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An exceptional deposit covering an area of ∼7700 km 2 , displaying morphology indicative of an originally fluid material, occurs near 42.2°N, 167.4°E on the lunar farside. The material occurs as smooth, flat deposits (here termed “ponds”) on the bottoms of many craters and in other topographic depressions, as well as a veneer across the majority of the region. Within this area, the ponded deposits and widespread-veneer have an estimated volume of ∼8 km 3 ; the veneer constitutes the great majority of this volume. This material appears to have flowed downslope across the surface, collecting in flat-surface accumulations. The surfaces of the ponds and veneer are only lightly cratered, indicating a young (i.e., late Copernican) age. Four possible modes of origin are investigated: basin ejecta, pyroclastic volcanism, effusive volcanism, and ballistically emplaced impact-melt. Volcanism and basin ejecta appear to be inconsistent with the observed morphology: an implausible number of vents are required for volcanism and the morphological properties do not resemble basin ejecta. We suggest that ballistically emplaced impact melt is most consistent with the observations. Possible source craters for impact melt, based on minimum required size (>20 km diameter) and age, are at least 250 km distant and cannot be definitively tied to the pond deposits. This discovery places important new constraints on our knowledge of the distribution of impact melt relative to the parent crater.

International audience; On February 14, 2011 Stardust-NExT (SN) flew by Comet Tempel 1, the target of the Deep Impact (DI) mission in 2005, obtaining dust measurements and high-resolution images of areas surrounding the 2005 impact site, and extending image coverage to almost two thirds of the nucleus surface. The nucleus has an average radius of 2.83 +/- 0.1 km and a uniform geometric albedo of about 6% at visible wavelengths. Local elevation differences on the nucleus reach up to 830 m. At the time of encounter the spin rate was 213 degrees per day (period = 40.6 h) and the comet was producing some 130 kg of dust per second. Some 30% of the nucleus is covered by smooth flow-like deposits and related materials, restricted to gravitational lows. This distribution is consistent with the view that the smooth areas represent material erupted from the subsurface and date from a time after the nucleus achieved its current shape. It is possible that some of these eruptions occurred after 1609 when the comet's perihelion distance decreased from 3.5 AU to the current 1.5 AU. Much of the surface displays evidence of layering: some related to the smooth flows and some possibly dating back to the accretion of the nucleus. Pitted terrain covers approximately half the nucleus surface. The pits range up to 850 m in diameter. Due to their large number, they are unlikely to be impact scars: rather they probably result from volatile outbursts and sublimational erosion. The DI impact site shows a subdued depression some 50 m in diameter implying surface properties similar to those of dry, loose snow. It is possible that the 50-m depression is all that remains of an initially larger crater. In the region of overlapping DI and SN coverage most of the surface remained unchanged between 2005 and 2011 in albedo, photometric properties and morphology. Significant changes took place only along the edges of a prominent smooth flow estimated to be 10-15 m thick, the margins of which receded in places by up to 50 m. Coma and jet activity were lower in 2011 than in 2005. Most of the jets observed during the SN flyby can be traced back to an apparently eroding terraced scarp. The dust instruments detected bursts of impacts consistent with a process by which larger aggregates of material emitted from the nucleus subsequently fragment into smaller particles within the coma. (c) 2012 Elsevier Inc. All rights reserved.

We consider the origin of ∼380 quasi-circular depressions (pits) seen to be distributed in a broad band across the surface of 9P/Tempel 1 and show that possibly ∼96% may be due to outburst activity. Of the rest, 4 kg and find that the resulting pit should have a diameter in the range 10–30 m. Published locations of mini-outbursts are revised to account for changes in the nucleus shape, rotation rate, and rotation pole that have resulted from observations made during the Stardust-NExT mission. Both of these locations are found to fall in, or on the edge of, the band of pits that encircles the nucleus. We have identified features in high-resolution images near one of these locations as the possible places of origin of the mini-outbursts. These features show close packing of multiple pits in the appropriate diameter range. We consider the distribution of pit diameters and show that the largest pits follow a power–law with exponent −2.24 ± 0.09. Using the June 14, 2005, mini-outburst and the Deep Impact crater to provide a calibration, we establish empirical relationships between pit diameter, D , the total outburst energy, E , and the visual magnitude change, Δ m abs , which is the visual amplitude of the outburst referenced to a standard initial brightness. We find Log 10 D ∼ 0.107(±0.004)Δ m abs + 1.3(±0.4) and Log 10 E ∼ 0.32(±0.01)Δ m abs + 10.1(±1.2) where the uncertainties represent the range of values for the coefficient rather than formal error. We apply these approximate relationships to the mega -outburst on 17P/Holmes and predict that it left a pit-like scar on the surface with a diameter in the range 160–1300 m, that the total energy released was in the range 7 × 10 12 –3 × 10 15 J, and that between 6 × 10 7 and 1.3 × 10 11 kg of material was ejected from the surface. While these predictions are crude they encompass, particularly near the upper end of the range, the results on kinetic energy release and mass loss found by Reach et al. (Reach, W.T., Vaubaillon, J., Lisse, C.M., Holloway, M., Rho, J. [2010]. Icarus 208, 276–292) based on IR observations of 17P.

Random bombardment by comets, asteroids and associated fragments form and alter the lunar regolith and other rocky surfaces. The accumulation of impact craters over time is of fundamental use in evaluating the relative ages of geologic units. Crater counts and radiometric ages from returned samples provide constraints with which to derive absolute model ages for unsampled units on the Moon and other Solar System objects. However, although studies of existing craters and returned samples offer insight into the process of crater formation and the past cratering rate, questions still remain about the present rate of crater production, the effect of early-stage jetting during impacts and the influence that distal ejecta have on the regolith. Here we use Lunar Reconnaissance Orbiter Camera (LROC) Narrow Angle Camera (NAC) temporal ('before and after') image pairs to quantify the contemporary rate of crater production on the Moon, to reveal previously unknown details of impact-induced jetting, and to identify a secondary impact process that is rapidly churning the regolith. From this temporal dataset, we detected 222 new impact craters and found 33 per cent more craters (with diameters of at least ten metres) than predicted by the standard Neukum production and chronology functions for the Moon. We identified broad reflectance zones associated with the new craters that we interpret as evidence of a surface-bound jetting process. We also observe a secondary cratering process that we estimate churns the top two centimetres of regolith on a timescale of 81,000 years-more than a hundred times faster than previous models estimated from meteoritic impacts (ten million years).

Lunar Lobate Scarps Revealed Lunar lobate scarps are relatively small-scale landforms that are thought to be formed by tectonic thrust faulting. Previously, lunar lobate scarps could only be identified clearly in high-resolution Apollo Panoramic Camera images confined to the lunar equatorial zone. Now, an analysis by Watters et al. (p. 936 ) of images returned by the Lunar Reconnaissance Orbiter Camera reveals 14 previously unknown lobate scarps and shows that lunar lobate scarps may be globally distributed. Their appearance suggests that lunar scarps are relatively young landforms (less than 1 Ga), possibly formed during a recent episode of global lunar radial contraction.

Deep Impact images of the nucleus of Comet Tempel 1 reveal pervasive layering, possible impact craters, flows with smooth upper surfaces, and erosional stripping of material. There are at least 3 layers 50–200 m thick that appear to extend deep into the nucleus, and several layers 1–20 m thick that parallel the surface and are being eroded laterally. Circular depressions show geographical variation in their forms and suggest differences in erosion rates or style over scales >1 km. The stratigraphic arrangement of these features suggests that the comet experienced substantial periods of little erosion. Smooth surfaces trending downslope suggest some form of eruption of materials from this highly porous object. The Deep Impact images show that the nucleus of Tempel 1 cannot be modeled simply as either an onion-layer or rubble pile structure.

Deep Impact collided with comet Tempel 1, excavating a crater controlled by gravity. The comet's outer layer is composed of 1- to 100-micrometer fine particles with negligible strength (1000 kelvins). A large increase in organic material occurred during and after the event, with smaller changes in carbon dioxide relative to water. On approach, the spacecraft observed frequent natural outbursts, a mean radius of 3.0 ± 0.1 kilometers, smooth and rough terrain, scarps, and impact craters. A thermal map indicates a surface in equilibrium with sunlight.

We present HST Faint Object Camera observations of the asteroid 1 Ceres at near-, mid-, and far-UV wavelengths (lambda = 3636, 2795, and 1621 A, respectively) obtained on 1995 June 25. The disk of Ceres is well-resolved for the first time, at a scale of ~50km. We report the detection of a large, ~250km diameter surface feature for which we propose the name ``Piazzi''; however it is presently uncertain if this feature is due to a crater, albedo variegation, or other effect. From limb fits to the images, we obtain semi-major and semi-minor axes of R_1=484.8+/-5.1km and R_2=466.4+/-5.9km, respectively, for the illumination-corrected projected ellipsoid. Although albedo features are seen, they do not allow for a definitive determination of the rotation or pole positions of Ceres, particularly because of the sparse sampling (two epochs) of the 9 hour rotation period. From full-disk integrated albedo measurements, we find that Ceres has a red spectral slope from the mid- to near-UV, and a significant blue slope shortward of the mid-UV. In spite of the presence of Piazzi, we detect no significant global differences in the integrated albedo as a function of rotational phase for the two epochs of data we obtained. From Minnaert surface fits to the near- and mid-UV images, we find an unusually large Minnaert parameter of k~0.9, suggesting a more Lambertian than lunar-like surface., Comment: 9 pages, 4 figures, in AASTeX v5.02 document class with "emulateapj5" and "onecolfloat" styles (onecolfloat.sty is included). To appear in the Astronomical Journal (January 2002). Also available at http://www.boulder.swri.edu/~joel/papers.html

Preliminary measurements of craters and boulders have been made in various locations on Eros from images acquired during the first nine months of NEAR Shoemaker's orbital mission, including the October 2000 low altitude flyover. (We offer some very preliminary, qualitative analysis of later LAF images and very high-resolution images obtained during NEAR's landing on 12 February 2001). Craters on Eros >100 m diameter closely resemble the saturated crater population of Ida; Eros is more heavily cratered than Gaspra but lacks the saturated giant craters of Mathilde. These craters and the other large-scale geological features were formed over a duration of very roughly 2 Gyr while Eros was in the main asteroid belt, between the time when its parent body was disrupted and Eros was injected into an Earth-approaching orbit (probably tens of Myr ago). Saturation equilibrium had been expected to shape Eros' crater population down to very small sizes, as on the lunar maria. However, craters

NEAR Shoemaker imaging of the asteroid Eros has revealed a variety of surface forms that probably reflect underlying structures. They range from global to local in scale. The most prominent positive topographic feature is a long ridge, Rahe Dorsum, which extends over 18 km and has the morphology of a thrust fault. Two large, ancient trough systems are also present. Grooves are ubiquitous, commonly occur in suborthogonal sets, and have linear or scalloped outlines. The grooves on Eros are similar in morphology and scale to those on other small bodies such as Gaspra, Ida, and Phobos and are likely to have a similar mode of origin, namely drainage of regolith into underlying fractures. Models that use groove dimensions as indicators of regolith thickness imply depths of up to several tens of meters. Many craters have been subjected to structural control, resulting in squared outlines. We find no evidence that specific groove sets are associated with individual impact events; however, evidence exists that preexisting structures have been reactivated by later impact activity. The large scale of Rahe Dorsum and many grooves on Eros's surface suggests that the asteroid is a largely coherent, but fractured body.

The NEAR-Shoemaker spacecraft was designed to provide a comprehensive characterization of the S-type asteroid 433 Eros (refs 1,2,3), an irregularly shaped body with approximate dimensions of 34 x 13 x 13 km. Following the completion of its year-long investigation, the mission was terminated with a controlled descent to its surface, in order to provide extremely high resolution images. Here we report the results of the descent on 12 February 2001, during which 70 images were obtained. The landing area is marked by a paucity of small craters and an abundance of 'ejecta blocks'. The properties and distribution of ejecta blocks are discussed in a companion paper. The last sequence of images reveals a transition from the blocky surface to a smooth area, which we interpret as a 'pond'. Properties of the 'ponds' are discussed in a second companion paper. The closest image, from an altitude of 129 m, shows the interior of a 100-m-diameter crater at 1-cm resolution.

During the Near Earth Asteroid Rendezvous (NEAR)–Shoemaker's low-altitude flyover of asteroid 433 Eros, observations by the NEAR Laser Rangefinder (NLR) have helped to characterize small-scale surface features. On scales from meters to hundreds of meters, the surface has a fractal structure with roughness dominated by blocks, structural features, and walls of small craters. This fractal structure suggests that a single process, possibly impacts, dominates surface morphology on these scales.

Eros is a very elongated (34 kilometers by 11 kilometers by 11 kilometers) asteroid, most of the surface of which is saturated with craters smaller than 1 kilometer in diameter. The largest crater is 5.5 kilometers across, but there is a 10-kilometer saddle-like depression with attributes of a large degraded crater. Surface lineations, both grooves and ridges, are prominent on Eros; some probably exploit planes of weakness produced by collisions on Eros and/or its parent body. Ejecta blocks (30 to 100 meters across) are abundant but not uniformly distributed over the surface. Albedo variations are restricted to the inner walls of certain craters and may be related to downslope movement of regolith. On scales of 200 meters to 1 kilometer, Eros is more bland in terms of color variations than Gaspra or Ida. Spectra (800 to 2500 nanometers) are consistent with an ordinary chondrite composition for which the measured mean density of 2.67 ± 0.1 grams per cubic centimeter implies internal porosities ranging from about 10 to 30 percent.

We determined the mass of asteroid 433 Eros, its lower order gravitational harmonics, and rotation state, using ground-based Doppler and range tracking of the Near Earth Asteroid Rendezvous (NEAR)–Shoemaker spacecraft and images of the asteroid's surface landmarks. The mass of Eros is (6.687 ± 0.003) × 10 18 grams, which, coupled with our volume estimate, implies a bulk density of 2.67 ± 0.03 grams per cubic centimeter. The asteroid appears to have a uniform density distribution. The right ascension and declination of the rotation pole are 11.37 ± 0.05 and 17.22 ± 0.05 degrees, respectively, and at least over the short term, the rotation state of Eros is stable with no measurable free precession of the spin pole. Escape velocities on the surface vary from 3.1 to 17.2 meters per second. The dynamical environment of Eros suggests that it is covered with regolith and that one might expect material transport toward the deepest potential wells in the saddle and 5.5-kilometer crater regions.

Four flybys of Phobos by the Mars Global Surveyor spacecraft in 1998 allowed imaging at resolutions (∼2–7 m) better than obtained for any other satellite or asteroid except Earth's Moon. The images show the interior and vicinity of the large crater Stickney. There are ∼2000 identifiable ejecta blocks, the largest about 85 m across. The great majority of these blocks come from Stickney as part of low-velocity ejecta spread eastward by the influence of Phobos's rapid rotation. They appear to have experienced only a few meters, at most, of burial or modification since emplacement. The images also show materials of different albedos within Stickney that mark downslope motion of regolith. The number of craters on the steep slopes of Stickney that are subject to the downslope motion suggests that 4 km) have been degraded to the point that they are no longer recognizable.

As the planet's principal cold traps, the martian polar regions have accumulated extensive mantles of ice and dust that cover individual areas of approximately 10(6) km2 and total as much as 3-4 km thick. From the scarcity of superposed craters on their surface, these layered deposits are thought to be comparatively young--preserving a record of the seasonal and climatic cycling of atmospheric CO2, H2O, and dust over the past approximately 10(5)-10(8) years. For this reason, the martian polar deposits may serve as a Rosetta Stone for understanding the geologic and climatic history of the planet--documenting variations in insolation (due to quasiperiodic oscillations in the planet's obliquity and orbital elements), volatile mass balance, atmospheric composition, dust storm activity, volcanic eruptions, large impacts, catastrophic floods, solar luminosity, supernovae, and perhaps even a record of microbial life. Beyond their scientific value, the polar regions may soon prove important for another reason--providing a valuable and accessible reservoir of water to support the long-term human exploration of Mars. In this paper we assess the current state of Mars polar research, identify the key questions that motivate the exploration of the polar regions, discuss the extent to which current missions will address these questions, and speculate about what additional capabilities and investigations may be required to address the issues that remain outstanding.

Spacecraft imaging data allow an accounting of the occurrence and characteristics of the largest craters on small asteroids and satellites. Data covering Phobos, Deimos, Gaspra, Ida, Mathilde, Vesta, Amalthea, Thebe, Janus, Epimetheus, Hyperion, and Proteus show that ∼50% of rocky objects support craters with diameters of 1 object mean radius (R m ), and the great majority have more than two craters with D>0.5 R m . The morphology of these craters is consistent with gravity (g −1 ) scaling from craters on larger objects. The amounts of visible peripheral damage, such as fracturing caused by these relatively large impacts, are usually small; the best possible example (Phobos) may have been assisted by tidal forces. The occurrence and characteristics of these craters supports modeling results that predict formation in gravity-controlled regimes with very large relative-sized craters formed before global fragmentation occurs.

The Near Earth Asteroid Rendezvous (NEAR) spacecraft images of C-type asteroid Mathilde show it to have a mean radius of 26.5±1.3 km. The asteroid is geologically distinguished by its low mean density (1.3±0.2 g cm −3 ) and by at least four craters with diameters equal to or greater than Mathilde's mean radius. The large craters and their unremarkable appearance suggest formation without significant damage to the asteroid outside the final crater volume and indicates the ability of at least this type of asteroid to survive very large impacts without disruption. There are suggestions of structures such as faults or layers. Preservation of such structures would imply that objects that are poor transmitters of impact seismic energy can still maintain some mechanical fabric, and their surface exposure may imply inefficient retention of ejecta from large craters.

Crater populations on Mathilde are assessed from NEAR flyby images. Small to intermediate-sized craters (1/2 to 5 km diameter) approach saturation equilibrium spatial densities and exhibit a range of degradation states similar to the crater populations on Ida. They probably represent a population in quasi-equilibrium between the creation and the destruction of craters caused by saturated impacts on Mathilde. Very large craters, with diameters approaching or exceeding the average radius of Mathilde itself, dominate the shape of Mathilde and represent a uniquely high spatial density of craters on a planetary surface. Formation of these large craters were remarkably ineffective in modifying, let alone destroying, preexisting topography. We speculate that Mathilde's retention of so many large, relatively well-preserved craters may be due to its bulk material properties and/or rubble-pile internal structure. The overall crater populations on Mathilde are consistent with (although do not uniquely require) the same production function (i.e., same projectile power-law size distribution) observed in the cratering records of the Moon, Gaspra, and (apparently) Ida.

Impact craters help characterize the age of a planetary surface, because they accumulate with time. They also provide useful constraints on the importance of surface erosion, as such processes will preferentially remove the smaller craters. Earlier studies of martian crater populations revealed that erosion and dust deposition are important processes on Mars. They disagreed, however, on the age of the youngest volcanism. These earlier studies were limited by image resolution to craters larger than a few hundred metres in diameter. Here we report an analysis, using new images obtained by the Mars Global Surveyor spacecraft, of crater populations that extend the size distribution down to about 16 m. Our results indicate a wide range of surface ages, with one region—lava flows within the Arsia Mons caldera—that we estimate to be no older than 40–100 million years. We suggest that volcanism is a continuing process on Mars.

The images of the four inner small jovian satellites obtained by the Galileo Solid State Imaging (SSI) experiment have much more detailed shape, color, and photometric information than were provided previously by Voyager images. The satellites are in synchronous rotation and show no binary or bifurcated shapes. Thebe and Amalthea have densities of large craters approximately at “empirical equilibrium” levels. The leading sides of Metis, Amalthea, and Thebe are all 25–35% brighter than their trailing sides; the global-average, clear-filter (λ = 0.64 μm) geometric albedos of these three satellites are 0.063, 0.091, and 0.049, respectively. A definite color gradient is observed, with the satellites closer to Jupiter being redder: the mean violet/green ratio (0.42/0.56 μm) decreases from Thebe to Metis. This ratio also is lower for the trailing sides of Thebe and Amalthea than for their leading sides. Bright spots on Amalthea and Thebe are small (

Clear-filter Viking images, and an accurate numerical model of the shape of Phobos, have been used to determine this satellite's photometric properties. A global-average Hapke function derived from disk-resolved data confirms previous indications that Phobos has a strong opposition surge. Photometrically corrected images were mosaicked into an albedo map; most of the resulting normal reflectances are in the range 0.06–0.10, and the brightest region on Phobos is the northeast rim of the crater Stickney, the portion of that rim with the highest concentration of grooves. Globally, there are three albedo classes, reasonably separated geographically: (1) Bright material is to the east and south of Stickney, corresponding approximately to the locations on Phobos having the highest, “bluest” visible/near-IR ratio (Murchie et al . 1991). (2) The darkest material is to Stickney's west, correlating with material having an intermediate visible/NIR ratio (Murchie et al .'s “bluish gray” unit). (3) Intermediate-albedo material dominates the anti-Stickney hemisphere, corresponding to material with a lower visible/NIR ratio (Murchie et al .'s “reddish gray” unit). A search for variations in phase behavior across Phobos' surface shows few such effects overall, limited to isolated areas: (1) Stickney's floor darkens with increasing phase faster than average Phobos. This crater floor is both slightly more backscattering and significantly rougher than the global average; the latter effect may be related to the slumping hinted at in low-resolution images of Stickney. (2) We confirm that in many cases, the contrast between the bright rims of small craters and grooves and their surroundings drops noticeably with increasing phase (phase angles in use = 10°–50°). However, these bright rims, overall, display a diversity of photometric behavior and are the most heterogeneous areas on Phobos in terms of regolith properties. (3) We confirm that dark deposits in the floors of smaller craters darken faster with increasing phase than their surroundings (cf., Goguen et al . 1978) and find that these deposits are more backscattering than average Phobos. Isolated regions with unusual phase behavior occur on Phobos but not on Deimos (Thomas et al . 1996). The variable photometric properties of Phobos' isolated craters and grooves, and the association of global albedo features with Stickney, reinforce interpretations that Phobos' regolith is emplaced and modified by discrete cratering events and is not mixed horizontally by extensive downslope creep as is apparently the case on Deimos.

Ejecta features on the martian satellites are compared to models of ballistic emplacement and downslope motion. The asymmetric distribution of ejecta around the large crater Stickney on Phobos appears most easily explained if the crater were formed when Phobos was only slightly farther from Mars than at present (due to tidal evolution), where rotational and tidal effects allow spreading of low velocity material several kilometers to the east of the crater, but which do not allow a large fraction of the crater's ejecta to be launched into temporary Mars orbit. An additional condition needed to explain the distribution of Stickney's ejecta is that the crater excavation occurred in a gravity-dominated regime. The distribution of Stickney's ejecta is not closely related to the patterns of grooves nor to potential long-distance downslope motion. On Deimos, ejecta from a large impact near the south pole also appear to have been produced by cratering largely in the gravity regime. The morphology of older craters, evidence of ejecta ponding, and crater density data indicate that the ejecta from this large impact have been subject to significant redistribution downslope. This crater is much larger relative to Deimos than Stickney is to Phobos; the resulting greater ejecta cover and possibly much greater seismic shaking on Deimos may help explain the strikingly different appearances of the martian moons.

Hubble Space Telescope images of asteroid 4 Vesta obtained during the favorable 1996 apparition show an impact crater 460 kilometers in diameter near the south pole. Color measurements within the 13-kilometer-deep crater are consistent with excavation deep into a high-calcium pyroxene-rich crust or olivine upper mantle. About 1 percent of Vesta was excavated by the crater formation event, a volume sufficient to account for the family of small Vesta-like asteroids that extends to dynamical source regions for meteorites. This crater may be the site of origin for the howardite, eucrite, and diogenite classes of basaltic achondrite meteorites.

Among the well-imaged small satellites and asteroids, Deimos displays a unique surface: very smooth with global-scale albedo features. We have examined the disk-resolved photometry of Deimos using Viking Orbiter images for clues to its distinctive appearance. Hapke parameters were fit to characterize the phase behavior and to compute normal reflectance. The opposition surge amplitude (B0) is smaller for Deimos than for Phobos. Outside the range of the opposition effect the two martian satellites have similarly shaped phase curves, but Deimos is about 20–30% brighter than Phobos from 10°–80° phase. The calculated mean normal reflectance of Deimos (λeff= 0.54 μm) is 0.068 ± 0.007. The brighter and darker areas on Deimos exhibit constant contrast between 0.6° and 81° phase; this characteristic allows a calculation of the range of normal reflectances over most of its surface, nearly all of which values are between 0.06 and 0.09. The trailing side of Deimos has a larger relative distribution of the brighter material, and is on average about 10% brighter than the leading side. The mean normal reflectance cannot be formally distinguished from that of Phobos (0.071 ± 0.012; Simonelli, D. P., M. Wisz, A. Switala, D. Adinolfi, J. Veverka, P. C. Thomas, and P. Helfenstein 1996. Submitted toIcarus). Although the statistical distribution of normal reflectances on the two satellites is similar, the geography of the albedo variations is very different. Deimos has gradational changes in albedo downslope from ridge crests, primarily manifested in long albedo “streamers.” On Phobos there is a more patchy global distribution of albedos, apparently related to ejecta from the large crater Stickney. Because of the similarity of mean density, spectral properties, mean normal reflectance, the range of normal reflectance, and phase function outside the opposition effect, Deimos appears to be made of materials with compositions very similar to those on Phobos, although the apparent wider distribution of ejecta on Deimos has been cited as indicating a greater role for strength scaling in cratering on Deimos (Lee, S. W., P. Thomas, and J. Veverka 1986.Icarus68, 77–86). Simple modeling of the formation of the albedo patterns by gardening, creep, and “weathering” of bright material from crater rims suggests that impact gardening contributes very little to the motion of the material downslope, and that vertical mixing and/or “weathering” must be important in addition to an unspecified creep process. The distinction of Deimos is primarily in the smooth surface that allows a particularly large scale of downslope movement of regolith on very gentle slopes. This smoothness is most easily explained by the effects from impact formation of a 10-km concavity at high southern latitudes in the latter half of Deimos' surface history. This impact scar is relatively much larger than is the crater Stickney on Phobos. The effects of this large impact probably include blanketing by an average of nearly 200 m of ejecta, but also may include seismic erasing of craters similar to that proposed for Ida by Asphauget al. (Asphaug, E., J. M. Moore, D. Morrison, W. Benz, and R. A. Sullivan 1996.Icarus120, 158–184).

Craters on Ida have a size distribution consistent with equilibrium under conditions of crater saturation over the range from 100 m to 1 km in diameter. The density of craters is high and they follow a −3 differential slope. The craters show a full spectrum of degradational morphologies at both large and small sizes. Craters >2 km in diameter (and especially those >5 km) are more prevalent on one part of Ida; a sub-equilibrium, production population is seen on the other part, which rises steeply and reaches equilibrium for diameters

Galileo's flyby of 243 Ida in August 1993 led to the discovery of a small satellite, Dactyl, some 85 km from the asteroid's center. From Earth at mean opposition, the satellite is a V = +20.3 mag object (some 6.7 magnitudes fainter than Ida). Forty-seven images of the satellite at 18 different observing times were played back, including one multicolor sequence in which the satellite is resolved adequately to distinguish surface markings (∼105 m/pxl) and three higher resolution single-color views (89, 39, and 24 m/pxl). The satellite, mean radius = 0.7 km, is an elongated, but not angular body with principal diameters of 1.6 × 1.4 × 1.2 km. In the highest resolution view, the longest axis points approximately in the direction of Ida, and its shortest axis is perpendicular to the orbital plane. The spin period is slow (> 8 hr?) and may be synchronous. The satellite shows no conspicuous sharp edges and is much less irregular in shape than Ida. Limb profiles are remarkably smooth over distances of 200–300 m. The geometric albedos of the two objects are similar (0.20 vs 0.21), as are the 0.4–1.0-μm colors. Like Ida, Dactyl is an S-asteroid, but has a slightly deeper 1-μm band than Ida (by 5–8%). While no identical regions (in color) are seen on Ida, the color difference is consistent with color variations reported within the Koronis family and may be due to a slightly higher pyroxene/olivine ratio on the satellite. More than a dozen craters ranging from ≲90 to 280 m diameter are visible in the best image (39 m/pxl at 47° phase). The largest contains an off-centered, positive relief feature some 75 m across. The image includes an intriguing crater chain, but no grooves, ridges, or sharp edges are evident. In terms of limb roughness, Dactyl is much smoother than Ida, but comparable to the two satellites of Mars, Phobos and Deimos. While the satellite's origin is uncertain, a likely scenario would have the satellite date from the breakup of the Koronis family. It is interesting that crater densities on the satellite are similar to those on Ida itself.

The surface of 243 Ida is dominated by the effects of impacts. No complex crater morphologies are observed. A complete range of crater degradation states is present, which also reveals optical maturation of the surface (darkening and reddening of materials with increasing exposure age). Regions of bright material associated with the freshest craters might be ballistically emplaced deposits or the result of seismic disturbance of loosely-bound surface materials. Diameter/depth ratios for fresh craters on Ida are ∼1:6.5, similar to Gaspra results, but greater than the 1:5 ratios common on other rocky bodies. Contributing causes include rim degradation by whole-body “ringing,” relatively thin ejecta blankets around crater rims, or an extended strength gradient in near-surface materials due to low gravitational self-packing. Grooves probably represent expressions in surface debris of reactivated fractures in the deeper interior. Isolated positive relief features as large as 150 m are probably ejecta blocks related to large impacts. Evidence for the presence of debris on the surface includes resolved ejecta blocks, mass-wasting scars, contrasts in color and albedo of fresh crater materials, and albedo streaks oriented down local slopes. Color data indicate relatively uniform calcium abundance in pyroxenes and constant pyroxene/olivine ratio. A large, relatively blue unit across the northern polar area is probably related to regolith processes involving ejecta from Azzurra rather than representing internal compositional heterogeneity. A small number of bluer, brighter craters are randomly distributed across the surface, unlike on Gaspra where these features are concentrated along ridges. This implies that debris on Ida is less mobile and/or consistently thicker than on Gaspra. Estimates of the average depth of mobile materials derived from chute depths (20–60 m), grooves (≥30 m), and shallowing of the largest degraded craters (20–50 m minimum, ∼100 m maximum) suggest a thickness of potentially mobile materials of ∼50 m, and a typical thickness for the debris layer of 50–100 m.

Seventeen positive relief features ∼45–150 m across are identified as probable blocks in Galileo high-resolution images of Ida. Their presence provides direct evidence for regolith retention on asteroids. The spatial distribution, maximum size, and integrated volume of the blocks are consistent with those of blocks associated with craters on the Earth, the Moon, Phobos, and Deimos. The concordance suggests that the features are impact ejecta fragments and that cratering mechanics on Ida, an object of average diameter ∼32 km, are similar to those applying on previously studied rocky bodies. The blocks that lie within or near the rims of craters Lascaux and Mammoth were likely mobilized in the low-velocity tail portion of the excavation flow that formed those craters. A few blocks located near smaller craters may have migrated some distance away from their source, possibly by impact-induced spallation, hopping, rolling, and/or sliding. Some blocks on Ida could be surviving fragments from the Koronis parent body, accreted after its breakup. The lifetime of 102-m sized boulders against collisional disruption is estimated to be in the 108–109year range, consistent with ages considered for the largest and oldest craters on Ida. Extrapolation of successful ejecta scaling laws to other asteroids suggests that blocks ∼15 and 70 m across could be present on Dactyl and Gaspra, respectively (in both cases too small to be identified in available Galileo images). Blocks 100 m in size could be present on 433 Eros, and km-sized megablocks on 4 Vesta.

Galileo SSI color data between 0.4 and 1.0 μm demonstrate that both Ida and Dactyl are S-type asteroids with similar, but distinct spectra. Small but definite color variations are also observed on Ida itself and involve both the blue part of the spectrum and the depth of the 1-μm pyroxene–olivine band. Ida's surface can be classified into two color terrains: Terrain A has a shallower 1-μm absorption and a steeper visible red slope than does Terrain B. Qualitatively, the color–albedo systematics of these two terrains follow those noted for color units on Gaspra and the variations in 1-μm band depth with weathering described by Gaffeyet al.(Gaffey, M. J., J. F. Bell, R. H. Brown, T. H. Burbine, J. Piatek, K. L. Reed, and D. A. Chaky 1993.Icarus106, 573–602). Terrain A, with its slightly lower albedo, its shallower 1-μm band, and its slightly steeper visible red slope relative to Terrain B could be interpreted as the “more processed,” “more mature,” or the “more weathered” of the two terrains. Consistent with this interpretation is that Terrain A appears to be the ubiquitous background on most of Ida, while Terrain B is correlated with some small craters as well as with possible ejecta from the 10-km Azzurra impact structure. Because of these trends, it is less likely that differences between Terrains A and B are caused by an original compositional inhomogeneity within the body of Ida, although they do fall within the range known to occur within the Koronis family. The spectrum of Dactyl is similar to, but definitely different from, that of Terrain B on Ida. It does not conform to the pattern that obtains between the colors and albedos of Terrains A and B: the satellite's 1-μm band is deeper than that of Terrain B, but its albedo is lower, rather than higher. By itself, the deeper band depth could be interpreted, following Gaffeyet al., to mean that Dactyl is a less weathered version of Terrain B on Ida, but such an interpretation is at odds with Dactyl's redder spectral slope. Thus, the explanation for the color difference between Dactyl and Ida is likely to be different from that which accounts for the differences between the two terrains on Ida. Given that Dactyl and Ida have very similar photometric properties (Helfenstein, P., J. Veverka, P. C. Thomas, D. P. Simonelli, K. Klassen, T. V. Johnson, F. Fanale, J. Granahan, A. S. McEwen, M. J. S. Belton, and C. R. Chapman 1996Icarus120, 48–65), thus ruling out any dramatic texture differences between the two surfaces, the most likely explanation is that the satellite has a slightly different composition (more pyroxene?) than Ida. The spectral difference is within the range reported by Binzelet al.(Binzel, R. P., S. Xu, and S. J. Bus 1993.Icarus106, 608–611.) for members of the Koronis family, and could be caused by compositional inhomogeneities of the Koronis parent body rather than by post-breakup regolith processes.

Voyager 2 images have been used to measure the instantaneous rotation state of Hyperion during the period of the spacecraft's 1981 flyby of Saturn, to define the satellite's shape, and to map geological features. For the 38 hr of image coverage at resolutions better than 17 km/pixel, control point calculations yield a rotation rate of 72°/day, with a pole direction of RA = 226 ± 4°, Dec = 35 ± 4°. Rotation was essentially about the long axis of Hyperion. A shape model with 5° × 5° spacing is developed using stereo, limb, and terminator data. About 95% of the surface of Hyperion is constrained within 20 km by the control point and lower resolution images. An ellipsoidal fit to this model gives principal axes (radii) of 164, 130, and 107 km. The shape model has a mean radius of 133 ± 8 km; the model moments of inertia (I/MR2) are A = 0.32 ± 0.02, B = 0.46 ± 0.02, and C = 0.52 ± 0.02. G. J. Black et al . (1995, Icarus) have used the shape model and the instantaneous rotation state to calculate the expected rotation of Hyperion throughout the Voyager 2 encounter. They predict significant changes in the rotation pole orientation, but only minor changes in rotation rate during the 18 days of low-resolution, disk-resolved imaging of Hyperion. We compare the Voyager-observed lightcurve of Hyperion to lightcurves predicted from the dynamical model. During this 18-day period the dynamical model matches the observed lightcurve very well; the best lightcurve fit is for a rotation rate of 72 + 3 - 4°/day (identical to the control point solution) at the time of the first of the high-resolution images. Much of the Hyperion lightcurve amplitude derives from wobble rather than from the spin. The good agreement of the dynamical model with observed behavior suggests that Hyperion's moments of inertia are well modeled by the shape and that it has no substantial internal density contrasts. The overall shape of Hyperion is typical for icy objects smaller than 200 km mean radius. Small icy satellites appear to have less structural control of their shapes than do rocky satellites and asteroids. A 300-km-long arcuate scarp and a more subdued, nearly parallel scarp suggest spallation of several kilometers of material from nearly 1/4 the area of the satellite due to the formation of a 140-km-diameter crater.

The first images of the asteroid 243 Ida from Galileo show an irregular object measuring 56-kilometers by 24 kilometers by 21 kilometers. Its surface is rich in geologic features, including systems of grooves, blocks, chutes, albedo features, crater chains, and a full range of crater morphologies. The largest blocks may be distributed nonuniformly across the surface; lineaments and dark-floored craters also have preferential locations. Ida is interpreted to have a substantial regolith. The high crater density and size-frequency distribution (-3 differential power-law index) indicate a surface in equilibrium with saturated cratering. A minimum model crater age for Ida-and therefore for the Koronis family to which Ida belongs-is estimated at 1 billion years, older than expected.

The surface of Gaspra can be divided into several facets separated by ridges. Superimposed on the facets and ridges are two populations of craters. Type 1 craters are crisply defined and constitute a production function of impact origin. Type 2 craters are mostly irregular shallow depressions and probably of various origins. Some of the type 2 craters may have formed by impact and be remnants of a crater population that predates the type 1 population. The number of type 1 craters suggests that they started to accumulate 2 × 107 to 3 × 108 years ago. The freshest craters have a depth/diameter ratio of 1:7, as compared with 1:5 typically observed on other bodies. The craters appear to have become shallower with age at a rate of 10-6 to 10-7 m/year. The shallow depth of newly formed craters and the progressive degradation with age are ascribed to downslope movement of poorly coherent surface materials, the movement being aided by seismic shaking as a result of impact. The surface shows subtle color differences. The most prominent differences are observed around craters on ridges, where the surface has a stronger 1-μm absorption than elsewhere. The contrast is not observed around craters on the facets. The color differences are attributed to alteration of the materials brought to the surface. The unaltered material is estimated to be located at least 50 m below the surface on the facets. Only on the ridges is unaltered material at a shallow enough depth to be excavated by the observed craters. The rate of degradation of craters, the rounded form of the ridges, the depth of altered material on the facets, and the presence of old degraded craters all suggest that Gaspra is covered with a regolith a few tens to several tens of meters thick.

[1] We use limb profiles to quantify the long‐wavelength topography of the Saturnian satellites. The degree 2 shapes of Mimas, Enceladus, and Tethys are not consistent with hydrostatic equilibrium. We derive 2‐D topographic maps out to spherical harmonic degree 8. There is a good correlation with topography derived from stereo techniques. If uncompensated, topography at degree 3 and higher is large enough to be detectable during close spacecraft flybys. If not properly accounted for, this topography may bias estimates of a satellite’s degree 2 gravity coefficients (which are used to determine the moment of inertia). We also derive a one‐dimensional variance spectrum (a measure of how roughness varies with wavelength) for each body. The short‐wavelength spectral slope is − 2t o−2.5, similar to silicate bodies. However, unlike the terrestrial planets, each satellite spectrum shows a reduction in slope at longer wavelengths. If this break in slope is due to a transition from flexural to isostatic support, the globally averaged elastic thickness Te of each satellite may be derived. We obtain Te values of ≥5 km, 1.5–5 km, ≈5 km, and ≥5 km for Tethys, Dione, Rhea, and Iapetus, respectively. For Europa, we obtain Te ≈ 1.5 km. These estimates are generally consistent with estimates made using other techniques. For Enceladus, intermediate wavelengths imply Te ≥ 0.5 km, but the variance spectrum at wavelengths greater than 150 km is probably influenced by long‐ wavelength processes such as convection or shell thickness variations. Impact cratering may also play a role in determining the variance spectra of some bodies. Citation: Nimmo, F., B. G. Bills, and P. C. Thomas (2011), Geophysical implications of the long‐wavelength topography of the Saturnian satellites, J. Geophys. Res., 116, E11001, doi:10.1029/2011JE003835.

[1] NEAR Shoemaker images reveal widespread occurrence of tectonic landforms on asteroid 433 Eros. Hinks Dorsum is a ridge that extends for about 18 km around the asteroid and strongly resembles thrust fault structures on the terrestrial planets. Tectonic landforms can provide information on the mechanical properties of asteroids, a subject of much controversy. Modeling constrained by topographic data shows that Hinks Dorsum can be accounted for by a shallow rooted thrust fault no greater than 250 m in depth with ∼90 m of cumulative slip. Strength envelopes based on frictional and rock mass strength criteria suggest the near-surface shear strength of Eros is from ∼1 to 6 MPa. A spatial correlation is found between Shoemaker crater, a transition from low to high crater density, and Hinks Dorsum. This spatial relation along with the estimated strength of the asteroid suggests the thrust fault was formed by impact induced compression.

The loose material--regolith--on the surfaces of asteroids is thought to represent ballistically emplaced ejecta from impacts but the identification of source craters and the detailed study of the regolith modification have been hampered by the limited spatial resolution and area coverage of the few asteroids imaged by spacecraft. Here we report the results of global mapping of the asteroid 433 Eros from high-resolution images obtained by the NEAR-Shoemaker spacecraft. Based on the images and ejecta-emplacement models, we suggest that most large ejecta blocks on Eros originate from a relatively young 7.6-km-diameter crater. A large fraction of the ejecta from impacts pre-dating that crater has apparently been buried or eroded. The images also show evidence for the action of a variety of sorting environments for regolith particles after they are deposited on the surface.

Iapetus Revealed The striking appearance of Saturn's moon Iapetus—half black and half white—has puzzled astronomers for over three centuries. Now Spencer and Denk (p. 432 , published online 10 December) present an explanation for this asymmetry: A thermally controlled runaway migration of water ice triggered by exogenic deposition of dark material on the moon's leading darker side, which faces the direction of motion of the moon in its orbit around Saturn. This mechanism is unique to Iapetus because it rotates slowly enough to allow large temperature variations to arise, it is small enough to allow long-range migration of water, and there is a source of dust to trigger the process. In a related paper, Denk et al. (p. 435 , published online 10 December) present data derived from the Cassini Imaging Science Subsystem that reveal that both dark and bright materials on the leading side of Iapetus are redder than their trailing-side counterparts. This asymmetry results from the deposition of dust and debris from other moons in the saturnian system—the very same process that initiates the thermal segregation proposed above.

We have used Cassini stereo images to study the topography of Iapetus' leading side. A terrain model derived at resolutions of 4–8 km reveals that Iapetus has substantial topography with heights in the range of −10 km to +13 km, much more than observed on the other middle-sized satellites of Saturn so far. Most of the topography is older than 4 Ga [Neukum, G., Wagner, R., Denk, T., Porco, C.C., 2005. Lunar Planet. Sci. XXXVI. Abstract 2034] which implies that Iapetus must have had a thick lithosphere early in its history to support this topography. Models of lithospheric deflection by topographic loads provide an estimate of the required elastic thickness in the range of 50–100 km. Iapetus' prominent equatorial ridge [Porco, C.C., and 34 colleagues, 2005. Science 307, 1237–1242] reaches widths of 70 km and heights of up to 13 km from their base within the modeled area. The morphology of the ridge suggests an endogenous origin rather than a formation by collisional accretion of a ring remnant [Ip, W.-H., 2006. Geophys. Res. Lett. 33, doi: 10.1029/2005GL025386 . L16203]. The transition from simple to complex central peak craters on Iapetus occurs at diameters of 11 ± 3 km . The central peaks have pronounced conical shapes with flanking slopes of typically 11° and heights that can rise above the surrounding plains. Crater depths seem to be systematically lower on Iapetus than on similarly sized Rhea, which if true, may be related to more pronounced crater-wall slumping (which widens the craters) on Iapetus than on Rhea. There are seven large impact basins with complex morphologies including central peak massifs and terraced walls, the largest one reaches 800 km in diameter and has rim topography of up to 10 km. Generally, no rings are observed with the basins consistent with a thick lithosphere but still thin enough to allow for viscous relaxation of the basin floors, which is inferred from crater depth-to-diameter measurements. In particular, a 400-km basin shows up-domed floor topography which is suggestive of viscous relaxation. A model of complex crater formation with a viscoplastic (Bingham) rheology [Melosh, H.J., 1989. Impact Cratering. Oxford Univ. Press, New York] of the impact-shocked icy material provides an estimate of the effective cohesion/viscosity at 0.04 ± 0.01 MPa/ 0.6 ± 0.2 GPa s . The local distribution of bright and dark material on the surface of Iapetus is largely controlled by topography and consistent with the dark material being a sublimation lag deposit originating from a bright icy substrate mixed with the dark components, but frost deposits are possible as well.

Saturn's moon Hyperion, an irregular shaped object in a tumbling orbit, looks odd: the Cassini flyby of September 2005 revealed a unique spongy surface. Two papers this week present the initial Cassini results. First, imaging and radio data suggest that the spongy appearance is caused by impact cratering on a porous body. And second, near-infrared and ultraviolet spectroscopy reveal details of the surface composition of the highly reflective areas that cover much of the surface, and also of the darker areas, mostly at the bottom of craters. The spectra are consistent with the presence of water ice contaminated with an organic solid. The low-albedo (dark) material is spectroscopically similar to that found on two other saturnian moons, Iapetus and Phoebe, containing a mixture of water ice, complex organics, carbon dioxide and nitriles. This cocktail of materials resembles those seen in comets and probably in Kuiper Belt objects. High resolution surface images of Hyperion, Saturn's largest known, irregularly shaped satellite are presented. These reveal a unique sponge-like appearance at scales of a few kilometres, with a high surface density of relatively well-preserved 2–10 km-sized craters. Hyperion is Saturn’s largest known irregularly shaped satellite and the only moon observed to undergo chaotic rotation1,2,3. Previous work has identified Hyperion’s surface as distinct from other small icy objects4,5 but left the causes unsettled. Here we report high-resolution images that reveal a unique sponge-like appearance at scales of a few kilometres. Mapping shows a high surface density of relatively well-preserved craters two to ten kilometres across. We have also determined Hyperion’s size and mass, and calculated the mean density as 544 ± 50 kg m-3, which indicates a porosity of >40 per cent. The high porosity may enhance preservation of craters by minimizing the amount of ejecta produced or retained6,7, and accordingly may be the crucial factor in crafting this unusual surface.

Impact cratering is a fundamental process on planets and asteroids. Many aspects of crater formation have been examined on the Earth and with spacecraft but one phenomenon that has been difficult to identify clearly is the seismic effect of large impacts. The detailed mapping of the entire asteroid Eros, orbited by the NEAR-Shoemaker spacecraft in 2000–01, has revealed a striking pattern of destruction of small craters. The creation of one large crater, for example, erased most smaller craters from nearly 40% of the asteroid's surface. This finding will allow inferences of interior and surface characteristics of asteroids and satellites. Impact cratering creates a wide range of topography on small satellites and asteroids. The population of visible craters evolves with impacts, and because there are no competing endogenic processes to modify the surface, determining the various ways younger craters add to or subtract from the population is a fundamental aspect of small-body geology1,2. Asteroid 433 Eros, the most closely studied small body, has regions of substantially different crater densities3,4,5 that remain unexplained. Here we show that the formation of a relatively young crater (7.6 km in diameter) resulted in the removal of other craters as large as 0.5 km over nearly 40 per cent of the asteroid. Burial by ejecta cannot explain the observed pattern of crater removal. The limitation of reduced crater density to a zone within a particular straight-line distance through the asteroid from the centre of the large crater suggests degradation of the topography by seismic energy6 released during the impact. Our observations indicate that the interior of Eros is sufficiently cohesive to transmit seismic energy over many kilometres, and the outer several tens of metres of the asteroid must be composed of relatively non-cohesive material.

OBSERVATIONS of stellar occultations by asteroids have suggested that some may have satellites1. But given the absence of any confirmatory evidence, the prevailing view has been that although such satellites probably do exist, they are likely to be rare2. Here we report the discovery3 by the Galileo spacecraft of a satellite associated with the asteroid 243 Ida. Although the satellite, Dactyl, is only 1.6km across, it has been imaged with sufficient resolution for geological analysis. We describe the physical properties of Dactyl, with emphasis on its notably smooth shape, its crater population (which includes a crater chain) and its photometric properties. We find that, spectroscopically, Dactyl resembles both Ida and the other members of the Koronis asteroid family, implying a similar composition; small spectral differences may reflect a space weathering process that slightly alters the colours with time. We argue that Dactyl originated during the breakup of the Koronis parent body, and that satellites could be common around other asteroids (particularly members of asteroid families).

On 25 October 2000, the Near Earth Asteroid Rendevous (NEAR)–Shoemaker spacecraft executed a low-altitude flyover of asteroid 433 Eros, making it possible to image the surface at a resolution of about 1 meter per pixel. The images reveal an evolved surface distinguished by an abundance of ejecta blocks, a dearth of small craters, and smooth material infilling some topographic lows. The subdued appearance of craters of different diameters and the variety of blocks and different degrees of their burial suggest that ejecta from several impact events blanketed the region imaged at closest approach and led to the building up of a substantial and complex regolith consisting of fine materials and abundant meter-sized blocks.

High-resolution images of the martian surface at scales of a few meters show ubiquitous erosional and depositional eolian landforms. Dunes, sandsheets, and drifts are prevalent and exhibit a range of morphology, composition (inferred from albedo), and age (as seen in occurrences of different dune orientations at the same location). Steep walls of topographic depressions such as canyons, valleys, and impact craters show the martian crust to be stratified at scales of a few tens of meters. The south polar layered terrain and superposed permanent ice cap display diverse surface textures that may reflect the complex interplay of volatile and non-volatile components. Low resolution regional views of the planet provide synoptic observations of polar cap retreat, condensate clouds, and the lifecycle of local and regional dust storms.

The first images of Jupiter, Io, Europa, and Ganymede from the Galileo spacecraft reveal new information about Jupiter's Great Red Spot (GRS) and the surfaces of the Galilean satellites. Features similar to clusters of thunderstorms were found in the GRS. Nearby wave structures suggest that the GRS may be a shallow atmospheric feature. Changes in surface color and plume distribution indicate differences in resurfacing processes near hot spots on Io. Patchy emissions were seen while Io was in eclipse by Jupiter. The outer margins of prominent linear markings (triple bands) on Europa are diffuse, suggesting that material has been vented from fractures. Numerous small circular craters indicate localized areas of relatively old surface. Pervasive brittle deformation of an ice layer appears to have formed grooves on Ganymede. Dark terrain unexpectedly shows distinctive albedo variations to the limit of resolution.

Galileo images of Gaspra reveal it to be an irregularly shaped object (19 by 12 by 11 kilometers) that appears to have been created by a catastrophic collisional disruption of a precursor parent body. The cratering age of the surface is about 200 million years. Subtle albedo and color variations appear to correlate with morphological features: Brighter materials are associated with craters especially along the crests of ridges, have a stronger 1-micrometer absorption, and may represent freshly excavated mafic materials; darker materials exhibiting a significantly weaker 1-micrometer absorption appear concentrated in interridge areas. One explanation of these patterns is that Gaspra is covered with a thin regolith and that some of this material has migrated downslope in some areas.

Voyager 2 images revealed six small satellites orbiting within five planetary radii of Neptune. The largest, Proteus (1989N1), has a mean radius of 208 ± 8 km; the smallest, Naiad (1989N6), is 29 ± 6 km in radius. The four satellites for which reasonable radius measurements can be made all have geometric albedos (not accounting for opposition effects) of 0.06 at 0.48 pm. The albedos and phase coefficients of these satellites are very similar to those measured for the inner satellites of Uranus but are quite distinct from the values for Nereid (p= 0.16-0.20, β = 0.028 magnitude/degree) as determined by Thomas et al. (this issue). Color measurements (0.4 to 0.6 μm) of Proteus show a neutral color. The two largest satellites are resolved as irregular objects and have limb roughnesses comparable to those of similar-sized satellites of other planets. One image of Proteus shows a 150-km crater with some overlapping smaller craters; this feature allows measurement of a rotation rate consistent with a synchronous period. The irregular shape and heavily cratered, photometrically bland surface of Proteus are consistent with characteristics of a primitive, undifferentiated object. Evidently, Proteus has escaped the tidal heating believed to have been responsible for the remarkable geologic processes on the similar-sized bodies Miranda and Enceladus and perhaps Mimas.

The characteristics of the grooves on Phobos indicate that these features were formed by a large impact, an observation which implies that grooves may occur on other small bodies. Published data on fragmentation experiments suggest that the energy of the impact which produced the grooves on Phobos was in a critical range, strong enough to produce fracturing at large distances from the crater, but too weak to fragment the object. Other small bodies, such as asteroids, which have experienced an impact flux sufficient to include one such critical cratering event should show grooves, unless they have been fragmented subsequently by a more severe impact. It is estimated that between 1/12th and 1/4th of asteroids smaller than 100 km in diameter may have surface features resembling grooves associated with large impact craters. Several morphologically distinct types of grooves can be expected, depending on the thickness and composition of the regolith of the parent body.

The paper considers Mariner 9 and Viking data that contradict Kuzmin's (1975) hypothesis that all crater-associated wind streaks on Mars are depositional and consist of unresolved barchan-like dunes. According to Kuzmin's hypothesis, any streak can appear either bright or dark relative to its surroundings depending on the sun's position. The spacecraft images, however, show examples of dark and light streaks visible at the same azimuth angle of the sun. Evidence that bright and dark streaks differ both in morphology and in character is considered. It is suggested that the common ragged dark streaks are probably erosion scars while most bright streaks probably represent accumulations of bright dust-storm fallout.

Ejecta block characteristics observed on Phobos, Deimos, and the moon are examined. The analyzed source craters on Deimos are 0.8-2.3 km in diameter, those on Phobos are 1.5-10 km, and the lunar craters are between 0.2-3.5 km in diameter. The size and radial distribution of the ejecta blocks for the three bodies are compared. It is observed that the size distribution of the ejecta blocks surrounding craters on the three objects are basically similar, and the radial distribution of the blocks for Phobos and the moon are the same (within 2 radii of the crater center); however, the ejecta on Deimos are more dispersed (greater than or equal to 2 radii from the crater center).

The density of craters larger than 1 km in diameter has been determined for the entire surface of Phobos, and half that of Deimos. Densities of craters as small as 10 m on Phobos and 5 m on Deimos have been measured for small areas of the satellites. On both objects, crater densities are similar and yield plots which have slopes close to -1.9 on both incremental and cumulative log-log graphs. These densities are close to those expected to obtain under equilibrium conditions. They are also near the maximum predicted, based on the fragmentation lifetimes of the two objects: that is, the densities are near to the maximum possible before such objects are likely to suffer an impact severe enough to disrupt them. While the observed crater densities cannot be converted to absolute ages in any rigorous fashion, they can be understood if the flux at Mars has been similar to that at the Moon and if the surfaces that we see today generally date back to the end of the period of heavy bombardment some 4 billion years ago. It is extremely unlikely that the surfaces are younger than 1 billion years. There are no large areas on Phobos for which crater densities differ by more than a factor of 3 from the average.