see Sun through the gas-rich circumplanetary disk, (3) break- Fig. 2.2) and Xe is also enclathrated at temperatures higher up, dissolution and recondensation of planetesimals in the than 50 K. However, as is further discussed in Section ∼ extended envelope of the forming planet, (4) collisional 18.104.22.168, both Xe and Kr may be sequestred in Titan’s icy capture of planetesimals . shell and interior owing to clathration process, partly dis- solved in the hydrocarbon seas of Titan (Cordier et al., Mechanism (1) is valid for micron- to millimeter-sized 2010) or trapped in aerosols (Jacovi and Bar-Nun, 2008), particles that remain coupled to the subdisk’s gas while thus potentially explaining their deﬁciency while predicted mechanism (2) requires larger size particles ( 1m)to to be present in signiﬁcant amounts in the building blocks. ∼ allow their capture and inward drift due to gas drag. Mech- anism (3) is probably not valid for Saturn’s subnebula This scenario seems also problematic regarding the ten- because the D/H measurement performed at Enceladus by tative detection of 22Ne by the Huygens GCMS (Niemann the Cassini spacecraft (Waite et al., 2009), which is close et al., 2010). According to this scenario, the building blocks to the cometary value, suggests that the building blocks of should be strongly depleted in Ne with 22Ne/36Ar << 1, this satellite were formed in the solar nebula. In contrast, while the ratio estimated in the atmosphere from the Huy- planetesimals formed via this mechanism would have a gens GCMS measurements is close to 1 (Niemann et al., protosolar D/H ratio. Mechanism (4) applies for gas-poor 2010). Note, however, that the GCMS detection of 22Ne or no gas environments and may have taken place at epochs is only tentative, and the errors are still very large. In subsequent to the subnebula phase. In any case, the relative situ measurements of Ne, Ar, Kr and Xe abundances in contribution of each of these mechanisms to the delivery Saturn’s atmosphere by a future probe as well as in surface 22 Tobie, Lunine, Monteux, Mousis & Nimmo sampling on Titan could constrain the origin of their unex- tion factor Q (< 2000). This latter condition is required to pected abundances in Titan’s atmosphere. This would help provide suﬃcient tidal expansion in order for the satellites in understanding the process of noble gas trapping in the to move from the edge of a massive ring system to their Saturnian system, thus providing key information on the present orbital position. They also show that the accretion composition of Titan’s building blocks. of the satellites from the outer edge of such a ring system may explain the diversity of rock/ice ratios among the Another issue concerns the satellite growth from the mid-sized moons. Alternatively, the variation in rock/ice solids present in the disk. In a gas-rich disk, the growing fraction among the mid-sized satellites may be attributed solid objects should migrate inward as a result of gas drag to the stochastic nature of late-stage accretion, character- for small objects and then of type I and/or II migration ized by giant impacts, as suggested by Sekine and Genda for largest objects ( 1000 km) (e.g. Estrada et al., 2009; (2011). Unfortunately, the diﬀerent proposed models lack ∼ Canup and Ward, 2009). A growing satellite may survive strong observational discriminants. Estimation of Saturn’s if it grows large enough to open a gap in the disk, limiting Qfactorfromastrometricmeasurementsisprobablyone the migration (Mosqueira and Estrada, 2003b), or if the gas of the most constraining observables. Preliminary analysis disk dissipates on a timescale comparable to the satellite by Lainey et al. (2011) indicates that the Q factor may formation timescale (Canup and Ward, 2002; Alibert et al., be lower than 2000. Future observations will be needed to 2005a). In the case of Saturn, only one big moon has sur- conﬁrm this unexpectedly low dissipation factor, which still vived this migration process. remains puzzling.
Nevertheless, several large moons may have existed orig- inally and been lost before the disk dissipates (Canup and 2.3.4 Accretion of Titan Ward, 2006; Sasaki et al., 2010). The diﬀerence of satellite architecture between Jupiter and Saturn may thus result There is some analogy between the formation of the ter- from diﬀerences in disk evolution and satellite migration restrial planets and that of the icy moons. As described (e.g. Sasaki et al., 2010). By considering two diﬀerent disk above, icy moons probably accreted within the ﬁrst mil- evolution/structure models for Jupiter and Saturn, Sasaki lion years of the Solar System history in circumplanetary et al. (2010) proposed that in the case of Jupiter, the satel- disks from aggregates of ice and rock (e.g. Canup and Ward, lite migration may have stopped due to quick truncation of 2002; Estrada et al., 2009). Among the major icy satellites, gas infall caused by Jupiter opening a gap in the solar neb- Ganymede and Callisto seem to exhibit a strong dichotomy ula, while the subdisk was still relatively hot. In contrast,in in their degree of diﬀerentiation. This has often been at- the case of Saturn, the gas infall would have never stopped tributed to diﬀerences in the accretion processes. Titan may abruptly, and rather slowly decayed. Most of the moons that represent an intermediate case. However, the presence of its formed collided with the planet, and only one last big moon massive atmosphere suggests that Titan formed from diﬀer- would have survived. This surviving satellite would have ent materials than did its Galilean cousins, which may have formed on long timescales ( Myrs), owing to reduced mass signiﬁcant consequences for the ﬁnal assemblage. ∼ infall. This is obviously only one scenario among others, and several alternative scenarios may be envisioned. 22.214.171.124 Assemblage sequence: Role of impactor distribution The accretion by Saturn of early formed satellites with sizes analogous to that of Titan may also provide an expla- Anumberofmodelshavebeendevelopedforaccretional nation for the formation of the ring and the inner regular heating of icy moons (e.g. Schubert et al., 1981; Squyres moons. Numerical simulations performed by Canup (2010) et al., 1988; Coradini et al., 1989; Barr et al., 2010)). Al- suggest that the tidal removal of mass from a diﬀerenti- though some ingredients change from one model to another, ated, Titan-sized satellite as it migrates inward towards all these models follow the same approach which was ini- Saturn may be the origin of the giant planet’s ring and tially proposed by Safronov (1978) and Kaula (1979) for inner moons. These simulations show that planetary tidal the accretion of the Earth. The main diﬀerences between forces preferentially strip material from the satellite’s outer the diﬀerent models concern the assumption for the impact icy layers, while its rocky core remains intact and is lost energy partitioning as well as for impactor size and velocity by collision with the planet. The result of these computa- distributions. The most recent study based on this approach tions is a pure ice ring much more massive than Saturn’s was developed by Barr et al. (2010). In their approach, they current rings. As the ring evolves, its mass decreases and assumed that during the satellite growth the impactors are icy moons are spawned from its outer edge with estimated small enough that the impact energy is deposited at shallow masses consistent with Saturn’s ice-rich moons interior to depth and is eﬃciently radiated away. In this condition, and including Tethys. they showed that if Titan accreted slowly from numerous cold and small embryos more than 5 Myrs after CAI forma- The latter scenario has been further investigated by tion, the melting of the outer layer may have been avoided Charnoz et al. (2011), who conﬁrmed that all the moons and the undiﬀerentiated proto-core may extend up to the interior to Titan’s orbit may have been formed by these surface. accretion processes, if Saturn has a suﬃciently low dissipa- The Origin and Evolution of Titan 23
induce signiﬁcant melting and vaporization (Artemieva and Lunine, 2003, 2005). As the volume of melted material scales with the impact velocity and the impactor volume (Tonks and Melosh, 1992), giant impact events occurring at the end of planetary accretion may generate oceans of melt (e.g. Sekine et al., 2011). Even for impact with mod- erate velocities, their cumulative eﬀect (especially for large impactors) may lead to progressive increase of temperature up to the melting point of water ice.
126.96.36.199 Energy budget during the early stage of Titan’s evolution During the early stages of Titan’s history, in addition to impact heating, three main sources of energy may have con- tributed to the internal thermal budget: radiogenic heating Figure 2.4 Assemblage sequence and consequences for thermal by short-lived radioactive isotopes (mainly 26Al), tidal heat- evolution of the growing planet ing associated with despinning and viscous heating due to ice-rock separation. In this section, we provide some simple In reality, the size and velocity distribution of the im- estimates of these sources of energy in order to evaluate pactors have probably varied during the accretion phase. their potential impact on the thermal budget of early Titan. Two reservoirs of bodies may have contributed to the satel- lite growth: bodies in orbit around the planet, formed or Impact heating results from the deposition of impactor captured within the circumplanetary disk, and bodies in kinetic energy during satellite accretion. It therefore de- orbit around the Sun, colliding directly with the growing pends on the velocity and the mass of the impactor as satellite (e.g. Squyres et al., 1988). The contribution of each well as the way kinetic energy is converted into heat. The of these two reservoirs probably varied as the circumplan- impact creates a shock wave that compresses the satellite etary disk evolved. When the accretion sequence started, beneath the impact site. Below the impact site, the peak centimeter-sized planetocentric particles were probably pre- pressure is almost uniform in a quasi-spherical region, called dominant, while kilometer-sized and larger bodies became the isobaric core (e.g. Croft, 1982; Senshu et al., 2002). As more and more frequent during the late-stage of accretion shock compression is an irreversible process, the entropy (see Figure 2.4). below the impact site increases, leading to a temperature increase. Most of existing shock simulations have been per- −1 The temporal evolution of the size of impactors played an formed at high impact velocity (> 10 km.s ,Pierazzo important role in the diﬀerentiation of icy satellites since et al. (1997); Barr and Canup (2010)), so that they cannot it determined the accretion rate and the volume of the be used directly to constrain the shock heating in the low heated region after an impact. The increase of temperature velocity regime expected during the satellite growth. The upon impact in the outer region of the satellite is mainly heating and associated temperature increase can however controlled by the eﬃciency of post-impact cooling. The be estimated from energy balance considerations. The tem- cooling is determined by radiation at the surface, subsur- perature increase is determined from the peak pressure face thermal conduction as well as local heat advection within the isobaric core, which is directly related to the associated with melting, excavation and impact-induced impactor velocity, and from the intrinsic properties of the convection (e.g. Senshu et al., 2002). If the impact rate is target materials (density, speciﬁc heat, thermal expansion, too high or if the impact energy is buried at depths larger etc.) (e.g. Senshu et al., 2002). than a few kilometers, radiation and conduction are not eﬃcient enough to cool the subsurface and the temperature The impact velocity vimp is a contribution of the inﬁnite 2 2 velocity v∞ √e + i with e the eccentricity and i the increases. ∝ inclination of the planetary embryos and of the escape ve- 2 As noted above, the averaged size of impactors is expected locity vesc = 8/3πGρR of the growing icy satellite with R its radius, ρ its density and G the gravitational constant. to increase as a function of time during the accretion, and p 2 2 2 large impactors will be more and more frequent, possibly vimp = vesc + v∞ (2.5) leading to giant impacts during the late-stage of accretion The contribution of inﬁnite and escape velocities to the (Sekine and Genda, 2011). This has major consequences impact velocity can be expressed under the form of the grav- for the eﬃciency of cooling of the satellite and the pos- itational focusing F (Kaula, 1979; Schubert et al., 1981; sible occurrence of melting during accretion. Heliocentric g Barr et al., 2010): impactors with high velocities may also become more and 2 more frequent (Agnor et al., 1999) resulting in larger in- vesc Fg =1+ (2.6) creases of temperature. High velocity impacts may locally v „ ∞ « 24 Tobie, Lunine, Monteux, Mousis & Nimmo
During the early stages of planetary formation, dynamical comes signiﬁcant for layers above a radius of about 1000 km. friction reduces eccentricities and inclinations of growing bodies. Hence, early impacts on Titan probably occurred An additional source of heating corresponds to dissipation near the two-body escape velocity (i.e. v∞ was negligible) due to enhanced tidal friction before the satellite reaches (Kokubo and Ida, 1996; Agnor et al., 1999). synchronous rotation. It is indeed likely that Titan has an initial spin period much shorter that its present-day pe- In this case and assuming that no signiﬁcant melting oc- riod of 15.95 days. Initial spin periods as low as 6-10 hours curs when the icy satellite is small, energy balance between are possible (Lissauer and Safronov, 1991; Robuchon et al., kinetic energy and impact heating leads to a local tempera- 2010). In this condition, the tidal bulge raised by Saturn ture increase ∆T0 that is proportional to the square of the rotates relative to the satellite body frame. Owing to the radius of the growing icy moon (Monteux et al., 2007): anelastic properties of the interior, the fast-rotating satel- lite does not respond perfectly to the tidal forcing, resulting 2 4π γ ρGR in a misalignment of the tidal bulge with respect to Saturn’s ∆T0 = (2.7) 9 f(m) Cp direction. This acts as a torque that tends to slowdown the satellite spin, until the satellite is locked in spin-orbit reso- where γ is the percentage of kinetic energy retained as nance. The gradual decrease of the spin rate ω as a function heat deep below the surface, f(m)isthevolumeeﬀectively of time t is determined from the tidal torque acting on the heated over the volume of the isobaric core (typically equal satellite and is given by: to 2.5-3, Pierazzo et al. (1997); Monteux et al. (2007)) and Cp is the averaged heat capacity (typically varying between dω 3 k GM 2R5 1000-1500 J.kg−1.K−1 for ice-rock mixtures). As long as the = 2 S , (2.9) dt 2 6 radius of the growing satellite is below 1000 km, the local DST QC temperature increase is less than 10-15 K after each impact. with k2 tidal Love number, MS Saturn’s mass, DST mean Above 2000 km, the increase of temperature is larger than distance between Titan and Saturn, Q dissipation factor of Titan and C its polar moment of inertia (C 0.4 M R2 50 K. & × T at the time of accretion). As the despinning rate scales with 5 Radiogenic heating by short-lived radioisotopes (26Al, R ,itrapidlyincreaseswhenthesatellitereachesitsﬁnal 60 53 size. For k 0.07, the expected value for an homogeneous Fe, Mn) may also signiﬁcantly contribute during the ac- 2 & cretion period. For simplicity, we consider here only 26Al, rigid interior (Sohl et al., 1995) and Q=100-500 and initial which is the principal contributor (e.g. Robuchon et al., rotation period of 10 hours, the time needed to reach the 2010; Barr et al., 2010). The increase of temperature due 16-days period range between 2 and 10 Myrs, suggesting 26 to the decay of Al at a radius r depends on the time, tf , that despinning is achieved shortly after the accretion. Note at which the layer at radius r has been formed and it is given that once large-scale melting occurs, a global liquid layer by (e.g. Barr et al., 2010): may form, thus increasing the Love number and the despin- ning rate, so that the current period is reached even faster. m q (0) τ t (r) The total amount of energy dissipated during the despinning ∆T (r)= r 26 26 exp ln 2 f , (2.8) 26 C ln 2 − τ stage is given by the diﬀerence of body rotation energy be- p „ 26 « tween the initial and ﬁnal periods and leads to the following where mr is the rock mass fraction in the accreting material temperature increase: (around 0.5 in the case of Titan), τ is the half-life time 26 2 2 of 26Al (τ =0.716 Myr) and q (0) is the initial 26Al 1 C(ω0 ωc ) 26 26 ∆Tdesp = − , (2.10) heating rate at the time of the calcium-aluminium-rich 2 CpMT inclusions (CAI’s). tf denotes the time after the forma- with ω0 and ωc the initial and current angular rotation tion of CAIs. The initial heating rate is determined by rates, respectively, and MT Titan’s mass. For initial periods 26 −1 the speciﬁc heat production of Al H26(0)=0.341 W.kg varying between 10 and 6 hours, the averaged temperature (Robuchon et al., 2010), the Al content in rocks, [Al] and increase would range between 25 and 75 K, respectively. the initial 26Al/27Al. Assuming a rock composition similar The localization of the dissipation is very sensitive to the to CI chondrites, [Al] = 0.865 wt% and initial 26Al/27Al internal structure (Tobie et al., 2005a). For a homogeneous of 5.85 10−5 (Thrane et al., 2006; Barr et al., 2010) interior, the dissipation will rather uniform, with a maxi- × leads to q (0) = 1.725 10−7 W.kg−1.Thisisabout20 mum of dissipation at about mid-radius. For a diﬀerentiated 26 × million times larger than the radiogenic heating rate due to structure, the dissipation will be mainly located in the outer the decay of long-lived isotopes expected at present within regions (liquid water and ice I) and would be more eﬃciently Titan. However, this huge heat production rapidly decays. transported to the surface. For an accretion time of 2 Myr after CAI formation, the temperature increase due to 26Al decay would be of the Ice-rock separation, which may be triggered already dur- order of 100 K, but it becomes less than 10 K for accre- ing the accretion phase, may also signiﬁcantly contribute to tion times above 5 Myrs. Radiogenic heating by short-lived the energy budget. The total increase of internal temper- radioisotopes may have been of importance for the ﬁrst ature due to ice-rock diﬀerentiation, ∆Tdiff ,canbeesti- materials to accrete, i.e. for the materials that formed the mated from the change of gravitational energy, ∆Eg,be- inner part of the satellite. By contrast, impact heating be- tween a homogeneous interior and a diﬀerentiated interior The Origin and Evolution of Titan 25 consisting of a rock core and a thick H2Omantle(e.g.Fried- the satellite growth. In this case, a very massive ocean in son and Stevenson, 1983): equilibrium with a very massive atmosphere (P>50 bars) would be generated. Rock blocks resulting from the melting 2 2 2 3 GMT ρr 5 ρi 5 of icy impactors would sediment at the base of the ocean ∆Eg 1 Xc + 1 Xc & 5 RT " − ρ¯ ρ¯ − forming a thick layer of unconsolidated rock (Kirk and „ « „ « “ ” (2.11) Stevenson, 1987; Lunine and Stevenson, 1987). The subse- 5 ρ ρ ρ quent evolution of these diﬀerent possible initial structure + i r i X3 1 X2 2 ρ¯ ρ¯ − ρ¯ c − c is discussed in Section 2.4.1. „ « “ ”– with ρr the density of rock, ρi the averaged den- sity of ices and liquid water, Xc the ratio between the rock core radius, Rc and Titan’s radius, RT .For −3 −3 ρr =2700kg.m and ρi =1250kg.m , Xc =0.75 and ∆T =∆Eg/cpM 100 150 K (for cp =1000 1500 diff & − − kg.m−3). If the diﬀerentiation process is slow (> 1Gyr), the convective heat transfer should be able to transport this additional energy. If the ice-rock separation is fast enough (< 0.5Gyr),thedissipationofpotentialenergymay induce runaway melting and thus may create a catastrophic 2.3.5 Summary diﬀerentiation (Friedson and Stevenson, 1983; Kirk and At the start of this Section we posed three questions; we Stevenson, 1987). In particular, giant impact events occur- will conclude by providing some provisional answers. ring at the end of planetary accretion (e.g. Sekine et al., 2011) may generate locally large volume of melts, leading 1. What was the composition of the primordial building to the formation of rock diapir migrating toward the center blocks that built Titan and the Saturnian system ? Most (Figure 2.4). The viscous heating associated to the migra- of the solids that formed the Saturnian system probably tion of large rock diapirs may generate signiﬁcant melting originate from the surrounding solar nebula. They are in the surrounding ice, potentially triggering diﬀerentiation predicted to be composed of rock and ice and to contain runaway in a way similar to what is expected in terrestrial signiﬁcant amounts of volatiles (NH3,CH4,CO2,H2S) planets (e.g Monteux et al., 2009; Ricard et al., 2009). trapped in the form of hydrate and pure condensate. The apparent deﬁcit in primordial CO, N2,ArinTitan relative to proto-solar composition may be explained by partial devolatilization of building blocks during their 188.8.131.52 Possible post-accretional interior structure migration in Saturn’s subnebula. Depending on the conditions under which Titan accreted, diﬀerent post-accretional interior structures are possible 2. What make the Jupiter and Saturn systems so diﬀerent ? just after accretion as illustrated in Figure 2.5. If the ac- The diﬀerence in satellite architectures may be explained cretion is very slow ( > 1Myr)andtheimpactorsremain by diﬀerent evolutions of the circumplanetary disk. In small all along the accretion period (< 100 m), the surface the case of Jupiter, the feeding of the disk may have should able to cool down between successive impacts, and stopped abruptly, while in the case of Saturn, the satel- diﬀerentiation may be avoided. A fully undiﬀerentiated lite accretion in the disk may have lasted longer owing structure also requires that the accretion starts about 3-4 to reduced mass infall. During the accretion process, Myr after the CAI formation, in order to limit the increase several Titan-size moons may have been formed and lost of internal temperature due to short-lived radiogenic ele- by migration into proto-Saturn. Titan would be the only ments, and that the initial spin period of Titan was larger surviving big moon of Saturn, while Jupiter was able to than 8-10 hours. In such a cold accretion case, no signiﬁcant retain four big ones. atmosphere would be generated, and the formation of a sig- niﬁcant atmosphere should occur later during the satellite 3. Can Titan accrete undiﬀerentiated ? evolution (for instance during the LHB event). It seems extremely diﬃcult to avoid ice melting during the accretion phase, owing to combined eﬀects of impact If the growing satellite is hit by larger impactors at the heating, radiogenic heating by short-lived isotopes and end of accretion, a surface water ocean may be generated tidal despinning. Melting may be prevented only if Titan (R>2300 2400 km) (Figure 2.5, middle). Even if only accreted slowly (> 1Myr)fromaswarmofsmallbuilding − 10-20% of the total volume of ice melts, release of gases blocks (< 100 m). Even for faster accretion with larger (mainly CO2 and CH4,Tobieetal.(2011))trappedinthe impactors, large-scale melting is only reached toward the ice phase may generate an atmosphere up to 5-10 bars. Even end of the accretional growth, so that the inner part of if a global ocean is generated, most of the mass of Titan the interior structure ( < 1500 km) remains undiﬀeren- would still be undiﬀerentiated. Finally, if the accretion is tiated. In any case, a totally undiﬀerentiated structure faster and/or if the impactors are large even during the after accretion is not needed to explain the present-day earliest stage of accretion, ice melting rapidly occur during moderately diﬀerentiated structure. 26 Tobie, Lunine, Monteux, Mousis & Nimmo
Figure 2.5 Possible evolution scenarios for the interior of Titan for diﬀerent initial states. Depending mostly on the eﬃciency of heat transfer in the interior, diﬀerent bifurcations in the evolutionnary path may have occurred. The Origin and Evolution of Titan 27
2.4 COUPLED EVOLUTION OF THE ever, limited diﬀerentiation requires not only slow accretion ATMOSPHERE, SURFACE AND INTERIOR but also implies that the convective heat transfer was ef- OF TITAN ﬁcient enough to remove heat from the interior (Friedson and Stevenson, 1983; Mueller and McKinnon, 1988; Nagel During and after the accretion, exchanges between the et al., 2004). This is possible if the viscosity of the ice-rock interior, surface and atmosphere are likely to have oc- mixtures is suﬃciently low. Nagel et al. (2004) showed that, curred. The icy shell has probably played a crucial role all in the case of Callisto, if the average size of rocks was of along Titan’s evolution by storing and transferring volatile the order of meters to tens of meters, the interior may compounds from the internal ocean to the surface. These have experienced a gradual, but still incomplete unmixing exchanges have probably inﬂuenced the composition of the of the two components. Signiﬁcant lateral heterogeneities atmosphere during most of Titan’s evolution, and they are among the ice/rock mixtures existing after accretion could maybe still operating at present. In this section, we dis- also have enhanced the diﬀerentiation processes (O’Rourke cuss the coupled evolution between the diﬀerent envelops and Stevenson, 2011). In these conditions, convective mo- of Titan, and try to address the three questions listed below. tions associated with ice-rock unmixing processes may be eﬃcient enough to avoid melting. Following this scenario, 1. To what extent did the interior of Titan inﬂuence the the interior of Titan would be relatively similar to Callisto, evolution of its atmosphere ? Evidence for exchange with but with a thinner layer of ice-rock mixed materials, as the interior is provided by the detection of a signiﬁcant suggested by its smaller MoI factor. amount of 40Ar in the atmosphere. However, the degree and timing of internal outgassing still remain uncon- Nevertherless, uncompleted ice-rock separation is not strained. In this section, we review the diﬀerent internal necessarily needed, as the MoI factor may also be consistent processes that may have occurred all along Titan’s evo- with a highly hydrated silicate interior, possibly dominated lution and discuss their implications for the evolution of by antigorite serpentine (Castillo-Rogez and Lunine, 2010; the atmosphere. Fortes, 2011). The formation of a highly hydrated core may result from a rapid diﬀerentiation process, implying large-scale melting and eﬃcient water-rock interactions (e.g. 2. Did Titan always possess a massive N2-CH4 atmosphere ? The composition and the mass of Titan’s atmosphere Fortes et al., 2007; Castillo-Rogez and Lunine, 2010; Fortes, have probably varied since the satellite accretion. De- 2011). The combined eﬀect of accretional heating, radio- pending on the initial state acquired during the accre- genic decay by short-lived isotopes, tidal heating associated tion and on the eﬃciency of internal outgassing, we to despinning, and viscous heating associated with negative will discuss the diﬀerent evolution scenarios that can be rock diapirs may increase the internal temperature above envisioned. the melting point of ice. The presence of freezing point depressing agents such as ammonia or methanol (Lunine 3. Where is the missing hydrocarbon inventory ? Hydrocar- and Stevenson, 1987; Grasset and Sotin, 1996; Grasset and bon lakes has been identiﬁed at the surface of Titan (Sto- Pargamin, 2005; Deschamps et al., 2010) would further fan et al., 2007). However, the estimated volume is much favor the generation of melts as they strongly decrease smaller than what is expected from continuous photolysis the melting point of ice. The aqueous melt, which is less of methane since the formation of Titan (Lorenz et al., dense than ice-rock mixtures, at least at pressure above 2008a). This suggests that either a large fraction of hy- 100 MPa, will percolate upward, thus interacting with drocarbon products has been removed from the surface or surrounding rock particles. During their ascent, aqueous that the photochemical production has not been constant solutions in proto-Titan may have been able to interact through time. We will discuss the implications of these chemically with silicate minerals (Engel and Lunine, 1994) two diﬀerent hypotheses for the long-term evolution of and with other substances in solution, such as sulfates Titan. (Fortes et al., 2007). Interactions of aqueous ﬂuids with initially anhydrous silicate minerals at low temperature may result in the formation of hydrated silicate minerals, 2.4.1 Diﬀerentiation and long-term evolution of called serpentine (Fortes et al., 2007; Castillo-Rogez and the interior Lunine, 2010; Fortes, 2011). The hydrated minerals may remain stable as long as the rocky core temperature re- 184.108.40.206 Ice-rock segregation and formation of a rocky mains below approximately 900 K (Grindrod et al., 2008; core Castillo-Rogez and Lunine, 2010). Above this temperature, As already discussed in Section 2, the moment of inertia dehydration would be expected, releasing warm supercrit- ( C/M R2 =0.33 0.34) inferred from Cassini gravity ical water ﬂuids which may interact again with overlying T T − measurements suggest that Titan’s diﬀerentiation was not rocks during their ascent (Castillo-Rogez and Lunine, 2010). complete. Titan probably does not have an iron core, and a layer consisting of a rock-ice mixture may still exist between As pointed out by Iess et al. (2010), the main challenge arockycoreandahigh-pressureicelayer.Thispossible with this hypothesis is to keep most of the core hydrated uncompleted ice-rock separation has been proposed to be (i.e. < 900 K until present). Thermal convection is unlikely the result of cold accretion process (Barr et al., 2010). How- to start in a cold core dominated by hydrated silicates 28 Tobie, Lunine, Monteux, Mousis & Nimmo
(Castillo-Rogez and Lunine, 2010). As thermal diﬀusion is Section 2.3.1, non-water ice components such as CH3OH, not eﬃcient enough to extract heat, the core temperature NH3,CO2,H2Sarelikelytohavebeenincorporatedinside should rise up to the dehydration temperature. The temper- Titan. In this condition, eutectic solutions at temperatures ature increase can however be limited if the core is depleted well below 273 K would have formed (e.g. Sohl et al., 2010). in radiogenic elements. Castillo-Rogez and Lunine (2010) These highly concentrated eutectic solutions should facili- showed that if at least 30% of potassium are leached into the tate dissolution of mineral phases and chemical reactions. ocean from the silicate minerals during the earlier hydration In the case of ammonia-rich ﬂuids, Engel and Lunine (1994) + event, the internal temperature would slowly increase and indicated that NH4 -ions contained in the aqueous solu- the outer part of the core would remain hydrated. As shown tion may eﬃciently replace ions such as K+,Na+,Rb+ by Castillo-Rogez and Lunine (2010), a rock core consisting and Ca2+ contained in the mineral phase. Ammonia-rich of only antigorite serpentine is not required to explain the liquids are also expected to react with magnesium sulfate, present-day moment of inertia. An inner anhydrous core contained in the chondritic phase and with carbon dioxide with radius up to 1300 km overlaid by a hydrated mantle is brought in the ice phase to form ammonium sulfate and compatible with the MoI factor of 0.33-0.34 (see Figure 2.5 ammonium carbonate, respectively (Kargel, 1992; Fortes bottom right). The scenario of a fully hydrated core initially et al., 2007). investigated by (Fortes et al., 2007) must be considered as averyextremecase,maybenotfullyrealistic.Nevertheless, During the migration of the enriched ﬂuids toward the the idea that a large fraction of the core remain hydrated surface, their composition should vary as new elements are during most of Titan’s evolution should be considered as extracted from primary minerals and other precipitate in apossiblescenario.Morecompletechemicalandthermal secondary phases (e.g. Zolotov, 2010). In turn, the circula- evolution models are probably needed in the future to test tion of aqueous solutions leads to oxydation and hydration in more detail this hypothesis. reactions, analogous to aqueous processes on chondrite par- ent bodies (e.g. Sohl et al., 2010). These reactions are well As illustrated on Figure 2.5, diﬀerent evolutionnary paths documented at low pressure conditions and have been used are possible, some of them combining both solid-state un- for instance to constrain the possible composition of Eu- mixing and ice melting processes. For instance, for an ropa and Enceladus’ ocean (e.g. Zolotov, 2007; Zolotov and initially warm and gravitationally unstable interior (Fig. Kargel, 2009), but they are still poorly known at the high 2.5, bottom), the rock core may form rapidly owing to run- pressure range of Titan’s proto-core (> 1GPa).Onlyone away melting of the undiﬀerentiated proto-core, resulting in study by Scott et al. (2002) tried to investigate the aqueous acatastrophiccoreoverturn(KirkandStevenson,1987),or alteration of chondritic materials at high pressure. However, more slowly with progressive ice-rock separation and limited Scott et al. (2002) did not take into account the probable ice melting (e.g. Estrada and Mosqueira, 2011). By con- presence of ammonia, methanol or other compounds, which trast, an initially cold undiﬀerentiated interior may slowly may signiﬁcantly aﬀect the chemical equilibrium. Future diﬀerentiate with no melting and no formation of ocean experimental and thermodynamical investigations will be during the entire evolution. Any intermediate evolution is needed to better understand the complex chemistry that also possible. probably occurred during the diﬀerentiation.
After the complete separation of rock and ice, water-rock interactions are limited to the ice-rock interface at the base 220.127.116.11 Water-rock interactions in Titan’s conditions of the high-pressure layer. Even though water is mostly in On Titan, the crystallization of the internal liquid water the form of high-pressure ice phase at this interface owing results in the formation of a low-pressure ice layer at its to the elevated pressure (> 0.8 1GPa),liquidwatermay − top and of a high-pressure ice layer at its base. On smaller locally exist and potentially circulate in the upper part of icy moons, such a high-pressure ice layer does not exist due the rock core. In order to eﬃciently transfer the heat com- to reduced pressure (high-pressure phases only exist above ing from the rocky core through the high-pressure ice layer, about 200 MPa). On Europa and Enceladus for instance, the temperature proﬁle in the high-pressure ice is probably the internal water ocean is permanently in contact with close to the melting curve of water ice (e.g. Sotin and Tobie, the rocky core, thus favoring water-rock interactions all 2004). In this context, the formation of a transient liquid along the satellite evolution. By contrast, on Titan, a thick reservoir at the rock-ice interface is likely. These liquidsmay layer of high-pressure ice rapidly forms during the interior potentially interact with the underlying rocks before perco- evolution (Figure 2.5). The water ocean is expected to be lating upward to the ocean. Warm liquid water may also be in direct contact with underlying rocks only during a brief released from the core when the dehydration temperature is period of time (< 100 Myr) just after accretion, and only in reached (Castillo-Rogez and Lunine, 2010). This hydration the case of relatively warm accretion (Figure 2.5). event may occur relatively late during Titan’s evolution (> 3 3.5Gyrafterformation),providedsuﬃcientamounts − However, even if there is no direct contact between the of radiogenic potassium were leached from the core and that water ocean and the rocky core, water-rock interactions thermal diﬀusivity in the core is not too low (Castillo-Rogez may have occurred at diﬀerent periods during Titan’s evo- and Lunine, 2010) (see Section 18.104.22.168). In summary, in lution, especially during the diﬀerentiation. As discussedin warm evolution scenarios, intensive water-rock interactions The Origin and Evolution of Titan 29 were likely during two main sequences of Titan’s evolution: surface of the rocky core (1 ? 1.5 GPa) (McKinnon, 2010). before 1 Gyr during the diﬀerentiation and after 3-3.5 Gyr during core dehydration. Moreover, reduced water-rock in- On Titan, the extraction of 40Ar is more likely the result teractions may also exist all along Titan’s evolution owing of a three step mechanism. 40Kmayhavebeenextracted to circulation of high-pressure ice melts at the rock-ice ﬁrst from the silicate minerals by water interactions (e.g. interface. Engel and Lunine, 1994). Then dissolved 40K+ ions in the ocean would have decayed into 40Ar. Finally, the 40Ar It has been proposed that water-rock interactions may contained in the ocean would be extracted through direct lead to eﬃcient production of methane from carbon diox- outgassing of the ocean or progressive incorporation in ide, carbon grains or organic matters by serpentinization the outer shell in the form of clathrate hydrate and then (Atreya et al., 2006). However, this mode of production destabilization near the surface by cryovolcanism (e.g. Tobie seems incompatible with the D/H ratio observed at present et al., 2006) or impact (e.g. Barr et al., 2010). Assuming in the atmospheric methane, which is about three times that 30% of potassium has been leached into the ocean greater than the value expected in hydrothermally pro- implies that about 25-30% of the produced 40Ar has been duced methane (Mousis et al., 2009a). It has also been outgassed through the outer icy shell, which seems a rea- suggested that the oxidation of primordial NH3 to N2 may sonable outgassing eﬃciency. also occur in hydrothermal systems on early Titan (Glein et al., 2009), as well as on early Enceladus (Matson et al., The presence of signiﬁcant amount of 22Ne in the at- 2007; Glein et al., 2008). The endogenic production of N2 mosphere, if conﬁrmed, may also be indicative of water- from NH3 is possible only if the hydrothermal systems rock interactions and eﬃcient outgassing processes (Tobie were suﬃciently oxidized with C existing predominately in et al., 2011). Indeed, a chondritic core may contain 22Ne the form of CO2,notCH4.Italsoimpliesthatprimordial in suﬃcient amounts to explain the tentatively observed 15 NH3 was enriched in Ninordertoexplaintheelevated abundance in the atmosphere (Niemann et al., 2010). As 15 14 N/ NobservedinatmosphericN2.Althoughwater-rock further discussed in Section 22.214.171.124, neon is the least sta- interactions and associated chemical reactions are thermally ble gas compound in the form of clathrate and the less plausible, there are essentially no constraints available to soluble in liquid water, and therefore it is expected to be assess their likelihood. For instance, constraints on the pos- very eﬃciently extracted during methane-driven outgassing 15 14 sible primordial value of N/ NinNH3 or D/H in water events. The accurate determination of the abundance of the are still missing. Future measurements in comets by Rosetta diﬀerent isotopes of Ne (20Ne, 21Ne and 22Ne) by a future for instance will bring new bear on these ideas. mission will permit conﬁrmation of the tentative detection of neon by Huygens,anditwillprovidepertinenttestson Circumstantial evidence for water-rock interactions is, the origin of neon and on its extraction mechanism from however, provided by the detection of 40Ar in the atmo- the interior. sphere (Niemann et al., 2005; Waite et al., 2005; Niemann et al., 2010). 40Ar is one of the decay products of 40Kini- tially contained in silicate minerals (e.g. Engel and Lunine, 126.96.36.199 Volatile storage in the H Omantle 1994). The mass of 40Ar in the atmosphere is estimated 2 between 4.0 and 4.6 1014 kg (Niemann et al., 2010), which AmajorissueonTitanconcernsthestorageandtransport × corresponds to about 7.5-9% of the total mass that may of gas compounds in the H2Omantle,whichconstitute be contained in the rocky core, assuming a CI chondritic about 60-70% of the satellite volume. The amount of gas composition (Tobie et al., 2011). As a comparison, on Earth, compounds that can be stored in the ocean and in the ice Mars and Venus, the fraction of 40Ar outgassed from the layers is determined by their solubility in water and their silicate mantle is estimated to about 50% on the Earth stability relative to the clathrate phase. NH3 and CH3OH, (All`egre et al., 1996) and about 25% on Mars and Venus for instance, are highly soluble in liquid water and do not (Hutchins and Jakosky, 1996; Kaula, 1999). The outgassing form clathrate structures. At low temperature (T < 177 K), process is probably very diﬀerent on Titan. On Earth and ammonia combines with water molecules to form ammonia terrestrial planets, 40Ar is progressively outgassed from the hydrate (not clathrate) and methanol solidiﬁes at relatively silicate mantle through volcanism (e.g. Xie and Tackley, similar temperatures. Such low temperature conditions can 2004). 40Kdecaysinto40Ar in the silicate minerals and be reached only near the cold surface, and both ammonia the latter is released when the minerals are brought near hydrates and methanol ice become unstable below a depth the surface by convective motions and melt. On Titan, as of 5-10 km (depending on the subsurface thermal gradient). suggested by the gravity data, Titan’s rocky core is prob- As a consequence, ammonia and methanol should be nearly ably relatively cold during most of its evolution (< 1000 exclusively stored dissolved in the internal ocean. K, Castillo-Rogez and Lunine (2010)), therefore it has probably never reached magmatic temperatures (at least Although the solubility of most of the other gases, for on a global scale). Moreover, even if silicate melting has instance methane, in liquid water is low, large quantities of occurred, outgassing of erupting silicate magmas should be gas can be stored dissolved in the liquid water as the ocean limited due to the high pressure conditions expected at the is very massive on Titan and it is conﬁned at relatively elevated pressures (> 25 MPa). For example, more than 30 Tobie, Lunine, Monteux, Mousis & Nimmo
late ocean crystallization may result in the accumulation of methane-rich clathrate hydrates in the outer shell. However, even in these conditions, the total thickness of methane-rich clathrates incorporated in the outer ice shell should not exceed 5-10 km, corresponding to about 20% of the total mass of methane that is potentially contained in the interior (Fig. 2.6, Tobie et al. (2011)). A fraction of CO2,COand Ar should also be incorporated in the outer shell in the form of clathrate, but to a lesser extent than methane due to diﬀerences in solubility and clathrate stability.
188.8.131.52 Volatile release from the H2Omantle Even though the interior may potentially contain large quantities of volatiles, their release into the atmosphere is limited owing to their solubility in the internal ocean and owing to clathration processes within the H2Omantle.Out- gassing from the interior may occur through two diﬀerent Figure 2.6 Possible present-day structure of Titan’s interior mechanisms: (1) destabilization of clathrates stored in the and estimated volatile distribution within the thick H2O outer ice shell by thermal convective plumes or by impacts, mantle. The radius, temperature and pressure of each interface (2) mobilization of dissolved gas molecules contained in the are only indicative. water ocean and subsequent bubble formation if the ocean gets saturated (Tobie et al., 2006, 2009). Both processes are controlled by methane, since it is the dominant gas in the 2000 times the present-day mass of atmospheric methane form of clathrate hydrate in the outer shell and since it is may be stored dissolved in the ocean (Tobie et al., 2011). the only gas compound that may reach the saturation point in the ocean during periods of reduced ice shell thicknesses. Depending on their stability relative to the clathrate The other dominant compounds, CO2,NH3 or H2S, are phase, a fraction of the gases dissolved in liquid water can more soluble in liquid water, so that they will never reach however be incorporated in the form of clathrate hydrate saturation. For noble gases such Ne or Ar, although they (Choukroun et al., 2010; Tobie et al., 2011; Choukroun have a low solubility, their concentration is very low and et al., 2011). The most stable species expected in the ocean, always remain below the saturation point. However, when such as H2SandXe,shouldbepreferentiallyincorporated methane reaches the saturation point, Ne and Ar owing to in the form of clathrate hydrate, whereas the least stable their low solubility in water may be preferentially incorpo- species such as CO and Ne remain mainly in the liquid rated in CH4-rich bubbles formed at the ocean-ice interface phase (Tobie et al., 2011). The fraction incorporated in and transported upward with the percolating ﬂuids. Fol- the clathrate phase should also increase as the ocean cools lowing this extraction mechanism, the detection of 40Ar down and crystallizes. Clathrate hydrates which have a (and potentially 22Ne) suggests that the atmospheric mass composition dominated by H2Shaveadensitylargerthat of methane has been renewed at least 20-30 times during liquid water, so that they are expected to sink to the base of Titan’s evolution (Tobie et al., 2011). the ocean. As an example, clathrate structures containing 75% of H2S, 20% of CO2 and 5 % of CH4 have a density −3 of about 1050 kg.m .TheseH2S-rich sinking clathrates will also preferentially incorporate Xe owing to its very high 2.4.2 Formation and evolution of the atmosphere stability in clathrate phase. As a consequence, H SandXe 2 184.108.40.206 Origin of the massive N -CH atmosphere are expected to be mostly stored in the form of clathrate 2 4 hydrate within the high-pressure ice layer. The formation of Titan’s atmosphere is certainly an event which occurred relatively early in its history, otherwise we During the clathration process, the dissolved fraction of should see a number of impact craters too small to have H2Sprogressivelyreduces,andCH4 and CO2 eventually been formed under the dense atmosphere of today. However, become the dominant clathrate former. When CH4 becomes whether it is an accretionally generated atmosphere (e.g. the dominant clathrate former gas, the clathrate density Lunine and Stevenson, 1987; Kuramoto and Matsui, 1994) gets lower than the density of liquid water, and clathrate or a late veneer generated during a subsequent late heavy hydrates accumulate at the top of the ocean. Diﬀerentiation bombardment (Zahnle et al., 1992; Sekine et al., 2011) is an of Titan’s interior during the early stage of its evolution issue that cuts to the heart of the question of whether Ti- may also lead to massive release of methane-rich clathrate tan’s formation environment was chemically and thermally hydrates, and hence accumulation in the outer shell above distinct from that of Jupiter (see above). If the disk around the ocean (Tobie et al., 2006). Both early diﬀerentiation and Saturn was colder than that at Jupiter, rich in volatiles like The Origin and Evolution of Titan 31
13 12 methane and ammonia, then the atmosphere most likely was C /C in CH4 should be much larger than the observed primordial. The likely presence of an interior water ocean, value, which is close to the solar value (Mandt et al., 2009; outgassing through time of methane needed to replace that Niemann et al., 2010). This nearly unmodiﬁed ratio suggests destroyed by photolysis, and the elevated nitrogen isotopic that methane released during the early stage of evolution ratio in Titan’s atmosphere (Niemann et al., 2010) all hint should have been totally lost or rapidly re-trapped in the at the likelihood that Titan acquired its volatile inventory subsurface without being aﬀected by atmospheric fractiona- throughout accretion, the present atmosphere being an tion processes. Today’s methane has probably been injected outgassed remnant of a volatile-rich Titan. much later during the evolution, possibly due to the desta- bilization of crustal reservoir of CH4-rich clathrate (Tobie AvolatilerichTitanformedunderallbutthecoldest et al., 2006). The only record of the accretional time lies accretional scenarios would have had an initially warm in the nitrogen isotopic ratio, modestly heavier than the ammonia-rich atmosphere, if indeed the dominant carrier terrestrial value (Mandt et al., 2009) , suggesting a source of nitrogen was ammonia. This seems to be required by the distinct form the Earth, or loss of nitrogen to space, or 36 measurement of Ar to nitrogen (Niemann et al., 2010), both. However, a single isotope ratio is a poor constraint on which can be interpreted as requiring that molecular ni- the rate and volume of the loss because isotopic enrichment trogen was initially in the form of ammonia (Owen, 1982). is a function of the rate and mechanism of escape. Further Arguments to the contrary–that somehow the argon is be- extensive discussion of the early Titan atmosphere may be ing sequestered by a condensed or crustal reservoir–don’t found elsewhere (Lunine et al., 2010a). seem to make sense in light of the presence of radiogenic argon at two orders of magnitude larger mixing ratio. Thus we can assume that ammonia was in abundance on early 220.127.116.11 Interactions with the icy shell Titan, initially in the form of an atmosphere and surface ammonia-water ocean supported by accretional heating and All along Titan’s evolution, interactions with the solid icy astronggreenhouseeﬀect(Lunineetal.,2010a).Conversion shell has likely inﬂuenced the composition of the atmo- of ammonia to molecular nitrogen by photolysis (Atreya sphere. During the post-accretional cooling leading to the et al., 1978) and shock heating (Jones and Lewis, 1987; crystallization of the primordial water ocean, part of the at- McKay et al., 1988) all provide a plausible story in which mospheric gases, notably CO2 and CH4,maybere-trapped the ammonia is rapidly consumed, the atmosphere cools and ﬁrst by clathration and then by condensation (Choukroun condenses out, and nitrogen takes its place. Ishimaru et al. et al., 2010; Lunine et al., 2010a). Large quantities of CO2 (2011), however, demonstrated that eﬃcient conversion of and CH4 were probably incorporated in the primitive icy NH3 into N2 by impact shock requires on oxidizing CO2-rich shell. Additional CH4-rich clathrates released during the proto-atmosphere on Titan. This is consistent with CO2 diﬀerentiation of the deep interior may have also accumu- as the dominant gas compound in the building blocks that lated at the base of this primitive outer shell, leading to the formed Titan as predicted by Hersant et al. (2008); Mousis formation of a thick ( > 10 km) CH4-rich clathrate layer et al. (2009c); Tobie et al. (2011). (Tobie et al., 2006). This shell would then thin in response to the release of heat associated to the core evolution (Tobie Sekine et al. (2011) also proposed, based on laser-gun et al., 2006), resulting in the release of large amounts of experiments, that ammonia ice contained in the outer icy methane which would be mainly stored dissolved in the shell may have been eﬃciently converted to N2 by high- ocean, only the gas in excess may have been outgassed velocity impacts during the LHB. The generation of the (Tobie et al., 2009). N2-rich atmosphere by this process requires the presence of several percents of ammonia in the outer shell, which im- According to this scenario, several kilometers of CH4-rich plies that Titan’s outer layer has remained undiﬀerentiated clathrates may still be present in the outer icy shell at and unmelted until the LHB era. As already discussed, it present, and recent thermal destabilizations of this reservoir is very diﬃcult to avoid melting and partial diﬀerentiation may explain the present-day atmospheric methane (Tobie during the accretion, so that it is unlikely that Titan’s icy et al., 2006). CH4-rich clathrates are probably not the only shell contained ammonia in suﬃcient amounts at the time crustal reservoir. Large volumes of liquid hydrocarbons of the LHB. While the LHB event has probably aﬀected may also be potentially stored in a porous crust (Kossacki the evolution of the atmosphere, notably by inducing a and Lorenz, 1996; Voss et al., 2007). The morphology and signiﬁcant atmospheric erosion (Korycansky and Zahnle, distribution of the hydrocarbon lakes on Titan suggest that 2011; Ishimaru et al., 2011), it seems unlikely that it has they are connected to a subsurface reservoir (Hayes et al., been the main contributor for the atmosphere generation, a 2008) (see also Aharonson et al. this book). However, no massive atmosphere probably already existed prior to that constraints exist on the extension of this subsurface reser- event. voir as well as on the methane content. The lakes seems dominated by ethane (Mitri et al., 2007; Brown et al., 2008; Outgassing of methane during the accretion and possibly Cordier et al., 2009), which suggests that the methane frac- during the LHB would have been vigorous. However, the tion in a subsurface liquid reservoir is limited at present. present-day methane in the atmosphere is unlikely to be the Based on the available observations, it is impossible to remnant of this primordial outgassing. If it was the case, the constrain the contribution of this potential reservoir to 32 Tobie, Lunine, Monteux, Mousis & Nimmo the methane inventory. Additional observations, such as 18.104.22.168 Long-term climate stability of Titan’s subsurface sounding by a future exploration mission, will atmosphere. be needed to better constrain the nature of the subsurface Titan’s atmosphere has been unstable on timescales of be- reservoir. tween tens and hundreds of millions of years unless large areas of methane seas and oceans, an order of magnitude The mysterious absence or rarity of the primordial no- or more larger than the present-day lakes and seas, have ble gases in Titan’s atmosphere may also be indicative of existed on the surface throughout geologic time. In the interactions of the atmosphere with the solid icy shell. It absence of such sustaining bodies of liquid methane, or has been shown recently that the apparent deﬁcit in noble an implausibly ﬁnely tuned outgassing rate that keeps the gases could be essentially due to the formation of multiple atmospheric methane content within a factor of several of guest clathrates at the surface-atmosphere interface of the the current abundance for billions of years, the methane in satellite (Mousis et al., 2011). These clathrates would prefer- Titan’s atmosphere must have been periodically depleted. entially store some species, including all heavy noble gases, For the ﬁrst half of Titan’s history, when the Sun’s lumi- over others. These results have been obtained from the nosity was less than 80-85% of its current value, loss of use of a statistical-thermodynamic model based on recent methane may have caused the nitrogen to rain out onto determinations of intermolecular potentials and supersede Titan’s surface (McKay et al., 1993; Lorenz et al., 1997a), the previous ones that suggested a limited trapping of argon producing a global sea of liquid nitrogen under a thinner (Osegovic and Max, 2005; Thomas et al., 2008). In fact, the clear atmosphere. clean water ice needed for the formation of these clathrates may be delivered by successive episodes of cryovolcanic This scenario, bizarre in its contrast with the present-day deposits. In this scenario, the formation of clathrates in state, raises all sorts of issues: (1) Is there geologic evi- the porous lavas and their propensity for trapping Ar, Kr, dence for a long-standing ancient global sea of either liquid and Xe would progressively remove these species from the nitrogen or (to avoid this catastrophe) a standing ancient atmosphere of Titan over its history. A global clathrate global sea of methane and ethane? (2) How could a thicker shell with an average thickness not exceeding a few meters atmosphere be restored, even with massive outgassing of may be suﬃcient on Titan for a complete removal of the methane (why wouldn’t the methane simply freeze out at heavy noble gases from the atmosphere. Note, however, that the lower global temperatures for which liquid nitrogen is there has been no convincing and unambiguous evidence stable)? (3) What would prevent under a thin atmosphere for cryovolcanic activities on Titan so far (see for instance the formation of massive ice caps of nitrogen at cold poles, Moore and Pappalardo, 2011), so that the proposed trap- possibly rendering as in point (2) a permanently frozen ping mechanism still remains hypothetic. Titan? It is possible that freeze out did not occur even under a fainter young Sun, and indeed for values of the If clathrate formation is an eﬃcient process as hypothe- solar luminosity within 15-20% of the present-day, it seems sized for the noble gases, then it is possible that clathrates that a nitrogen atmosphere not much diﬀerent from that at are also a sink for the prodigious amounts of ethane, and present could have been sustained in the absence of methane lesser amounts of propane, expected to have been produced (Lorenz et al., 1997a). over the history of Titan by photochemical processing of methane. Episodes of crustal volcanism could encourage for- As crustal and deeper sources of methane continue to be mation of ethane clathrate (Mousis and Schmitt, 2008), depleted by photolysis, it seems inevitable that a state will as could impacts (Lunine et al., 2010b), or simple per- be reached in which Titan will lose its surface methane and colation through a heavily fractured crust might progres- revert to a nearly pure nitrogen atmosphere, largely clear of sively allow the substitution of ethane for methane in crustal aerosols, but with a surface temperature not much diﬀerent clathrate hydrate, and reconversion of water ice that once from today or even larger as the Sun continues to grow in was clathrate prior to losing methane by outgassing. Re- luminosity. Eventually, when the Sun becomes a red giant, gardless of the process, while methane clathrate hydrate Titan will have lost much of its atmosphere thanks to a is slightly less dense than the underlying water-ammonia strongly enhanced solar wind, but will for a time support a ocean, ethane clathrate is not. Hence any ethane clathrate surface water ocean until the red giant’s luminosity begins will have a tendency to move downward in a predominantly to fade (Lorenz et al., 1997b) . Perhaps any life extant in the methane clathrate shell, and once at the bottom will sink in deep water interior might then adapt to the surface, develop the ocean–to exist stably at the bottom. Timescales might aformofphotosynthesis,etc.,iftimeisnottoobrief.Such be an issue here, leaving most of the ethane in the outer shell, speculations would be in the realm of science ﬁction were it but sequestered away from the surface-atmosphere system. If not for the likelihood that Titan is representative of a class this view is correct, Titan’s outer shell is laden with ethane of planets that would stably exist within 1 AU of M to K clathrate from billions of years of methane photolysis–the dwarfs, with a spectrum of environments from the present- lakes and seas being only the most recently produced ethane. day state to an ancient nitrogen sea state, to the ultimate warmer version discussed here, simply by virtue of adjust- ing the stellar spectral type from late M to K (Lunine, 2010). The Origin and Evolution of Titan 33
2.4.3 Summary some of the remaining problems include: At the start of this Section we posed three questions; we 1. What was the composition of the building blocks that will conclude by providing some provisional answers. formed Titan ? Theorotical models of planetary forma- tion suggest that the building blocks that formed Titan 1. To what extent did the interior of Titan inﬂuence the contained a few percents of various gas compounds, evolution of its atmosphere ? The probable existence of mainly CO ,NH,CH,aswellassulfurcompounds, an internal ocean at present suggests that the diﬀer- 2 3 4 methanol and traces of noble gases. Most of the gas entiation of Titan’s interior has been accompanied by compounds initially expected have not been detected so signiﬁcant water-rock interactions and massive releases far. Xe and Kr, for instance, appear to be depleted rela- of gas compounds initially incorporated in clathrate tive to the predicted abundance. By contrast, Ne, which structure. Although most of gas species should be stored was not expected, has been tentatively detected. CO , dissolved in the massive internal ocean and in the form 2 which is predicted to be the dominant carbon-bearing of clathrate in both the outer ice layer and the high- species, has only been identiﬁed at Huygens’landingsite pressure ice layer, a signiﬁcant fraction of CH was likely 4 (Niemann et al., 2010) and tentatively in some other outgassed through time. speciﬁc areas (McCord et al., 2008). Constraining the CO fraction on Titan as well as the noble gas content 2. Did Titan always possess a massive N -CH atmosphere 2 2 4 will provide key constraints on the conditions under ? Titan likely acquired its atmosphere during its ﬁrst which the building blocks formed. Atmospheric sampling hundred millions of years, at the latest during the LHB. by a future exploration mission will permit us to deter- The atmospheric composition and mass have probably mine more accurately the noble gas content of Titan’s varied along Titan’s evolution. The early atmosphere atmosphere, in particular to conﬁrm the detection of may have been more massive and dominated by CO . 2 neon. Determination of noble gas abundance in Saturn’s Following the post-accretional cooling phase, nitrogen atmosphere by an entry probe will also be extremely probably rapidly became the dominant gas. It is still useful to determine the eﬃciency of noble gas trapping unclear how nitrogen has been produced from ammonia in planetesimals embedded in the feeding zone of Saturn. and how atmospheric escape has inﬂuenced the nitro- Moreover, in situ analysis in diﬀerent materials collected gen inventory. The methane abundance has probably at the surface will also provide pertinent information signiﬁcantly varied through time depending on the bal- on the possible CO content in the outer shell as well ance between replenishing and destruction. Long periods 2 as on the possible sequestration of noble gases in the without atmospheric methane may even be envisioned. subsurface. 3. Where is the missing hydrocarbon inventory ? Large vol- 2. How did Titan’s N -dominating atmosphere form and umes of hydrocarbon may be stored in the form of liquid 2 evolve ? It is likely that Titan’s atmosphere has been or clathrate hydrate in the solid icy shell. However, there acquired relatively early during the satellite evolution. are still no direct constraints on the volume of these However, it is still unclear how the N -dominating at- putative reservoirs. Alternatively, the lack of extensive 2 mosphere has been produced and how its composition hydrocarbon reservoirs may indicate that the production has evolved through time. The elevated 15N/14Nratio of hydrocarbons from methane photolysis has been less observed in Titan’s nitrogen is commonly attributed eﬃcient than previously anticipated. This may suggest to enhancement due to massive atmospheric escape. that methane has been present in the atmosphere during However, we are still missing constraints on the ini- limited periods of time. tial value in Titan. Determination of 15N/14Nratioin NH3 in comets, in particular by the Rosetta mission on Comet 67P/Churyumov-Gerasimenko, will provide the ﬁrst direct constraint. Such measurements together with determination of isotopic ratios in non-radiogenic 2.5 CONCLUSION AND QUESTIONS FOR isotopes of argon (38Ar/36Ar) and neon (22Ne/20Ne) by FUTURE MISSIONS in situ sampling will allow us to determine how much mass Titan’s atmosphere have lost since its formation, Since the arrival of the Cassini-Huygens mission to the thus providing constraints on the early state of Titan’s Saturn system, a large amount of data has been collected atmosphere. on Titan and various models have been developed to inter- pret them in an evolution context. In spite of considerable 3. How eﬃcient was the interaction between rock and H2O progresses in our understanding of Titan’s origin and evo- in the interior ? The moment of inertia inferred from lution, many questions remain, and they will probably not Cassini gravity data as well as the detection of 40Ar in be answered by the end of the Cassini-Huygens mission. A the atmosphere is indicative of prolonged interactions series of investigations in situ and from the orbit will be between silicate minerals and H2O(liquidwaterorwa- needed in the future to address all the mysteries of Titan ter ice). Such an interaction has key implications for that Cassini-Huygens just started to unveil. Speciﬁcally, the chemical evolution of Titan’s interior, in particular 34 Tobie, Lunine, Monteux, Mousis & Nimmo
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