Search

The primary objective of the Europa Clipper mission is to assess the habitability of Europa, an overarching goal that rests on improving our understanding of Europa’s interior structure, composition, and geologic activity. Here we describe the Gravity and Radio Science (G/RS) investigation. The primary measurement, the gravitational tidal Love number $k_{2}$ k 2 , will be an independent diagnostic of the presence of a global subsurface ocean, but G/RS will make a number of other key measurements related to Europa’s deep interior, silicate mantle-ocean interface, ice shell, ionosphere, and plasma environment. Although radio science is common to many missions, Europa Clipper’s orbit and spacecraft configuration during flybys present special challenges for the design of this experiment. The information obtained through G/RS will be complementary to the measurements by the other instruments onboard Europa Clipper, and their combined analysis will refine the geophysical understanding of Europa necessary to best assess its potential habitability.

This is a white paper submitted to the Planetary Science and Astrobiology Decadal Survey. The deep atmosphere of Venus is largely unexplored and yet may harbor clues to the evolutionary pathways for a major silicate planet with implications across the solar system and beyond. In situ data is needed to resolve significant open questions related to the evolution and present-state of Venus, including questions of Venus' possibly early habitability and current volcanic outgassing. Deep atmosphere "probe-based" in situ missions carrying analytical suites of instruments are now implementable in the upcoming decade (before 2030), and will both reveal answers to fundamental questions on Venus and help connect Venus to exoplanet analogs to be observed in the JWST era of astrophysics.

Understanding the divergent evolution of Venus and Earth is a fundamental problem in planetary science. Although Venus today has a hot, dry atmosphere, recent modeling suggests that Venus may have had a clement surface with liquid water until less than 1 billion years ago [1]. Venus today has a nearly stagnant lithosphere. However, Ishtar Terra’s folded mountain belts, 8-11 km high, morphologically resemble Tibet and the Himalaya mountains on Earth and apparently require several thousand kilometers of surface motion at some time in Venus’s past. Loss of liquid surface water increases the coefficient of friction in fault zones, favoring a transition from an early mobile lithosphere to a present-day stagnant lithosphere [2]. Solar-driven climate evolution could contribute to a prolonged epoch of water loss on Venus and may be the ultimate cause of the divergent evolution of both the climate history and the interior dynamics of Venus and Earth. In addition to being a problem of first-order importance for Solar System evolution, understanding the divergent evolution of Venus and Earth is also important for understanding the temporal and spatial distribution of habitable environments in the Solar System. Understanding the evolution of Venus is also a key test for models that interpret Earth-sized exoplanets. Testing evolutionary hypotheses requires interpreting clues that were left behind in both the isotopic composition of the Venus atmosphere and in the rock record of the Venus surface. Although several mission concepts are currently competing for possible flights to Venus, only the Deep Atmosphere Venus Investigation of Noble gases, Chemistry, and Imaging Plus (DAVINCI+) mission [3] can examine both the atmospheric isotopic record and the rock record of Venus. DAVINCI+ is therefore a compelling choice for selection in the current NASA Discovery Program Phase A competition. DAVINCI+ includes an atmospheric entry probe and a carrier spacecraft (Figure 1). The probe measures atmospheric composition using a mass spectrometer and tunable laser spectrometer, performs descent imaging, and measures atmospheric structure. Following completion of the probe mission, the carrier spacecraft enters Venus orbit and images the Venus surface in the 1 micron atmospheric window. This payload is ideally suited for testing models of Venus evolution. Figure 1: The DAVINCI+ entry probe studies the atmosphere while imaging the Alpha Regio landing site. Afterward, the carrier probe performs infrared imaging of selected targets from orbit. The History of Water: Isotopic Record Pioneer Venus measured the D/H value of an H2SO4 cloud droplet at ~55 km as 157±30 times the terrestrial value, which was interpreted as the signature of escape of water from the Venus atmosphere [4]. Terrestrial spectroscopy produced a similar range, whereas Venus Express measured a value up to three times larger between 70-90 km [5]. The large value relative to Earth shows that Venus lost a substantial amount of water, but the large uncertainty and the lack of data below the clouds makes it difficult to quantitatively model the history of water loss [6]. DAVINCI+ will measure D/H to high precision from above the clouds down to the surface, greatly improving our ability to interpret the history of water loss on Venus. In addition, the abundance and isotopic ratios of Xe and Kr, together with the Kr/Ar and Xe/Ar ratios, which DAVINCI+ will measure with high precision, will be instrumental in revealing whether Venus and Earth formed in the same way and how their climates diverged. The History of Water: Rock Record A key but poorly answered question is the extent to which Venus has produced granitic or felsic (SiO2-rich) volcanism. Small amounts of felsic magma can be generated by lithospheric processes [7], but large amounts of felsic material requires the presence of water in the melting zone [8], as in terrestrial subduction zones. Tessera, which are regions of old, thick, highly tectonized crust, are widely accepted as the most likely location for felsic material on Venus. Venus Express observations suggest that tessera in Alpha Regio has a felsic composition [9]. DAVINCI+ will explore the presence and distribution of felsic rock in two ways. Comparison of the reflectivity at 1 micron and in panchromatic descent images at the Alpha Regio landing site will test the presence of felsic rock at patch sizes much smaller than can be observed from orbit. Descent imaging will also explore the landing zone geomorphology, and stereo topography will enable quantitative modelling of faulting and folding. Orbital imaging in the 1 micron atmospheric window will test for the presence of felsic rock in other tessera, including Tellus Regio, Fortuna Tessera, Maxwell Montes, Ovda Regio, and Thetis Regio. Our approach is similar to VIRTIS on Venus Express [9] but focuses on regions in the northern hemisphere and near the equator that were not imaged by VIRTIS. The History of Volcanism: Isotopic Record Volcanic outgassing releases radioactive decay products such as 40Ar and 4He to the atmosphere. DAVINCI+ measurements of their atmospheric abundance can be used to estimate volcanic outgassing over time. 40Ar provides an integrated record of volcanism over Venus history. Because 4He escapes from the atmosphere to space, its atmospheric abundance constrains geologically recent (last billion years) volcanism. Existing measurements of 40Ar and 4He are too imprecise to strongly constrain the volcanic history of Venus [10] but will be measured with much greater accuracy by DAVINCI+. Because DAVINCI+ will constrain the history of both water and volcanism on Venus, it will provide new insights into the feedbacks that shaped the divergent evolution of Venus and Earth. References [1] Way and Del Genio, JGR 125, e2019JE006276, 2020. [2] Weller and Kiefer, JGR 125, e2019JE005960, 2020. [3] Garvin et al., LPSC 51, abstract 2599, 2020. [4] Donahue et al., Science 216, 630-633, 1982. [5] Bertaux et al., Nature 450, 646-649, 2007. [6] Donahue et al., Venus II, 385-414, 1997. [7] Elkins-Tanton et al., JGR 112, E04S06, 2007. [8] Campbell and Taylor, GRL 10, 1061-1064, 1983. [9] Gilmore et al., Icarus 254, 350-361, 2015. [10] Namiki and Solomon, JGR 103, 3655-3677, 1998.

If Venus were a newly discovered exoplanet, it would be one of the most Earth-like yet identified. Its similarity in size, bulk density and cloud top temperatures give no clue to the hellish temperatures at its surface. Venus hosts an array of geological features as complex as Earth, but without its organisation, and sustains a chemically reactive atmosphere, but without life. Proposed in response to ESA’s M5 call, with enabling support from NASA, EnVision is currently in Phase A study for the next medium-class mission opportunity proceeding towards a final mission selection in summer 2021 for a launch in 2032. Here we discuss the science questions motivating the mission in more detail. Building on discoveries from Magellan, Venus Express and Akatsuki, EnVision will focus on three overarching questions: Is Venus geodynamically active? How did Venus arrive at its current state? How does the Venus climate work and how do the interior, surface and atmosphere interact? 1 – Is Venus geodynamically active? Venus should be geologically active today and its globally young surface implies extensive volcanic resurfacing, but whether this happened in episodic global events or in a continuous process of small-scale resurfacing is uncertain. Constraining the rate of volcanic and tectonic activity can reveal whether Venus occupies one of these end members, or lies somewhere on the spectrum between. Magellan radar image mosaic of a 550 km wide region of Lada Terra (47°S, 25°E) showing a system of radar-bright and dark lava flows breaching a ridge belt and ponding in a large radar-bright deposit covering 100,000 km². Ammavaru, the source caldera, lies 300 km to the west (JPL/Caltech). Understanding the tectonic regime driving resurfacing and heat loss is as important as knowing the rate of activity. The surface appears partitioned into areas of low strain bounded by narrow high strain margins. The abundance of steep slopes and landslides implies active uplift in these high strain areas, but existing data provide no constraint on the frequency of landslides or rates of tectonic movement. Are the low strain regions actively created and destroyed, like Earth’s oceanic plates, or simply mobilised locally? What is the significance of the global network of elevated rift systems and linear lowland wrinkle-ridged plains? Unique to Venus are coronae, quasi-circular volcano-tectonic features, typically 100–500 km across. Are they the surface expression of plumes, or magmatic intrusions, or subduction zones? The processes of weathering, mass wasting and aeolian transport are critical for understanding both the geological history and climatic evolution of Venus. Venus Express found anomalously high IR emissivities near suspected active volcanoes, interpreted as fresh, unweathered lava flows, but both their mineralogy and the weathering processes involved are unknown. Magellan detected several, probably impact-related, dune fields at the limit of its resolution, and found indirect evidence for globally distributed smale-scale ripples. 2 – How Did Venus Arrive at its Current State? The cratering record reveals that the Venus surface is, on average, under 1 Ga old. While some areas are likely active today, other regions, e.g. the tessera highlands, may be considerably older. Magellan data imply a variety of age relationships and long-term activity, with a non-random association between geological features and elevation, e.g. the uplands are consistently more deformed than the lowlands. Impact craters themselves show alteration to dark floors, perhaps from airfall deposits or magmatic resurfacing. Densely fractured plains (right) abutting tessera (bottom); both are embayed by younger plains (dark areas). A steep-sided volcanic dome (upper left) may be the source of the plains or unrelated to them, and older or younger than the tesserae. Magellan image at 46°N, 360°E, after Basilevsky and McGill (2007). Constraining the history is critical to understanding not only when and how resurfaced, but whether that activity has changed systematically through time. Were the plains formed in a short period by massive outpourings or by many thousands of small flows over their entire history? Or did many different mechanisms – including sedimentary processes – operate at different times and places? Globally, Venus exhibits perhaps even more tectonic deformation than Earth: what is its role in resurfacing the planet and have the regimes of tectonic deformation changed over time? How did tesserae, especially, accumulate their extraordinary degree of deformation? The interior of Venus is probably Earth-like, but not the same; its core size is poorly constrained by Magellan gravity data. These data are consistent with an organised pattern of mantle convection but lack the resolution necessary to connect it with geological-scale features. 3 - How does the Venus climate work and how do the Interior, Surface and Atmosphere Interact? How and why the atmospheres of two Earth-like planets evolved so differently is one of the many compelling reasons to study Venus, directly addressing the question of how our own world became habitable. Understanding its evolution depends on knowing the exchanges between its interior, surface and space. The lower/middle clouds of Venus as imaged by Akatsuki IR2 camera. Dark areas represent thicker clouds, which may represent volcanic plumes of ash or sulphate particulates. (JAXA/ISAS/DART/Damia Bouic) Models suggest that the clouds are maintained by a constant input of H2O and SO2, both of which vary considerably, perhaps because of volcanic emissions. Elevated D/H ratios indicate the loss of early oceans, but the ‘starting’ mantle D/H ratio is unknown. Identification of granites in the ancient highlands would support this inference, but the direct detection and characterisation of volcanic volatiles is key to inferring past climate evolution. Beyond the specific investigations outlined above, however, the most important unknown factor in determining the present day state of Venus and its atmosphere, and how it arrived at that state, is the interaction between the interior, surface and atmosphere: the whole being more than the sum of the parts. References Basilevsky and McGill (2007), Surface evolution of Venus, in Exploring Venus as a Terrestrial Planet, Geophysical Monograph Series 176, edited by L. W. Esposito, E. R. Stofan and T. E. Cravens, pp. 23-43, American Geophysical Union, Washington, DC

The last time scientists got a detailed view of the surface of Venus was with the revolutionary NASA Magellan Mission more than 30 years ago. Two new mission concepts under study by NASA and ESA are designed to return to Venus and obtain refined imagery and mapping data to answer some of the major questions concerning the evolution and current geologic activity of Venus. Both these missions would employ synthetic aperture radars to view the surface through the thick optically opaque atmosphere. Each radar design was optimized to meet the science objectives of their respective missions and bring back new data to the planetary science community to unravel why Venus, our twin planet, is so very different than the Earth. This paper discusses the complementary nature of these missions, the design and performance of these two radar instruments, and how each will further Venus and planetary science.

The Apollo Heat Flow Experiment (HFE) was conducted at landing sites 15 and 17. On Apollo 15, surface and subsurface temperatures were monitored from July 1971 to January 1977. On Apollo 17, monitoring took place from December 1972 to September 1977. The investigators involved in the HFE examined and archived only data from the time of deployment to December 1974. The present authors recovered and restored major portions of the previously un-archived HFE data from January 1975 through September 1977. The HFE investigators noted that temperature of the regolith well below the reach of insolation cycles (approx. 1 m) rose gradually through December 1974 at both sites. The restored data showed that the subsurface warming continued until the end of observations in 1977. Simultaneously, the thermal gradient decreased, because the warming was more pronounced at shallower depths. The present study has examined potential causes for the warming. Recently acquired images of the Lunar Reconnaissance Orbiter Camera over the two landing sites show that the regolith on the paths of the astronauts turned darker, lowering the albedo. We suggest that, as a result of the astronauts' activities, solar heat intake by the regolith increased slightly on average, and that resulted in the observed warming. Simple analytical heat conduction models with constant regolith thermal properties can show that an abrupt increase in surface temperature of 1.6 K to 3.5 K at the time of probe deployment best duplicate the magnitude and the timing of the observed subsurface warming's at both Apollo sites.

Orbital measurements of gravity and topography demonstrate that the crust of Mars varies in thickness by about 100 km across the surface. As a result, the heat flux due to crustal radioactivity varies laterally by nearly 15 mW m−2. In an important new paper, Plesa et al. (2016, doi:10.1002/2016JE005126) have assessed how these variations in crustal thickness and heat flow affect the spatial pattern of convection in the Martian mantle. Their results have important implications for the interpretation of the Martian heat flux that will be measured by the upcoming InSight mission to Mars.

Based on meteorite evidence, the present-day Martian mantle has a combined abundance of up to a few hundred ppm of H2O, Cl, and F, which lowers the solidus and enhances the magma production rate. Adiabatic decompression melting in upwelling mantle plumes is the best explanation for young (last 200 Myr) volcanism on Mars. We explore water undersaturated mantle plume volcanism using a finite element mantle convection model coupled to a model of hydrous peridotite melting. Relative to a dry mantle, the reduction in solidus temperature due to water increases the magma production rate by a factor of 1.3–1.7 at 50 ppm water and by a factor of 1.9–3.2 at 200 ppm water. Mantle water also decreases the viscosity and increases the vigor of convection, which indirectly increases the magma production rate by thinning the thermal boundary layer and increasing the flow velocity. At conditions relevant to Mars, these indirect effects can cause an order of magnitude increase in the magma production rate. Using geologic and geophysical observations of the Late Amazonian magma production rate and geochemical observations of melt fractions in shergottite meteorites, present-day Mars is constrained to have a core–mantle boundary temperature of ~1750 to 1800 °C and a volume-averaged thermal Rayleigh number of 2 × 106 to 107, indicating that moderately vigorous mantle convection has persisted to the present day. Melting occurs at depths of 2.5–6 GPa and is controlled by the Rayleigh number at the low pressure end and by the mantle water concentration at high pressure.

International audience; The interior structure of the Moon is constrained by its mass, moment of inertia, and k 2 and h 2 tidal Love numbers. We infer the likely radius, density, and (elastic limit) rigidity of all interior layers by solving the inverse problem using these observational constraints assuming spherical symmetry. Our results do not favor the presence of a low rigidity transition layer between a liquid outer core and mantle. If a transition layer exists, its rigidity is constrained to 43 +26 −9 GPa, with a preference for the high rigidity values. Therefore, if a transition layer exists, it is more likely to have a rigidity similar to that of the mantle (∼70 GPa). The total (solid and liquid) core mass fraction relative to the lunar mass is constrained to 0.0098 +0.0066 −0.0094 and 0.0198 +0.0026 −0.0049 for interior structures with and without a transition layer, respectively, narrowing the range of possible giant impact formation scenarios.

We define lunar cryptovolcanism as volcanic deposits on the Moon hidden by overlying material. Notably, cryptovolcanism includes both cryptomaria (subsurface extrusive basaltic deposits that are obscured by overlying higher albedo basin and crater ejecta) and earlier candidate extrusives, such as the Mg-suite. Knowledge of the volume and extent of cryptovolcanism is necessary for a comprehensive understanding of lunar volcanic history, particularly in early (pre 3.8 Ga) epochs when abundant impact craters and basins obscured surface volcanic deposits by lateral emplacement of ejecta. We use Gravity Recovery and Interior Laboratory (GRAIL) gravity and Lunar Orbiter Laser Altimeter (LOLA) topography data to construct maps of the Moon's positive Bouguer and isostatic gravity anomalies, and explore the possibility that these features are due to mass excesses associated with cryptovolcanism by cross-referencing the regions with geologic data such as dark halo craters. We model the potential cryptovolcanic deposits as buried high-density rectangular prisms at depth in the upper crust, and find a volume of candidate buried cryptovolcanism between 0.4 × 10 6 km 3 and 4.8 × 10 6 km 3 , depending on assumptions about density and crustal compensation state. These candidate deposits correspond to a surface area of between 0.50 × 10 6 km 2 and 1.14 × 10 6 km 2 , which would increase the amount of the lunar surface containing volcanic deposits from 16.6% to between 17.9% and 19.5%. The inferred volume of cryptovolcanism is comparable to the smallest estimates of the volume of visible mare basalts and up to ∼50% of the largest estimates; the high-resolution GRAIL and LOLA observations thus would collectively indicate that early (pre 3.8 Ga) lunar volcanism is an important element of lunar thermal evolution. Alternatively, the buried material could represent the presence of intrusive Mg-suite sills or plutons.

An exciting and novel science mission concept called Seismic and Atmospheric Exploration of Venus (SAEVe) has been developed which uses high-temperature electronics to enable a three-order magnitude increase in expected surface life (120 Earth days) over what has been achieved to date. This enables study of long-term, variable phenomena such as the seismicity of Venus and near surface weather, near surface energy balance, and atmospheric chemical composition. SAEVe also serves as a critical pathfinder for more sophisticated landers in the future. For example, first order seismic measurements by SAEVe will allow future missions to deliver better seismometers and systems to support the yet unknown frequency and magnitude of Venus events. SAEVe is focused on science that can be realized with low data volume instruments and will most benefit from temporal operations. The entire mission architecture and operations maximize science while minimizing energy usage and physical size and mass. The entire SAEVe system including its protective entry system is estimated to be around 45 ​kg and approximately 0.6 ​m diameter. These features allow SAEVe to be relatively cost effective and be easily integrated onto a Venus orbiter mission. The technologies needed to implement SAEVe are currently in development by several funded activities. Component and system level work is ongoing under NASA’s Long Lived Insitu Solar System Explorer (LLISSE) project and the HOTTech program. . LLISSE, is a NASA project to develop a small Venus lander that will operate on the surface of Venus for 60 days and measure variations in meteorology, radiance, and atmospheric chemistry. LLISSE is developing a full-function engineering model of a Venus lander that contains essentially all the core capabilities of SAEVe thus greatly reducing the technology risk to SAEVe. The SAEVe long duration Venus lander promises exciting new science and is an ideal complimentary element to many future Venus orbiter missions being proposed or planned today.

Fourteen types of geophysical instruments deployed at the Apollo 12, 14, 15, 16, and 17 sites by the astronauts for long-term observation were collectively called the Apollo Lunar Surface Experiments Package (ALSEP). These instruments were active from the times of their deployment (November 1969–December 1972) to September 1977. At the conclusion of the experiments, the raw instrument data received from the Moon prior to March 1976 were left unarchived. Portions of the data processed by the principal investigators (PIs) of these experiments had been archived at the NASA Space Science Data Coordinated Archive (NSSDCA) in various formats. The unarchived data, residing then on open-reel magnetic tapes, became lost in the decades since, along with much of the metadata (the supporting documents for these data). We have recently recovered 440 of the previously lost tapes, containing raw ALSEP instrument data from April through June of 1975. Here we describe the data extracted from these tapes and summarize the data products generated for archiving at the NASA Planetary Data System (PDS) and NSSDCA, along with their historical narrative. In addition, we have reformatted many of the datasets delivered to NSSDCA by the PIs in the 1970s for archiving at the PDS. Finally, we have compiled an online searchable repository of ALSEP-related documents by optically scanning tens of thousands of pages of them kept at the Lunar and Planetary Institute in Texas.

© 2017 Elsevier Inc. The Orientale basin is the youngest and best-preserved multiring impact basin on the Moon, having experienced only modest modification by subsequent impacts and volcanism. Orientale is often treated as the type example of a multiring basin, with three prominent rings outside of the inner depression: the Inner Rook Montes, the Outer Rook Montes, and the Cordillera. Here we use gravity data from NASA's Gravity Recovery and Interior Laboratory (GRAIL) mission to reveal the subsurface structure of Orientale and its ring system. Gradients of the gravity data reveal a continuous ring dike intruded into the Outer Rook along the plane of the fault associated with the ring scarp. The volume of this ring dike is ∼18 times greater than the volume of all extrusive mare deposits associated with the basin. The gravity gradient signature of the Cordillera ring indicates an offset along the fault across a shallow density interface, interpreted to be the base of the low-density ejecta blanket. Both gravity gradients and crustal thickness models indicate that the edge of the central cavity is shifted inward relative to the equivalent Inner Rook ring at the surface. Models of the deep basin structure show inflections along the crust–mantle interface at both the Outer Rook and Cordillera rings, indicating that the basin ring faults extend from the surface to at least the base of the crust. Fault dips range from 13° to 22° for the Cordillera fault in the northeastern quadrant, to 90° for the Outer Rook in the northwestern quadrant. The fault dips for both outer rings are lowest in the northeast, possibly due to the effects of either the direction of projectile motion or regional gradients in pre-impact crustal thickness. Similar ring dikes and ring faults are observed around the majority of lunar basins.

At a given pressure, terrestrial peridotites of varying composition may have solidus temperatures that differ by up to 100 °C. Based on meteorite evidence, the mantle of Mars is believed to be enriched in Na and K and to have a higher Fe/Mg ratio (lower magnesium number, Mg#) than Earth. These differences all favor a mantle solidus temperature on Mars that is lower than on Earth and are important in understanding the volcanic history of Mars. We parameterize the peridotite solidus at 1 and 3 GPa as a function of Mg# and total alkali content, using existing measurements of peridotite melting at 1 and 3 GPa for Mg# between 75 and 91 and total alkali content between 0.06 and 1.17 weight percent. The solidus on early Mars was likely 30–40 °C lower than on Earth, which increases the predicted crustal production by about 20% over martian history, relative to a Mars model that uses a solidus calculated for terrestrial peridotite composition. Because Na is incompatible and migrates to the crust over time, the present-day martian solidus is higher than the primitive solidus but is still ∼15 °C less than on Earth. This enhances the present-day magma production rate at martian mantle plumes by a factor of 2–3.

Lunar roughness measurements derived from the Lunar Orbiter Laser Altimeter are compared to 12.6 cm wavelength radar data collected by the Miniature Radio Frequency instrument and 70 cm wavelength radar data collected by the Arecibo Observatory. These data are compared to assess how surface and subsurface roughness are correlated and affected by parameters including age and composition at length scales between 0.1 and 100 m. A range of features are analyzed including volcanic domes (Marius Hills, Rumker Hills, Gruithuisen, and Mairan Domes); mare (Imbrium, Serenitatis, and Oceanus Procellarum); pyroclastic dark mantle deposits (Sinus Aestuum, Sulpicius Gallus, and Mare Vaporum); and two young craters (Copernicus and Tycho). Statistically significant positive correlations exist between topographic roughness and both P- and S-band circular polarization ratios. The strongest correlation is observed at the longest length scales. Correlations weaken as length scales become less similar, potentially due to distinct processes controlling surface modification. Roughness is not significantly correlated with local slope. Although the Marius Hills are compositionally distinct from the Gruithuisen and Mairan domes, they are indistinguishable in roughness characteristics. Conversely, the Rumker Hills, mare, and dark mantle deposits are smoother at the length scales examined, possibly due to fine-grained mantling of regolith or pyroclastic deposits. The floor and ejecta of Tycho are the roughest surfaces measured in this study, while the floor and ejecta of Copernicus overlap the roughness distribution of the volcanic features. This study shows that many factors control the evolution of roughness over time on various length scales.

International audience; We analyzed data from the Gravity Recovery and Interior Laboratory (GRAIL) mission using a localized admittance approach to map out spatial variations in the vertical density structure of the lunar crust. Mare regions are characterized by a distinct decrease in density with depth, while the farside is characterized by an increase in density with depth at an average gradient of ∼35 kg m −3 km −1 and typical surface porosities of at least 20%. The Apollo 12 and 14 landing site region has a similar density structure to the farside, permitting a comparison with seismic velocity profiles. The interior of the South Pole-Aitken (SP-A) impact basin appears distinct with a near-surface low-density (porous) layer 2-3 times thinner than the rest of the farside. This result suggests that redistribution of material during the large SPA impact likely played a major role in sculpting the lunar crust.

The Gravity Recovery and Interior Laboratory (GRAIL) mission has sampled lunar gravity with unprecedented accuracy and resolution. The lunar GM, the product of the gravitational constant G and the mass M, is very well determined. However, uncertainties in the mass and mean density, 3345.56 ± 0.40 kg/m3, are limited by the accuracy of G. Values of the spherical harmonic degree-2 gravity coefficients J2 and C22, as well as the Love number k2 describing lunar degree-2 elastic response to tidal forces, come from two independent analyses of the 3 month GRAIL Primary Mission data at the Jet Propulsion Laboratory and the Goddard Space Flight Center. The two k2 determinations, with uncertainties of ~1%, differ by 1%; the average value is 0.02416 ± 0.00022 at a 1 month period with reference radius R = 1738 km. Lunar laser ranging (LLR) data analysis determines (C − A)/B and (B − A)/C, where A

[1] The Orientale impact basin is the youngest and best-preserved lunar multi-ring basin and has, thus, been the focus of studies investigating basin-forming processes and final structures. A consensus about how multi-ring basins form, however, remains elusive. Here we numerically model the Orientale basin-forming impact with the aim of resolving some of the uncertainties associated with this basin. By using two thermal profiles estimating lunar conditions at the time of Orientale's formation and constraining the numerical models with crustal structures inferred from gravity data, we provide estimates for Orientale's impact energy (2–9 × 1025 J), impactor size (50–80 km diameter), transient crater size (∼320–480 km), excavation depth (40–55 km), and impact melt volume (∼106 km3). We also analyze the distribution and deformation of target material and compare our model results and Orientale observations with the Chicxulub crater to investigate similarities between these two impact structures.

[1] The volcanic style of the Moon has long been understood to consist almost exclusively of flood basalts erupted from fissures along with minor pyroclastic activity; large central vent shield volcanoes that characterize basaltic volcanism on the other terrestrial planets appeared to be absent. Small (few kilometers diameter) central vent constructs have long been recognized in the lunar maria and often are found clustered in fields throughout the lunar maria. New global topographic data from the LOLA and LROC instruments on LRO reveal that almost all of these volcanic complexes on the Moon occur on large, regional topographic rises in the lunar maria, tens to hundreds of kilometers in extent and between several hundred to several thousand meters high. We propose that these topographic swells are shield volcanoes and are the lunar equivalents of the large basaltic shields found on the Earth, Venus, and Mars. The newly recognized lunar shields are found peripheral to the large, deeply flooded impact basins Imbrium and Serenitatis, suggesting a genetic relation to those features. Loading of the lithosphere by these basalt-filled basins may be responsible for inducing a combination of flexural and membrane stress, inducing a pressure distribution on vertically oriented dikes favorable to magma ascent. This condition would occur in a zone annular to the large circular loads produced by the basins, where the shield volcanoes occur.

[1] The Marius Hills, the Moon's largest volcanic dome field, has more than 250 basaltic domes and cones in an area 200 × 250 km across. It is a major free-air gravity anomaly, 236 mGal in the north and 150 mGal in the south. In the northern half of the structure, the topography can only explain about half of the gravity anomaly, and in the south, there is virtually no topographic relief associated with the gravity anomaly. High-density material must be present at depth, most likely as mare basalt intruded into the underlying porous feldspathic highland crust. The gravity anomaly is modeled using two spherical caps. The northern cap is 160–180 km in diameter and at least 3.0 km thick. The southern cap is 100–140 km in diameter and at least 6.2 km thick. The intruded basalt may have served as the magma chambers that fed the overlying surface volcanism. Magma crystallization within these chambers provided a source of crystal-rich, high viscosity lava that fed the volcanic domes. The volume of intruded basalt is 1.6 × 104 km3. The total volcanic volume, including both intruded and extruded material, is 2.6 × 104 km3, indicating that the Marius Hills is a major volcanic center. Intrusion of hot magma may cause thermal annealing of the porous feldspathic host rock, significantly reducing the host rock porosity. This would allow a large volume of magma to be intruded into the crust with little change in overall crustal volume.

The Holy GRAIL? The gravity field of a planet provides a view of its interior and thermal history by revealing areas of different density. GRAIL, a pair of satellites that act as a highly sensitive gravimeter, began mapping the Moon's gravity in early 2012. Three papers highlight some of the results from the primary mission. Zuber et al. (p. 668 , published online 6 December) discuss the overall gravity field, which reveals several new tectonic and geologic features of the Moon. Impacts have worked to homogenize the density structure of the Moon's upper crust while fracturing it extensively. Wieczorek et al. (p. 671 , published online 6 December) show that the upper crust is 35 to 40 kilometers thick and less dense—and thus more porous—than previously thought. Finally, Andrews-Hanna et al. (p. 675 , published online 6 December) show that the crust is cut by widespread magmatic dikes that may reflect a period of expansion early in the Moon's history.

The Orientale basin is the youngest and best-preserved major impact structure on the Moon. We used the Gravity Recovery and Interior Laboratory (GRAIL) spacecraft to investigate the gravitational field of Orientale at 3- to 5-kilometer (km) horizontal resolution. A volume of at least (3.4 ± 0.2) × 106 km3 of crustal material was removed and redistributed during basin formation. There is no preserved evidence of the transient crater that would reveal the basin’s maximum volume, but its diameter may now be inferred to be between 320 and 460 km. The gravity field resolves distinctive structures of Orientale’s three rings and suggests the presence of faults associated with the outer two that penetrate to the mantle. The crustal structure of Orientale provides constraints on the formation of multiring basins.

Data from the Gravity Recovery and Interior Laboratory (GRAIL) mission have revealed that ∼98% of the power of the gravity signal of the Moon at high spherical harmonic degrees correlates with the topography. The remaining 2% of the signal, which cannot be explained by topography, contains information about density variations within the crust. These high-degree Bouguer gravity anomalies are likely caused by small-scale (10′s of km) shallow density variations. Here we use gravity inversions to model the small-scale three-dimensional variations in the density of the lunar crust. Inversion results from three non-descript areas yield shallow density variations in the range of 100–200 kg/m 3 . Three end-member scenarios of variations in porosity, intrusions into the crust, and variations in bulk crustal composition were tested as possible sources of the density variations. We find that the density anomalies can be caused entirely by changes in porosity. Characteristics of density anomalies in the South Pole-Aitken basin also support porosity as a primary source of these variations. Mafic intrusions into the crust could explain many, but not all of the anomalies. Additionally, variations in crustal composition revealed by spectral data could only explain a small fraction of the density anomalies. Nevertheless, all three sources of density variations likely contribute. Collectively, results from this study of GRAIL gravity data, combined with other studies of remote sensing data and lunar samples, show that the lunar crust exhibits variations in density by ± 10% over scales ranging from centimeters to 100′s of kilometers.

International audience; A new method of determining the thickness of mare basalts on the Moon is introduced that is made possible by high-resolution gravity data acquired from NASA's Gravity Recovery and Interior Laboratory (GRAIL) mission. Using a localized multitaper spherical-harmonic analysis, an effective density spectrum is calculated that provides an estimate of the average crustal density as a function of spherical harmonic degree. By comparing the observed effective density spectrum with one generated from a theoretical model, the thickness of mare basalts can be constrained. We assume that the grain density of the basalts is known from remote sensing data and petrologic considerations, we assign a constant porosity to the basalts, and we let both the thickness of the basalts and the density of the underlying crust vary. Using this method, the total thickness of basalts was estimated on the nearside hemisphere, yielding an average of 0.74 km with 1 upper and lower bounds of 1.62 km and 100 m, respectively. The region of Marius Hills, which is a long-lived volcanic complex, is found to have the thickest basalts, with an average of 2.86 km and 1 limits of 3.65 and 1.02 km, respectively. The crust beneath the Mare Imbrium basalts is found to have an atypically high density of about 3000 kg m −3 that we interpret as representing a mafic, unfractured, impact melt sheet.

Asteroid 4 Vesta is differentiated into a crust, mantle, and core, as demonstrated by studies of the eucrite and diogenite meteorites and by data from NASA's Dawn spacecraft. Most models for the differentiation and thermal evolution of Vesta assume that the metal phase completely melts within 20 degrees of the eutectic temperature, well before the onset of silicate melting. In such a model, core formation initially happens by Darcy flow, but this is an inefficient process for liquid metal and solid silicate. However, the likely chemical composition of Vesta, similar to H chondrites with perhaps some CM or CV chondrite, has 13-16 weight percent S. For such compositions, metal-sulfide melting will not be complete until a temperature of at least 1350 degrees Centigrade. The silicate solidus for Vesta's composition is between 1100 and 1150 degrees Centigrade, and thus metal and silicate melting must have substantially overlapped in time on Vesta. In this chemically and physically more likely view of Vesta's evolution, metal sulfide drops will sink by Stokes flow through the partially molten silicate magma ocean in a process that can be envisioned as "iron rain". Measurements of eucrites show that moderately siderophile elements such as Ni, Mo, and W reached chemical equilibrium between the metal and silicate phases, which is an important test for any Vesta differentiation model. The equilibration time is a function of the initial metal grain size, which we take to be 25-45 microns based on recent measurements of H6 chondrites. For these sizes and reasonable silicate magma viscosities, equilibration occurs after a fall distance of just a few meters through the magma ocean. Although metal drops may grow in size by merger with other drops, which increases their settling velocities and decreases the total core formation time, the short equilibration distance ensures that the moderately siderophile elements will reach chemical equilibrium between metal and silicate before metal drop merger becomes important. In this model, there must be at least 30 percent melting of the silicate phase when metal melting is complete, corresponding to a crust thickness of at least 30 kilometers on Vesta, consistent with Dawn gravity observations. Greater degrees of silicate melting and a correspondingly thicker crust are possible if Vesta accreted sufficiently rapidly.

The South Pole-Aitken (SPA) basin is the largest and oldest definitive impact structure on the Moon. To understand how this immense basin formed, we conducted a suite of SPA-scale numerical impact simulations varying impactor size, impact velocity, and lithospheric thermal gradient. We compared our model results to observational SPA basin data to constrain a best-fit scenario for the SPA basin-forming impact. Our results show that the excavation depth-to-diameter ratio for SPA-scale impacts is constant for all impact scenarios and is consistent with analytical and geological estimates of excavation depth in smaller craters, suggesting that SPA-scale impacts follow proportional scaling. Steep near-surface thermal gradients and high internal temperatures greatly affected the basin-forming process, basin structure and impact-generated melt volume. In agreement with previous numerical studies of SPA-scale impacts, crustal material is entirely removed from the basin center which is instead occupied by a large melt pool of predominantly mantle composition. Differentiation of the melt pool is needed to be consistent with observational data. Assuming differentiation of the thick impact-generated melt sheet occurred, and using observational basin data as constraints, we find the best-fit impact scenario for the formation of the South Pole-Aitken basin to be an impact with an energy of ∼4 × 1026 J (our specific model considered an impactor 170 km in diameter, striking at 10 km/s).

Reliable measurements of the Moon's global heat flow would serve as an important diagnostic test for models of lunar thermal evolution and would also help to constrain the Moon's bulk abundance of radioactive elements and its differentiation history. The two existing measurements of lunar heat flow are unlikely to be representative of the global heat flow. For these reasons, obtaining additional heat flow measurements has been recognized as a high priority lunar science objective. In making such measurements, it is essential that the design and deployment of the heat flow probe and of the parent spacecraft do not inadvertently modify the near-surface thermal structure of the lunar regolith and thus perturb the measured heat flow. One type of spacecraft-related perturbation is the shadow cast by the spacecraft and by thermal blankets on some instruments. The thermal effects of these shadows propagate by conduction both downward and outward from the spacecraft into the lunar regolith. Shadows cast by the spacecraft superstructure move over the surface with time and only perturb the regolith temperature in the upper 0.8 m. Permanent shadows, such as from thermal blankets covering a seismometer or other instruments, can modify the temperature to greater depth. Finite element simulations using measured values of the thermal diffusivity of lunar regolith show that the limiting factor for temperature perturbations is the need to measure the annual thermal wave for 2 or more years to measure the thermal diffusivity. The error induced by permanent spacecraft thermal shadows can be kept below 8% of the annual wave amplitude at 1 m depth if the heat flow probe is deployed at least 2.5 m away from any permanent spacecraft shadow. Deploying the heat flow probe 2 m from permanent shadows permits measuring the annual thermal wave for only one year and should be considered the science floor for a heat flow experiment on the Moon. One way to meet this separation requirement would be to deploy the heat flow and seismology experiments on opposite sides of the spacecraft. This result should be incorporated in the design of future lunar geophysics spacecraft experiments. Differences in the thermal environments of the Moon and Mars result in less restrictive separation requirements for heat flow experiments on Mars.

New gravity measurements greatly improve the Moon’s preserved impact basin inventory., Observations from the Gravity Recovery and Interior Laboratory (GRAIL) mission indicate a marked change in the gravitational signature of lunar impact structures at the morphological transition, with increasing diameter, from complex craters to peak-ring basins. At crater diameters larger than ~200 km, a central positive Bouguer anomaly is seen within the innermost peak ring, and an annular negative Bouguer anomaly extends outward from this ring to the outer topographic rim crest. These observations demonstrate that basin-forming impacts remove crustal materials from within the peak ring and thicken the crust between the peak ring and the outer rim crest. A correlation between the diameter of the central Bouguer gravity high and the outer topographic ring diameter for well-preserved basins enables the identification and characterization of basins for which topographic signatures have been obscured by superposed cratering and volcanism. The GRAIL inventory of lunar basins improves upon earlier lists that differed in their totals by more than a factor of 2. The size-frequency distributions of basins on the nearside and farside hemispheres of the Moon differ substantially; the nearside hosts more basins larger than 350 km in diameter, whereas the farside has more smaller basins. Hemispherical differences in target properties, including temperature and porosity, are likely to have contributed to these different distributions. Better understanding of the factors that control basin size will help to constrain models of the original impactor population.

Devana Chasma is a rift system on Venus formed in association with the Beta Regio and Phoebe Regio volcanic highlands, which are interpreted as mantle plumes. We present a new analysis of a 2500-km-long segment of Devana. Based on the rift topography, the horizontal extension across the rift boundary faults is 3–9 km. This is a lower bound on the total rift extension because the altimetry does not resolve the topographic relief across the numerous faults that are visible in radar images of the rift floor. The total extension across Devana is approximately 20 km, similar in magnitude to continental rift systems on Earth. Rift flank elevations are up to 3.1 km in the regions nearest the mantle plumes and decay strongly with increasing distance from the plumes, indicating a strong thermal component to the rift flank topography, unlike the situation usually reported for terrestrial rifts. As on Earth, there is also a flexural uplift component to the flank topography. Rift depths are up to 2.5 km below the surrounding plains, with considerable along-strike variability. There is a 600 km lateral offset along Devana Chasma near the mid-point between the two mantle plumes. Devana most likely formed as two distinct rifts due to the horizontal stresses created by outflow from the upwelling plumes. The offset zone formed as a result of the interaction between the two rift tips, which requires that upwelling at the two mantle plumes overlapped in time.

The Martian meteorite Yamato (Y-) 980459 is an olivine-phyric shergottite. It has a very primitive character and may be a primary melt of the Martian mantle. We have conducted crystallization experiments on a synthetic Y-980459 composition at Martian upper mantle conditions in order to test the primary mantle melt hypothesis. Results of these experiments indicate that the cores of the olivine megacrysts in Y-980459 are in equilibrium with a melt of bulk rock composition, suggesting that these megacrysts are in fact phenocrysts that grew from a magma of the bulk rock composition. Multiple saturation of the melt with olivine and a low-calcium pyroxene occurs at approximately 12 ± 0.5 kbar and 1540 ±10 °C, suggesting that the meteorite represents a primary melt that separated from its mantle source at a depth of ~100 km. Several lines of evidence suggest that the Y-980459 source underwent extensive melting prior to and/or during the magmatic event that produced the Y-980459 parent magma. When factored into convective models of the Martian interior, the high temperature indicated for the upper Martian mantle and possibly high melt fraction for the Y-980459 magmatic event suggests a significantly higher temperature at the core-mantle boundary than previously estimated.

Syrtis Major is an ancient basaltic shield volcano on Mars with a basal diameter of 1100 km. The free-air gravity anomaly is 126 mGal at spherical harmonic degree 50 and reaches its maximum amplitude over the 2 km deep topographic caldera. The observed gravity anomaly cannot be explained by flexurally supported surface topography and requires the presence of a buried, high-density load. The geologically most reasonable interpretation of this high-density load is that it represents the magma chamber of Syrtis Major, now solidified and filled at least in part by dense igneous cumulates. Pyroxene is likely to be the dominant cumulate mineral in this system, although olivine may also be present. Gravity models presented here define the structure of the buried load and in essence provide a look at the magmatic plumbing system of this volcano. The preferred model involves a buried load that is approximately 300×600 km across, roughly twice as large as the topographic caldera. Both the buried load and the caldera are elongated in the north–south direction. In the center of the buried load, the minimum thickness is 2.8 km for an olivine-dominated cumulate system or 3.9 km for a pyroxene-dominated system. The best terrestrial analog for this structure is the Bushveld Complex, an igneous cumulate body that is similar in size and thickness to the Syrtis Major structure. Assuming that the mean crustal density is 2600 kg m −3 due to impact brecciation, the elastic lithosphere at Syrtis Major was 10–15 km thick at the time when the topographic load was emplaced. This corresponds to a lithospheric thermal gradient of 28–52 K/km and a surface heat flux of 70–130 mW m −2 . Higher resolution gravity data, such as that which is planned for the 2005 Mars Reconnaissance Orbiter, will permit further refinement of the dimensions of this structure.

Radiometric age dating of the shergottite meteorites and cratering studies of lava flows in Tharsis and Elysium both demonstrate that volcanic activity has occurred on Mars in the geologically recent past. This implies that adiabatic decompression melting and upwelling convective flow in the mantle remains important on Mars at present. I present a series of numerical simulations of mantle convection and magma generation on Mars. These models test the effects of the total radioactive heating budget and of the partitioning of radioactivity between crust and mantle on the production of magma. In these models, melting is restricted to the heads of hot mantle plumes that rise from the core-mantle boundary, consistent with the spatially localized distribution of recent volcanism on Mars. For magma production to occur on present-day Mars, the minimum average radioactive heating rate in the martian mantle is 1.6 x 10^(-12) W/kg, which corresponds to 39% of the Wnke and Dreibus (1994) radioactivity abundance. If the mantle heating rate is lower than this, the mean mantle temperature is low, and the mantle plumes experience large amounts of cooling as they rise from the base of the mantle to the surface and are, thus, unable to melt. Models with mantle radioactive heating rates of 1.8 to 2.1 x 10^(-12) W/kg can satisfy both the present-day volcanic resurfacing rate on Mars and the typical melt fraction observed in the shergottites. This corresponds to 43-50% of the Wnke and Dreibus radioactivity remaining in the mantle, which is geochemically reasonable for a 50 km thick crust formed by about 10% partial melting. Plausible changes to either the assumed solidus temperature or to the assumed core-mantle boundary temperature would require a larger amount of mantle radioactivity to permit present-day magmatism. These heating rates are slightly higher than inferred for the nakhlite source region and significantly higher than inferred from depleted shergottites such as QUE 94201. The geophysical estimate of mantle radioactivity inferred here is a global average value, while values inferred from the martian meteorites are for particular points in the martian mantle. Evidently, the martian mantle has several isotopically distinct compositions, possibly including a radioactively enriched source that has not yet been sampled by the martian meteorites. The minimum mantle heating rate corresponds to a minimum thermal Rayleigh number of 2 x 10^6, implying that mantle convection remains moderately vigorous on present-day Mars. The basic convective pattern on Mars appears to have been stable for most of martian history, which has prevented the mantle flow from destroying the isotopic heterogeneity.

The Procellarum region is a broad area on the nearside of the Moon that is characterized by low elevations, thin crust, and high surface concentrations of the heat-producing elements uranium, thorium, and potassium. The region has been interpreted as an ancient impact basin approximately 3,200 kilometres in diameter, although supporting evidence at the surface would have been largely obscured as a result of the great antiquity and poor preservation of any diagnostic features. Here we use data from the Gravity Recovery and Interior Laboratory (GRAIL) mission to examine the subsurface structure of Procellarum. The Bouguer gravity anomalies and gravity gradients reveal a pattern of narrow linear anomalies that border Procellarum and are interpreted to be the frozen remnants of lava-filled rifts and the underlying feeder dykes that served as the magma plumbing system for much of the nearside mare volcanism. The discontinuous surface structures that were earlier interpreted as remnants of an impact basin rim are shown in GRAIL data to be a part of this continuous set of border structures in a quasi-rectangular pattern with angular intersections, contrary to the expected circular or elliptical shape of an impact basin. The spatial pattern of magmatic-tectonic structures bounding Procellarum is consistent with their formation in response to thermal stresses produced by the differential cooling of the province relative to its surroundings, coupled with magmatic activity driven by the greater-than-average heat flux in the region.

Measuring the seismic activity of Venus is critical to understanding its composition and interior dynamics. Because Venus has an average surface temperature of 462 °C and the challenge of providing cooling to multiple seismometers, a high temperature, wireless sensor using a wide bandgap semiconductor is an attractive option. This paper presents progress towards a seismometer sensor with wireless capabilities for Venus applications. A variation in inductance of a coil caused by a 1 cm movement of a ferrite probe held in the coil and attached to a balanced leaf-spring seismometer causes a variation of 80 MHz in the transmitted signal from the oscillator/sensor system at 420 °C, which correlates to a 10 kHz/mm sensitivity when the ferrite probe is located at the optimum location in the coil.