Titan impacts and escape

We report on hydrodynamic calculations of impacts of large (multi-kilometer) objects on Saturn’s moon Titan. We assess escape from Titan, and evaluate the hypothesis that escaping ejecta blackened the leading hemisphere of Iapetus and peppered the surface of Hyperion.We carried out two- and three-dimensional simulations of impactors ranging in size from 4 to 100 km diameter, impact velocities between 7 and 15 km s⁻¹, and impact angles from 0° to 75° from the vertical. We used the ZEUSMP2 hydrocode for the calculations. Simulations were made using three different geometries: three-dimensional Cartesian, two-dimensional axisymmetric spherical polar, and two-dimensional plane polar. Three-dimensional Cartesian geometry calculations were carried out over a limited domain (e.g. 240 km on a side for an impactor of size d_i = 10 km), and the results compared to ones with the same parameters done by Artemieva and Lunine (2005); in general the comparison was good. Being computationally less demanding, two-dimensional calculations were possible for much larger domains, covering global regions of the satellite (from 800 km below Titan’s surface to the exobase altitude 1700 km above the surface). Axisymmetric spherical polar calculations were carried out for vertical impacts. Two-dimensional plane-polar geometry calculations were made for both vertical and oblique impacts. In general, calculations among all three geometries gave consistent results.Our basic result is that the amount of escaping material is less than or approximately equal to the impactor mass even for the most favorable cases. Amounts of escaping material scaled most strongly as a function of velocity, with high-velocity impacts generating the largest amount, as expected. Dependence of the relative amount of escaping mass f_esc = m_esc/M_i on impactor diameter d_i was weak. Oblique impacts (impact angle θ_i > 45°) were more effective than vertical or near-vertical impacts; ratios of m_esc/M_i ∼ 1-2 were found in the simulations.     

Scaling of oblique impacts in frictional targets: Implications for crater size and formation mechanisms

Almost every meteorite impact occurs at an oblique angle of incidence, yet the effect of impact angle on crater size or formation mechanism is only poorly understood. This is, in large part, due to the difficulty of inferring impactor properties, such as size, velocity and trajectory, from observations of natural craters, and the expense and complexity of simulating oblique impacts using numerical models. Laboratory oblique impact experiments and previous numerical models have shown that the portion of the projectile’s kinetic energy that is involved in crater excavation decreases significantly with impact angle. However, a thorough quantification of planetary-scale oblique impact cratering does not exist and the effect of impact angle on crater size is not considered by current scaling laws. To address this gap in understanding, we developed iSALE-3D, a three-dimensional multi-rheology hydrocode, which is efficient enough to perform a large number of well-resolved oblique impact simulations within a reasonable time. Here we present the results of a comprehensive numerical study containing more than 200 three-dimensional hydrocode-simulations covering a broad range of projectile sizes, impact angles and friction coefficients. We show that existing scaling laws in principle describe oblique planetary-scale impact events at angles greater than 30° measured from horizontal. The displaced mass of a crater decreases with impact angle in a sinusoidal manner. However, our results indicate that the assumption that crater size scales with the vertical component of the impact velocity does not hold for materials with a friction coefficient significantly lower than 0.7 (sand). We found that increasing coefficients of friction result in smaller craters and a formation process more controlled by impactor momentum than by energy.     

Intra-crater glacial processes in central Utopia Planitia, Mars

We describe and interpret a series of previously unidentified glacial-like lobes (34-43°N; 107-125°E) that were discovered as part of a survey of large (D > 5 km) impact craters in Utopia Planitia, one of the great northern plains of Mars. The lobes have characteristics that are consistent with a glacial origin. Evidence includes curvilinearity of form, lineations and ridges, and surface textures that are thought to form by the sublimation of near-surface volatiles. The lobes display morphologies that range from wedge-shaped to near-circular to elongate. The flow directions are towards the northern walls in the case of craters with large single lobes, and in all directions in the case of the largest (D > 30 km) craters. Concentric crater fill is also interspersed within craters of our study region, with such craters having much higher filling rates than those with flow lobes. We suggest that the impact crater population in south-west Utopia Planitia demonstrates a spectrum of glacial modifications, from low levels of filling in the case of craters with elongate lobes at one extreme, to concentric crater fill in highly-filled craters at the other.     

Nonuniform cratering of the Moon and a revised crater chronology of the inner Solar System

We model the cratering of the Moon and terrestrial planets from the present knowledge of the orbital and size distribution of asteroids and comets in the inner Solar System, in order to refine the crater chronology method. Impact occurrences, locations, velocities and incidence angles are calculated semi-analytically, and scaling laws are used to convert impactor sizes into crater sizes. Our approach is generalizable to other moons or planets. The lunar cratering rate varies with both latitude and longitude: with respect to the global average, it is about 25% lower at (±65°N, 90°E) and larger by the same amount at the apex of motion (0°N, 90°W) for the present Earth-Moon separation. The measured size-frequency distributions of lunar craters are reconciled with the observed population of near-Earth objects under the assumption that craters smaller than a few kilometers in diameter form in a porous megaregolith. Varying depths of this megaregolith between the mare and highlands is a plausible partial explanation for differences in previously reported measured size-frequency distributions. We give a revised analytical relationship between the number of craters and the age of a lunar surface. For the inner planets, expected size-frequency crater distributions are calculated that account for differences in impact conditions, and the age of a few key geologic units is given. We estimate the Orientale and Caloris basins to be 3.73 Ga old, and the surface of Venus to be 240 Ma old. The terrestrial cratering record is consistent with the revised chronology and a constant impact rate over the last 400 Ma. Better knowledge of the orbital dynamics, crater scaling laws and megaregolith properties are needed to confidently assess the net uncertainty of the model ages that result from the combination of numerous steps, from the observation of asteroids to the formation of craters. Our model may be inaccurate for periods prior to 3.5 Ga because of a different impactor population, or for craters smaller than a few kilometers on Mars and Mercury, due to the presence of subsurface ice and to the abundance of large secondaries, respectively. Standard parameter values allow for the first time to naturally reproduce both the size distribution and absolute number of lunar craters up to 3.5 Ga ago, and give self-consistent estimates of the planetary cratering rates relative to the Moon.     

Variability in the small crater population on Callisto

Previous analyses of Galileo images showed the small (≈1 km and smaller) crater population on Callisto to be lower than had been expected (Moore, J.M. et al. [1999]. Icarus 140, 294-312; Bierhaus E.B. et al. [2000]. Lunar Planet. Sci. 31. Abstract #1996). In this paper we examine the small crater population using high-resolution imagery from Callisto flybys during Galileo orbits C3, C10, C21, and C30, including several C30 regions not previously analyzed. Our findings confirm that most small craters are depleted relative to a presumed equilibrium of R = 0.22, and we find that there is significant variability in the small crater counts. While some of the variability in the small crater population on Callisto can be attributed to secondary cratering, some variability also may be explained by resetting of portions of Callisto’s surface by larger impactors. This is expected where the differential size frequency distribution of the crater production population b < 3 (where b represents the exponent of a differential power-law crater-size distribution), such that large impacts affect a greater planetary surface area than smaller craters.     

High-resolution simulations of the impacts of asteroids into the venusian atmosphere III: further 3D models

We report on high-resolution three-dimensional calculations of oblique impacts into planetary atmospheres, specifically the atmosphere of Venus, extending the results of Korycansky et al. (2000, Icarus 146, 387-403; 2002, Icarus 157, 1-23). We have made calculations for impacts at 0°, 45°, and 60° from the vertical, different impactor velocities (10, 20, and 40 km s⁻¹), and different impactor masses and orientations. We present results for porous impactors using a simple model of porosity. We have investigated the sensitivity to initial conditions of the calculations [as a follow-up to the results found in Korycansky et al. (2002)] and resolution effects. For use in cratering calculations, we fit simple functions to the numerical results for mass and momentum that penetrate to a given altitude (column mass) and investigate the behavior of the fit coefficients as functions of impactor parameters such as mass, velocity, and impact angle. Generally speaking, the mass and momentum (and hence resulting crater diameters) depend primarily on impactor mass and mass of atmosphere encountered and weakly or not at all on other parameters such as impactor velocity, impact angle, or porosity. The column mass to which the last portion of the impactor penetrates is approximately equal to the mass of impactor at the top of the atmosphere before the impact takes place. Finally, we present the beginnings of a simplified but physically based model for the impactor and its fragments to reproduce the mass and momentum fluxes as a function of height during the impact.     

Impact cratering records of the mid-sized, icy saturnian satellites

Resolution of Voyager 1 and 2 images of the mid-sized, icy saturnian satellites was generally not much better than 1 km per line pair, except for a few, isolated higher resolution images. Therefore, analyses of impact crater distributions were generally limited to diameters (D) of tens of kilometers. Even with the limitation, however, these analyses demonstrated that studying impact crater distributions could expand understanding of the geology of the saturnian satellites and impact cratering in the outer Solar System. Thus to gain further insight into Saturn’s mid-sized satellites and impact cratering in the outer Solar System, we have compiled cratering records of these satellites using higher resolution CassiniISS images. Images from Cassini of the satellites range in resolution from tens m/pixel to hundreds m/pixel. These high-resolution images provide a look at the impact cratering records of these satellites never seen before, expanding the observable craters down to diameters of hundreds of meters. The diameters and locations of all observable craters are recorded for regions of Mimas, Tethys, Dione, Rhea, Iapetus, and Phoebe. These impact crater data are then analyzed and compared using cumulative, differential and relative (R) size-frequency distributions. Results indicate that the heavily cratered terrains on Rhea and Iapetus have similar distributions implying one common impactor population bombarded these two satellites. The distributions for Mimas and Dione, however, are different from Rhea and Iapetus, but are similar to one another, possibly implying another impactor population common to those two satellites. The difference between these two populations is a relative increase of craters with diameters between 10 and 30 km and a relative deficiency of craters with diameters between 30 and 80 km for Mimas and Dione compared with Rhea and Iapetus. This may support the result from Voyager images of two distinct impactor populations. One population was suggested to have a greater number of large impactors, most likely heliocentric comets (Saturn Population I in the Voyager literature), and the other a relative deficiency of large impactors and a greater number of small impactors, most likely planetocentric debris (Saturn Population II). Meanwhile, Tethys’ impact crater size-frequency distribution, which has some similarity to the distributions of Mimas, Dione, Rhea, and Iapetus, may be transitional between the two populations. Furthermore, when the impact crater distributions from these older cratered terrains are compared to younger ones like Dione’s smooth plains, the distributions have some similarities and differences. Therefore, it is uncertain whether the size-frequency distribution of the impactor population(s) changed over time. Finally, we find that Phoebe has a unique impact crater distribution. Phoebe appears to be lacking craters in a narrow diameter range around 1 km. The explanation for this confined “dip” at D = 1 km is not yet clear, but may have something to do with the interaction of Saturn’s irregular satellites or the capture of Phoebe.     

Impact megadomes and the origin of the martian crustal dichotomy

We show that a sufficiently energetic impact can generate a melt volume which, after isostatic adjustment and differentiation, forms a spherical cap of crust with underlying depleted mantle. Depending on impact energy and initial crustal thickness, a basin may be retained or impact induced crust may be topographically elevated. Retention of a martian lowland scale impact basin at impact energies ∼3 × 10²⁸-3 × 10²⁹ J requires an initial crustal thickness greater than 10 km. Formation of impact induced crust with size comparable to the martian highlands requires a larger impact energy, ∼1-3 × 10³⁰ J, and initial crustal thickness <20 km. Furthermore, we show that the boundary of impact induced crust can be elliptical due to a spatially asymmetric impact melt volume caused by an oblique impact. We suggest the term “impact megadome” for topographically elevated, impact induced crust and propose that processes involved in megadome formation may play an important role in the origin of the martian crustal dichotomy.     

Impacts onto H2O ice: Scaling laws for melting, vaporization, excavation, and final crater size

Shock-induced melting and vaporization of H₂O ice during planetary impact events are widespread phenomena. Here, we investigate the mass of shock-produced liquid water remaining within impact craters for the wide range of impact conditions and target properties encountered in the Solar System. Using the CTH shock physics code and the new 5-phase model equation of state for H₂O, we calculate the shock pressure field generated by an impact and fit scaling laws for melting and vaporization as a function of projectile mass, impact velocity, impact angle, initial temperature, and porosity. Melt production nearly scales with impact energy, and natural variations in impact parameters result in only a factor of two change in the predicted mass of melt. A fit to the π-scaling law for the transient cavity and transient-to-final crater diameter scaling are determined from recent simulations of the entire cratering process in ice. Combining melt production with π-scaling and the modified Maxwell Z-model for excavation, less than half of the melt is ejected during formation of the transient crater. For impact energies less than about 2 × 10²⁰ J and impact velocities less than about 5 km s⁻¹, the remaining melt lines the final crater floor. However, for larger impact energies and higher impact velocities, the phenomenon of discontinuous excavation in H₂O ice concentrates the impact melt into a small plug in the center of the crater floor.     

Impact penetration of Europa's ice crust as a mechanism for formation of chaos terrain

Abstract— Ice thickness estimates and impactor dynamics indicate that some impacts must breach Europa's ice crust; and outcomes of impact experiments using ice‐over‐water targets range from simple craters to chaos‐like destroyed zones, depending on impact energy and ice competence. First‐order impacts‐into thick ice or at low impact energy‐produce craters. Second‐order impacts punch through the ice, making holes that resemble raft‐free chaos areas. Third‐order impacts‐into thinnest ice or at highest energy‐produce large irregular raft‐filled zones similar to platy chaos. Other evidence for an impact origin for chaos areas comes from the size‐frequency distribution of chaos+craters on Europa, which matches the impact production functions of Ganymede and Callisto; and from small craters around the large chaos area Thera Macula, which decrease in average size and density per unit area as a function of distance from Thera's center. There are no tiny chaos areas and no craters >50 km diameter. This suggests that small impactors never penetrate, whereas large ones (ÜberPenetrators: >2.5 km diameter at average impact velocity) always do. Existence of both craters and chaos areas in the size range 2–40 km diameter points to spatial/temporal variation in crust thickness. But in this size range, craters are progressively outnumbered by chaos areas at larger diameters, suggesting that probability of penetration increases with increasing scale of impact. If chaos areas do represent impact sites, then Europa's surface is older than previously thought. The recalculated resurfacing age is 480 (‐302/+960) Ma: greater than prior estimates, but still very young by solar system standards.     

Numerical simulations of very large impacts on the Earth

Vertical impacts on the Earth of asteroids 500-3000 km in diameter at 15 km/s have been numerically modelled using the hydrodynamic SOVA code. This code has been modified for the spherical system of coordinates well suited for simulations of very large impacts when the entire Earth is involved in motion. The simulations include cratering process, upward motion of deep mantle layers, fall of ejecta on the Earth, escape of matter to space, and formation of rock vapour atmospheres. The calculations were made for the period preceding disappearance of rock vapour atmospheres caused by radiation several years after the largest impacts. For very large vertical impacts at 15 km/s, escaping masses proved to be negligibly small. Quantities of kinetic, internal, potential, and radiated away energies are obtained as functions of time and space. After the impacts, a global layer of condensed ejecta covers the whole of the Earth's surface and the ejecta energy is sufficient to vaporise an ocean 3 km deep. The mass of rock vapour atmosphere is 10-23% of the impactor mass. This atmosphere has a greater mass than the water atmosphere if impactor is 2000 km in diameter or larger.     

The transition from complex crater to peak-ring basin on the Moon: New observations from the Lunar Orbiter Laser Altimeter (LOLA) instrument

Impact craters on planetary bodies transition with increasing size from simple, to complex, to peak-ring basins and finally to multi-ring basins. Important to understanding the relationship between complex craters with central peaks and multi-ring basins is the analysis of protobasins (exhibiting a rim crest and interior ring plus a central peak) and peak-ring basins (exhibiting a rim crest and an interior ring). New data have permitted improved portrayal and classification of these transitional features on the Moon. We used new 128 pixel/degree gridded topographic data from the Lunar Orbiter Laser Altimeter (LOLA) instrument onboard the Lunar Reconnaissance Orbiter, combined with image mosaics, to conduct a survey of craters >50 km in diameter on the Moon and to update the existing catalogs of lunar peak-ring basins and protobasins. Our updated catalog includes 17 peak-ring basins (rim-crest diameters range from 207 km to 582 km, geometric mean = 343 km) and 3 protobasins (137-170 km, geometric mean = 157 km). Several basins inferred to be multi-ring basins in prior studies (Apollo, Moscoviense, Grimaldi, Freundlich-Sharonov, Coulomb-Sarton, and Korolev) are now classified as peak-ring basins due to their similarities with lunar peak-ring basin morphologies and absence of definitive topographic ring structures greater than two in number. We also include in our catalog 23 craters exhibiting small ring-like clusters of peaks (50-205 km, geometric mean = 81 km); one (Humboldt) exhibits a rim-crest diameter and an interior morphology that may be uniquely transitional to the process of forming peak rings. A power-law fit to ring diameters (D_ring) and rim-crest diameters (D_r) of peak-ring basins on the Moon [D_ring = 0.14 ± 0.10(D_r)^1.21±0.13] reveals a trend that is very similar to a power-law fit to peak-ring basin diameters on Mercury [D_ring = 0.25 ± 0.14(D_rim)^1.13±0.10] [Baker, D.M.H. et al. [2011]. Planet. Space Sci., in press]. Plots of ring/rim-crest ratios versus rim-crest diameters for peak-ring basins and protobasins on the Moon also reveal a continuous, nonlinear trend that is similar to trends observed for Mercury and Venus and suggest that protobasins and peak-ring basins are parts of a continuum of basin morphologies. The surface density of peak-ring basins on the Moon (4.5 × 10⁻⁷ per km²) is a factor of two less than Mercury (9.9 × 10⁻⁷ per km²), which may be a function of their widely different mean impact velocities (19.4 km/s and 42.5 km/s, respectively) and differences in peak-ring basin onset diameters. New calculations of the onset diameter for peak-ring basins on the Moon and the terrestrial planets re-affirm previous analyses that the Moon has the largest onset diameter for peak-ring basins in the inner Solar System. Comparisons of the predictions of models for the formation of peak-ring basins with the characteristics of the new basin catalog for the Moon suggest that formation and modification of an interior melt cavity and nonlinear scaling of impact melt volume with crater diameter provide important controls on the development of peak rings. In particular, a power-law model of growth of an interior melt cavity with increasing crater diameter is consistent with power-law fits to the peak-ring basin data for the Moon and Mercury. We suggest that the relationship between the depth of melting and depth of the transient cavity offers a plausible control on the onset diameter and subsequent development of peak-ring basins and also multi-ring basins, which is consistent with both planetary gravitational acceleration and mean impact velocity being important in determining the onset of basin morphological forms on the terrestrial planets.     

A giant crater on 90 Antiope?

Mutual event observations between the two components of 90 Antiope were carried out in 2007-2008. The pole position was refined to λ₀ = 199.5 ± 0.5° and β₀ = 39.8 ± 5° in J2000 ecliptic coordinates, leaving intact the physical solution for the components, assimilated to two perfect Roche ellipsoids, and derived after the 2005 mutual event season (Descamps, P., Marchis, F., Michalowski, T., Vachier, F., Colas, F., Berthier, J., Assafin, M., Dunckel, P.B., Polinska, M., Pych, W., Hestroffer, D., Miller, K., Vieira-Martins, R., Birlan, M., Teng-Chuen-Yu, J.-P., Peyrot, A., Payet, B., Dorseuil, J., Léonie, Y., Dijoux, T., 2007. Figure of the double Asteroid 90 Antiope from AO and lightcurves observations. Icarus 187, 482-499). Furthermore, a large-scale geological depression, located on one of the components, was introduced to better match the observed lightcurves. This vast geological feature of about 68 km in diameter, which could be postulated as a bowl-shaped impact crater, is indeed responsible of the photometric asymmetries seen on the “shoulders” of the lightcurves. The bulk density was then recomputed to 1.28 ± 0.04 g cm⁻³ to take into account this large-scale non-convexity. This giant crater could be the aftermath of a tremendous collision of a 100-km sized proto-Antiope with another Themis family member. This statement is supported by the fact that Antiope is sufficiently porous (∼50%) to survive such an impact without being wholly destroyed. This violent shock would have then imparted enough angular momentum for fissioning of proto-Antiope into two equisized bodies. We calculated that the impactor must have a diameter greater than ∼17 km, for an impact velocity ranging between 1 and 4 km/s. With such a projectile, this event has a substantial 50% probability to have occurred over the age of the Themis family.     

Linking the collisional history of the main asteroid belt to its dynamical excitation and depletion

The main belt is believed to have originally contained an Earth mass or more of material, enough to allow the asteroids to accrete on relatively short timescales. The present-day main belt, however, only contains ∼5×¹⁰⁻⁴ Earth masses. Numerical simulations suggest that this mass loss can be explained by the dynamical depletion of main belt material via gravitational perturbations from planetary embryos and a newly-formed Jupiter. To explore this scenario, we combined dynamical results from Petit et al. [Petit, J. Morbidelli, A., Chambers, J., 2001. The primordial excitation and clearing of the asteroid belt. Icarus 153, 338-347] with a collisional evolution code capable of tracking how the main belt undergoes comminution and dynamical depletion over 4.6 Gyr [Bottke, W.F., Durda, D., Nesvorny, D., Jedicke, R., Morbidelli, A., Vokrouhlický, D., Levison, H., 2005. The fossilized size distribution of the main asteroid belt. Icarus 175, 111-140]. Our results were constrained by the main belt's size-frequency distribution, the number of asteroid families produced by disruption events from diameter D>100 km parent bodies over the last 3-4 Gyr, the presence of a single large impact crater on Vesta's intact basaltic crust, and the relatively constant lunar and terrestrial impactor flux over the last 3 Gyr. We used our model to set limits on the initial size of the main belt as well as Jupiter's formation time. We find the most likely formation time for Jupiter was 3.3±2.6 Myr after the onset of fragmentation in the main belt. These results are consistent with the estimated mean disk lifetime of 3 Myr predicted by Haisch et al. [Haisch, K.E., Lada, E.A., Lada, C.J., 2001. Disk frequencies and lifetimes in young clusters. Astrophys. J. 553, L153-L156]. The post-accretion main belt population, in the form of diameter D?1000 km planetesimals, was likely to have been 160±40 times the current main belt's mass. This corresponds to 0.06-0.1 Earth masses, only a small fraction of the total mass thought to have existed in the main belt zone during planet formation. The remaining mass was most likely taken up by planetary embryos formed in the same region. Our results suggest that numerous D>200 km planetesimals disrupted early in Solar System history, but only a small fraction of their fragments survived the dynamical depletion event described above. We believe this may explain the limited presence of iron-rich M-type, olivine-rich A-type, and non-Vesta V-type asteroids in the main belt today. The collisional lifetimes determined for main belt asteroids agree with the cosmic ray exposure ages of stony meteorites and are consistent with the limited collisional evolution detected among large Koronis family members. Using the same model, we investigated the near-Earth object (NEO) population. We show the shape of the NEO size distribution is a reflection of the main belt population, with main belt asteroids driven to resonances by Yarkovsky thermal forces. We used our model of the NEO population over the last 3 Gyr, which is consistent with the current population determined by telescopic and satellite data, to explore whether the majority of small craters (D<0.1-1 km) formed on Mercury, the Moon, and Mars were produced by primary impacts or by secondary impacts generated by ejecta from large craters. Our results suggest that most small craters formed on these worlds were a by-product of secondary rather than primary impacts.     

The effect of the oceans on the terrestrial crater size‐frequency distribution: Insight from numerical modeling

Abstract— On Earth, oceanic impacts are twice as likely to occur as continental impacts, yet the effect of the oceans has not been previously considered when estimating the terrestrial crater size‐frequency distribution. Despite recent progress in understanding the qualitative and quantitative effect of a water layer on the impact process through novel laboratory experiments, detailed numerical modeling, and interpretation of geological and geophysical data, no definitive relationship between impactor properties, water depth, and final crater diameter exists. In this paper, we determine the relationship between final (and transient) crater diameter and the ratio of water depth to impactor diameter using the results of numerical impact models. This relationship applies for normal incidence impacts of stoney asteroids into water‐covered, crystalline oceanic crust at a velocity of 15 km s^?1. We use these relationships to construct the first estimates of terrestrial crater size‐frequency distributions (over the last 100 million years) that take into account the depth‐area distribution of oceans on Earth. We find that the oceans reduce the number of craters smaller than 1 km in diameter by about two‐thirds, the number of craters ?30 km in diameter by about one‐third, and that for craters larger than ?100 km in diameter, the oceans have little effect. Above a diameter of ?12 km, more craters occur on the ocean floor than on land; below this diameter more craters form on land than in the oceans. We also estimate that there have been in the region of 150 impact events in the last 100 million years that formed an impact‐related resurge feature, or disturbance on the seafloor, instead of a crater.     

Volatile retention from cometary impacts on the Moon

Impacts of comets and asteroids play an important role in volatile delivery on the Moon. We use a novel method for tracking vapor masses that reach escape velocity in hydrocode simulations of cometary impacts to explore the effects of volatile retention. We model impacts on the Moon to find the mass of vapor plume gravitationally trapped on the Moon as a function of impact velocity. We apply this result to the impactor velocity distribution and find that the total impactor mass retained on the Moon is approximately 6.5% of the impactor mass flux. Making reasonable assumptions about water content of comets and the comet size-frequency distribution, we derive a water flux for the Moon. After accounting for migration and stability of water ice at the poles, we estimate a total 1.3×10⁸-4.3×10⁹ metric tons of water is delivered to the Moon and remains stable at the poles over 1 Ga. A factor of 30 uncertainty in the estimated cometary impact flux is primarily responsible for this large range of values. The calculated mass of water is sufficient to account for the neutron fluxes poleward of 75° observed by Lunar Prospector. A similar analysis for water delivery to the Moon via asteroid impacts shows that asteroids provide six times more water mass via impacts than comets.     

A geomorphic analysis of Hale crater, Mars: The effects of impact into ice-rich crust

Hale crater, a 125 × 150 km impact crater located near the intersection of Uzboi Vallis and the northern rim of Argyre basin at 35.7°S, 323.6°E, is surrounded by channels that radiate from, incise, and transport material within Hale’s ejecta. The spatial and temporal relationship between the channels and Hale’s ejecta strongly suggests the impact event created or modified the channels and emplaced fluidized debris flow lobes over an extensive area (>200,000 km²). We estimate ∼10¹⁰ m³ of liquid water was required to form some of Hale’s smaller channels, a volume we propose was supplied by subsurface ice melted and mobilized by the Hale-forming impact. If 10% of the subsurface volume was ice, based on a conservative porosity estimate for the upper martian crust, 10¹² m³ of liquid water could have been present in the ejecta. We determine a crater-retention age of 1 Ga inside the primary cavity, providing a minimum age for Hale and a time at which we propose the subsurface was volatile-rich. Hale crater demonstrates the important role impacts may play in supplying liquid water to the martian surface: they are capable of producing fluvially-modified terrains that may be analogous to some landforms of Noachian Mars.     

Topographic modeling of Phoebe using Cassini images

High-resolution Cassini stereo images of Saturn's moon Phoebe have been used to derive a regional digital terrain model (DTM) and an orthoimage mosaic of the surface. For DTM-control a network of 130 points measured in 14 images (70-390 m/pixel resolution) was established which was simultaneously used to determine the orientation of the spin-axis. The J2000 spin-axis was found at Dec=78.0°±0.1° and RA=356.6°±0.3°, substantially different from the former Voyager solution. The control points yield a mean figure radius of 107.2 km with RMS residuals of 6.2 km demonstrating the irregular shape of this body. The DTM was computed from densely spaced conjugate image points determined by methods of digital image correlation. It has a horizontal resolution of 1-2 km and vertical accuracies in the range 50-100 m. It is limited in coverage, but higher in resolution than the previously derived global shape model of Phoebe [Porco et al., 2005. Cassini imaging science: initial results on Phoebe and Iapetus. Science 307, 1237-1242] and allows us to study the morphology of the surface in more detail. There is evidence for unconsolidated material from a steep and smooth slope at the rim of a 100 km impact feature. There are several conically shaped craters on Phoebe, which may hint at highly porous and low compaction material on the surface.     

Ejecta size-velocity relation derived from the distribution of the secondary craters of kilometer-sized craters on Mars

The relation between the size and velocity of impact crater ejecta has been studied by both laboratory experiments and numerical modeling. An alternative method, used here, is to analyze the record of past impact events, such as the distribution of secondary craters on planetary surfaces, as described by Vickery (Icarus 67 (1986) 224; Geophys. Res. Lett. 14 (1987) 726). We first applied the method to lunar images taken by the CLEMENTINE mission, which revealed that the size-velocity relations of ejecta from craters 32 and 40 km in diameter were similar to those derived by Vickery for a crater 39 km in diameter. Next, we studied the distribution of small craters in the vicinity of kilometer-sized craters on three images from the Mars Orbiter Camera (MOC) on board the Mars Global Surveyor (MGS). If these small craters are assumed to be secondaries ejected from the kilometer-sized crater in each image, the ejection velocities are of hundreds of meters per second. These data fill a gap between the previous results of Vickery and those of laboratory studies.