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Mimas has a high eccentricity and an anomalously high physical libration like its neighbor, Enceladus, but does not appear to have a geologically active surface. We investigate Mimas' interior with a technique that infers spatial variations in tidal heating from its global shape. To account for its hydrostatic shape, we find Mimas' normalized moment of inertia is 0.375 ± 0.0025, indicating a relatively undifferentiated world. Its remaining topography is consistent with a ∼30 km thick conductive ice shell in Airy isostasy atop a weakly convecting ∼30 km thick layer that itself mantles a ∼140 km radius ice‐rock interior. The convective shell's density must be closer to the interior density to satisfy our moment of inertia and provide a denser compensating layer for Airy isostasy. This ice‐rock interior is elongated along the Mimas‐Saturn axis, which can match Mimas' observed physical libration without appealing to an ocean. The inferred ice shell thickness variations indicate a high obliquity (≈1.7°). We suggest that the obliquity damped rapidly, after which topography froze in when internal heat was conducted out of Mimas quicker than isostatic ice shell thickness variations could relax. We speculate on several possible explanations for this transient high obliquity, including excitation by ring‐forming material following the recent tidal disruption of an eccentric satellite. We cannot rule out a young Mimantean ocean, but our inferred moment of inertia favors a Mimas that was solid when it experienced a period of high obliquity, did not significantly melt during a recent resonance with Enceladus, and is solid today. Plain Language Summary: While Saturn's moon Mimas appears to have an ancient surface covered in craters, it has a highly elliptical orbit, as well as a wobble in its rotation that is higher than expected for a solid, spherically symmetric moon. This wobble could indicate Mimas has either an ocean or an elongated core. To investigate the possibility of an ocean, we examine differences from sphericity in Mimas' shape. Some of this shape is from Mimas' tidal bulge and rotational flattening. Assuming the rest is topography due to differences in heat that arise due to tidal forces, we can relate Mimas' topography to its tidal heating pattern. Because a moon's tidal heating pattern is influenced by its interior structure, we can then make inferences on Mimas' interior. What we find indicates that the tidal heating conditions that produced Mimas' shape indicate its spin pole was once tilted much higher than one would expect at present day. This could have occurred long ago, after which the topography froze in when Mimas rapidly cooled. The tidal heating pattern we infer also indicates that Mimas did not have an ocean at the time of this strong tidal heating, and likely remains solid to this day. Key Points: We infer spatial variations in tidal heating and ice‐shell thickness from Mimas' non‐hydrostatic shapeOur preferred Mimas interior has no sub‐surface ocean but requires high obliquity tides in Mimas' pastMimas' past obliquity may have been excited by gravitational interaction with an inward‐migrating debris disk that formed Saturn's rings [ABSTRACT FROM AUTHOR]

In order to improve our understanding of the interior structure of Saturn's small moon Enceladus, we reanalyze radiometric tracking and onboard imaging data acquired by the Cassini spacecraft during close encounters with the moon. We compute the global shape, gravity field, and rotational parameters of Enceladus in a reference frame consistent with the International Astronomical Union's definition, where the center of the Salih crater is located at −5° East longitude. We recover a quadrupole gravity field with J3 and a forced libration amplitude of 0.091° ± 0.009° (3‐σ). We also compute a global shape model using a stereo‐photoclinometry technique with a global resolution of 500 m, although some local maps have higher resolutions ranging from 25 to 100 m. While our overall results are generally consistent with previous studies, we infer a thicker 27–33 km mean ice shell, a thinner 21–26 km mean ocean thickness, and a mean core density range of 2,270–2,330 kg/m3. Plain Language Summary: Geodetic data, such as shape, gravity, and rotation, provide important constraints for probing a planetary body's interior structure. We analyze radiometric tracking and onboard imaging data acquired during close encounters of Enceladus by the Cassini spacecraft to compute geodetic products including topographic and gravitational fields in a common reference frame. The recovered Enceladus topography has a global resolution of 500 m, with some local regions having 25–100 m resolution. Our study suggests that Enceladus has a 27–33 km mean ice shell thickness, a 21–26 km mean ocean thickness, and a mean core density range of 2,270–2,330 kg/m3. Key Points: A full quadrupole gravity field with J3 and the forced libration amplitude of 0.091° ± 0.009° are recoveredA 500‐m resolution global topography model was computed, with some local regions having 25−100 m resolutionThe results suggest that Enceladus has a 27–33 km mean ice shell thickness, a 21–26 km ocean thickness, and a mean core density range of 2,270–2,330 kg/m3 [ABSTRACT FROM AUTHOR]

We compare several recent Martian interior models and evaluate how these are impacted by the tidal constraints provided by the Love number k2 and the secular acceleration in longitude s of its main moon, Phobos. The expression of the latter is developed up to harmonic degree 5 to match the accuracy of the current observations. We match a number of current interior structure models to the recent measurements of the tidal parameters and derive estimations of the possible core radius, temperature profile, and attenuation in the Martian interior. Our estimation of the core radius is 1,820 ± 80 km, consistent with recent seismic measurements. The attenuation profiles in the Martian interior at the main tidal period of Phobos are similar between the considered models, giving a range for the degree‐2 bulk tidal attenuation Q2 = 93.0 ± 8.40 but diverge at seismic frequencies. At seismic frequencies, model shear attenuation Qμ ranges between 100 and 4,000 in the lower mantle, so that a measurement of seismic shear attenuation could be used as an effective means for distinguishing between the models considered. Other constraints such as elastic lithosphere thickness and Chandler Wobble period favor a thicker elastic lithosphere and models with a frequency dependence α of the shear attenuation between 0.15 and 0.4. Improved constraints on the Martian interior should be possible with additional seismic and radio observations from the InSight mission. Plain Language Summary: The largest moon of Mars, Phobos, raises tides on the red planet causing its surface to be deformed. The size of this tidal bulge depends on the elastic properties of Mars, while viscosity and anelasticity inside the planet cause the tidal bulge to be misaligned with the position of Phobos on its orbit. This misalignment between Phobos and the tidal bulge of Mars creates an acceleration in longitude for the moon, which can be measured and used with the size of the tidal bulge to constraint the interior of Mars. Both the size and orientation of the tidal bulge can be used as tidal constraints to probe the size of the core of Mars, the temperature profile in its interior, and how strong the viscous dissipation is in its mantle. Our current estimates give a wide range of possible attenuation in the Martian mantle, favoring larger core sizes with larger elastic thickness of the lithosphere, the rigid outermost layer of a solid planet. Knowledge of the Martian interior could be further improved with future seismic and radio measurements from the InSight mission. Key Points: We develop a formulation to calculate the secular acceleration of Phobos due to tidal dissipation in MarsUsing the Love number k2 and the secular acceleration s, we obtain for Mars a core size of 1,820 ± 80 km and a bulk Q of 93.0 ± 8.40We discuss several means to distinguish between Mars models using seismological, rotational, and lithospheric observations [ABSTRACT FROM AUTHOR]

We derive a topography data set from images of Pluto and Charon that contain the body edge (i.e., limb profiles) which will help in understanding the comparative history of the binary system. We use the profiles to derive topographic variance spectra and find that while the variance spectrum of Pluto fits a single power law, Charon's spectrum displays a clear breakpoint at ∼150 km wavelength. Assuming the breakpoint is a result of topographic flexure, we find that Charon's elastic thickness must have been 20 ± 10 km during topography formation. A lack of a breakpoint for Pluto sets a minimum elastic thickness for Pluto of 60 km. We use these elastic thickness estimates to calculate a maximum heat flux of ∼13 m Wm−2 on Pluto during and after topography formation. On Charon, however, we find that the heat flux during topography formation was 35−15+44 m Wm−2. This range of values far exceeds the likely radiogenic heat production and is consistent with either heat released following the Charon‐forming impact event or (more likely) tidal heating during Charon's early history. Plain Language Summary: Studying the topography of planetary bodies provides key insights into the geologic processes of their surfaces and interiors. In this work we develop a topography data set for Pluto and Charon by mapping variations in the height along the worlds' edges in images from New Horizons. We analyze the data to determine roughness using the mean amplitude of mountains and valleys for a range of widths. Pluto shows the expected result of a single slope decreasing in roughness at shorter widths, but Charon has a change in the slope at ∼150 km. Mountains and valleys on Charon wider than this are respectively shorter and shallower than expected. This gives insight into how the landforms on Charon formed as well as the ability of Charon's crust to support variations in elevation. Charon's landforms must have formed at the observed size or decreased over time to have modern amplitudes. Either case implies that Charon had a thinner ice shell, and was relatively hotter, than Pluto in the ancient past. This extra heat is consistent with a Charon‐forming impact or (more likely) tidal heating during the Charon's initial history. Key Points: We derived new limb profile topography datasets of Pluto and Charon based on the body edge in imagesFor Charon, the topographic variance spectrum displays a distinct change in slope at ∼150 km wavelengthCharon's topography records high heat fluxes from either tidal heating or a giant impact [ABSTRACT FROM AUTHOR]

The textures and accretion ages of chondrites have been used to argue that their parent asteroids never differentiated. Without a core, undifferentiated planetesimals could not have generated magnetic fields through dynamo activity, so chondrites are not expected to have experienced such fields. However, the magnetic remanence carried by the CV chondrites is consistent with dynamo‐generated fields, hinting that partially differentiated asteroids consisting of an unmelted crust atop a differentiated interior may exist. Here, we test this hypothesis by applying synchrotron X‐ray microscopy to metallic veins in the slowly cooled H6 chondrite Portales Valley. The magnetic remanence carried by nanostructures in these veins indicates that this meteorite recorded a magnetic field over a period of tens to hundreds of years at ∼100 Myr after solar system formation. These properties are inconsistent with external field sources such as the nebula, solar wind, or impacts, but are consistent with dynamo‐generated fields, indicating that the H chondrite parent body contained an advecting metallic core and was therefore partially differentiated. We calculate the thermal evolution of the chondritic portions of partially differentiated asteroids that form through incremental accretion across 105 to 106 years, finding this can agree with the measured ages and cooling rates of multiple H chondrites. We also predict that the cores of these bodies could have been partially liquid and feasibly generating a dynamo at 100 Myr after solar system formation. These observations contribute to a growing body of evidence supporting a spectrum of internal differentiation within some asteroids with primitive surfaces. Plain Language Summary: Asteroids formed during the first few million years of the solar system through the accretion of billions of millimeter‐sized solids. If this process occurred within the first ∼2 Myr of the solar system, the asteroid is thought to have partially melted, while if it occurred after this time, the asteroid is thought to have remained completely unmelted. Partial melting is an easy mechanism allowing an asteroid to differentiate into a rocky mantle and metallic core. Recently, this discrete nature of asteroid melting has been challenged by magnetic measurements of a group of unmelted meteorites that suggest they experienced magnetic fields generated in an asteroid core, hinting that their parent asteroid contained both melted and unmelted material and was therefore partially differentiated. Here, we show that a previously unmeasured type of unmelted meteorite recorded a magnetic field over a period of tens to hundreds of years at ∼100 million years after solar system formation. These timings make this a particularly robust observation that some unmelted meteorites experienced dynamo fields and originate from partially differentiated asteroids. This observation favors the episodic formation of some asteroids, potentially impacting our understanding of the thermal and structural history of the first planetary bodies in our solar system. Key Points: The Portales Valley H6 chondrite experienced a magnetic field with properties consistent with dynamo fields at  100 Myr after CAI formationThis observation indicates that the H chondrite parent body contained an advecting metallic core, so was partially differentiatedWe model the thermal evolution of such bodies, finding that they can reproduce the measured ages and cooling rates of multiple H chondrites [ABSTRACT FROM AUTHOR]

Some large impact basins on the Moon are covered with younger basalts, making it hard to determine their age by crater counting. Evans et al. (2018, https://doi.org/10.1029/2017JE005421) use Gravity Recovery and Interior Laboratory gravity data to identify buried craters larger than 90 km and construct a relative chronology for the impact basins. This relative chronology is broadly consistent with the degree to which the basins have relaxed: Older basins tend to be more relaxed. Pre‐Nectarian basins are likely saturated, so the roughly constant measured crater density does not require a brief formation interval. Some of the results of Evans et al. (2018) are at odds with other determinations; for example, they suggest that Imbrium could be older than Crisium, and Serenitatis as young as Imbrium. Plain Language Summary: To tell the age of an impact basin on the Moon, you count the number of craters. But some basins are flooded by later lavas, making crater counting difficult. Evans et al. (2018) identify buried craters by looking for small changes in the Moon's gravity. This allows them to decide how old each impact basin is. The older basins are also generally shallower. This is because the older basins formed when the Moon was hotter and the lunar rocks were able to flow and partly fill the basin in. Key Points: Evans et al. (2018) use gravity data to identify buried craters within large basinsThey reinterpret the relative ages of lunar basins using these observationsThe relative ages correlate quite well with the degree of basin relaxation [ABSTRACT FROM AUTHOR]

Abstract: We invert the Martian tidal response and mean mass and moment of inertia for chemical composition, thermal state, and interior structure. The inversion combines phase equilibrium computations with a laboratory‐based viscoelastic dissipation model. The rheological model, which is based on measurements of anhydrous and melt‐free olivine, is both temperature and grain size sensitive and imposes strong constraints on interior structure. The bottom of the lithosphere, defined as the location where the conductive geotherm meets the mantle adiabat, occurs deep within the upper mantle (∼200–400 km depth) resulting in apparent upper mantle low‐velocity zones. Assuming an Fe‐FeS core, our results indicate (1) a mantle with a Mg# (molar Mg/Mg+Fe) of ∼0.75 in agreement with earlier geochemical estimates based on analysis of Martian meteorites; (2) absence of bridgmanite‐ and ferropericlase‐dominated basal layer; (3) core compositions (15–18.5 wt% S), core radii (1,730–1,840 km), and core‐mantle boundary temperatures (1620–1690°C) that, together with the eutectic‐like core compositions, suggest that the core is liquid; and (4) bulk Martian compositions with a Fe/Si (weight ratio) of 1.66–1.81. We show that the inversion results can be used in tandem with geodynamic simulations to identify plausible geodynamic scenarios and parameters. Specifically, we find that the inversion results are largely reproducible by stagnant lid convection models for a range of initial viscosities (∼1018–1020 Pa s) and radioactive element partitioning between crust and mantle around 0.01–0.1. The geodynamic models predict a mean surface heat flow between 15 and 25 mW/m2. [ABSTRACT FROM AUTHOR]