Saturn's largest satellite, Titan, has a massive nitrogen atmosphere containing up to 5 per cent methane near its surface. Photochemistry in the stratosphere would remove the present-day atmospheric methane in a few tens of millions of years. Before the Cassini-Huygens mission arrived at Saturn, widespread liquid methane or mixed hydrocarbon seas hundreds of metres in thickness were proposed as reservoirs from which methane could be resupplied to the atmosphere over geologic time. Titan fly-by observations and ground-based observations rule out the presence of extensive bodies of liquid hydrocarbons at present, which means that methane must be derived from another source over Titan's history. Here we show that episodic outgassing of methane stored as clathrate hydrates within an icy shell above an ammonia-enriched water ocean is the most likely explanation for Titan's atmospheric methane. The other possible explanations all fail because they cannot explain the absence of surface liquid reservoirs and/or the low dissipative state of the interior. On the basis of our models, we predict that future fly-bys should reveal the existence of both a subsurface water ocean and a rocky core, and should detect more cryovolcanic edifices.

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Saturn's moon Enceladus harbours a global water ocean , which lies under an ice crust and above a rocky core . Through warm cracks in the crust a cryo-volcanic plume ejects ice grains and vapour into space that contain materials originating from the ocean. Hydrothermal activity is suspected to occur deep inside the porous core, powered by tidal dissipation . So far, only simple organic compounds with molecular masses mostly below 50 atomic mass units have been observed in plume material. Here we report observations of emitted ice grains containing concentrated and complex macromolecular organic material with molecular masses above 200 atomic mass units. The data constrain the macromolecular structure of organics detected in the ice grains and suggest the presence of a thin organic-rich film on top of the oceanic water table, where organic nucleation cores generated by the bursting of bubbles allow the probing of Enceladus' organic inventory in enhanced concentrations.

[1] The thickness of Europa's ice shell is constrained with numerical experiments of thermal convection, including heterogeneous tidal heating. Thermal convection occurs in the stagnant lid regime with most of the tidal heating located in the bottom half of the layer. The addition of tidal heating mainly results in the increase of the temperature of the well-mixed interior and in the decrease of the heat flux at the base of the ice layer. In many simulations, the ice in hot plumes is heated up to its melting point. This induces episodic upwellings (0.5 Ma) of partially molten ice up to the base of the conductive lid, with possible implications for the formation of lenticulae and chaos regions. The thickness of the convective ice shell in equilibrium with the heat flow from the silicate core is estimated to be about 20-25 km. Tidal dissipation and surface temperature variations create lateral variations of the ice shell thickness of about 5 km, with maxima near the equator at the Jovian and anti-Jovian points and minima at midlatitudes. Surface heat flux is about 35-40 mW.m À2; it is almost constant all over Europa's surface, even though the tidal dissipation rate is four times larger at the poles than at the equator.

TRAPPIST-1 planets are invaluable for the study of comparative planetary science outside our Solar System and possibly habitability. Both Time Transit Variations (TTV) of the planets and the compact, resonant architecture of the system suggest that TRAPPIST-1 planets could be endowed with various volatiles today. First, we derive from N-body simulations possible planetary evolution scenarios, and show that all the planets are likely in synchronous rotation. We then use a versatile 3-D Global Climate Model (GCM) to explore the possible climates of cool planets around cool stars, with a focus on the TRAPPIST-1 system. We look at the conditions required for cool planets to prevent possible volatile species to be lost permanently by surface condensation, irreversible burying or photochemical destruction. We also explore the resilience of the same volatiles (when in condensed phase) to a runaway greenhouse process. We find that background atmospheres made of N 2 , CO or O 2 are rather resistant to atmospheric collapse. However, even if TRAPPIST-1 planets were able to sustain a thick background atmosphere by surviving early X/UV radiation and stellar wind atmospheric erosion, it is difficult for them to accumulate significant greenhouse gases like CO 2 , CH 4 or NH 3 . CO 2 can easily condense on the permanent nightside, forming CO 2 ice glaciers that would flow toward the substellar region. A complete CO 2 ice surface cover is theoretically possible on TRAPPIST-1g and h only, but CO 2 ices should be gravitationally unstable and get buried beneath the water ice shell in geologically short timescales. Given TRAPPIST-1 planets large EUV irradiation (at least ∼ 10 3 × Titan's flux), CH 4 and NH 3 are photodissociated rapidly and are thus hard to accumulate in the atmosphere. Photochemical hazes could then sedimentate and form a surface layer of tholins that would progressively thicken over the age of the TRAPPIST-1 system. Regarding habitability, we confirm that few bars of CO 2 would suffice to warm the surface of TRAPPIST-1f and g above the melting point of water. We also show that TRAPPIST-1e is a remarkable candidate for surface habitability. If the planet is today synchronous and abundant in water, then it should always sustain surface liquid water at least in the substellar region, whatever the atmosphere considered.

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Tidal interactions between Saturn and its satellites play a crucial role in boththe orbital migration of the satellites and the heating of their interiors. Therefore constraining the tidal dissipation of Saturn (here the ratio k 2 /Q) opens the door to the past evolution of the whole system. If Saturn's tidal ratio can be determined at different frequencies, it may also be possible to constrain the giant planet's interior structure, which is still uncertain. Here, we try to determine Saturn's tidal ratio through its current effect on the orbits of the main moons, using astrometric data spanning more than a century. We find an intense tidal dissipation (k 2 /Q= (2.3 ± 0.7) × 10 -4 ), which is about ten times higher than the usual value estimated from theoretical arguments. As a consequence, eccentricity equilibrium for Enceladus can now account for the huge heat emitted from Enceladus' south pole. Moreover, the measured k 2 /Q is found to be poorly sensitive to the tidal frequency, on the short frequency interval considered. This suggests that Saturn's dissipation may not be controlled by turbulent friction in the fluid envelope as commonly believed. If correct, the large tidal expansion of the moon orbits due to this strong Saturnian dissipation would be inconsistent with the moon formations 4.5 Byr ago above the synchronous orbit in the Saturnian subnebulae. But it would be compatible with a new model of satellite formation in which the Saturnian satellites formed possibly over longer time scale at the outer edge of the main rings. In an attempt to take into account for possible significant torques exerted by the rings on Mimas, we fitted a constant rate da/dt on Mimas semi-major axis, also. We obtained an unexpected large acceleration related to a negative value of da/dt= -(15.7 ± 4.4) × 10 -15 au/day. Such acceleration is about an order of magnitude larger than the tidal deceleration rates observed for the other moons. If not coming from an astrometric artifact associated to the proximity of Saturn's halo, such orbital decay may have significant implications on the Saturn's rings.

Geophysical measurements can reveal the structures and thermal states of icy ocean worlds. The interior density, temperature, sound speed, and electrical conductivity thus characterize their habitability. We explore the variability and correlation of these parameters using 1‐D internal structure models. We invoke thermodynamic consistency using available thermodynamics of aqueous MgSO4, NaCl (as seawater), and NH3; pure water ice phases I, II, III, V, and VI; silicates; and any metallic core that may be present. Model results suggest, for Europa, that combinations of geophysical parameters might be used to distinguish an oxidized ocean dominated by MgSO4 from a more reduced ocean dominated by NaCl. In contrast with Jupiter's icy ocean moons, Titan and Enceladus have low‐density rocky interiors, with minimal or no metallic core. The low‐density rocky core of Enceladus may comprise hydrated minerals or anhydrous minerals with high porosity. Cassini gravity data for Titan indicate a high tidal potential Love number (k2>0.6), which requires a dense internal ocean (ρocean>1,200 kg m−3) and icy lithosphere thinner than 100 km. In that case, Titan may have little or no high‐pressure ice, or a surprisingly deep water‐rock interface more than 500 km below the surface, covered only by ice VI. Ganymede's water‐rock interface is the deepest among known ocean worlds, at around 800 km. Its ocean may contain multiple phases of high‐pressure ice, which will become buoyant if the ocean is sufficiently salty. Callisto's interior structure may be intermediate to those of Titan and Europa, with a water‐rock interface 250 km below the surface covered by ice V but not ice VI.