Eros: Shape, Topography, and Slope Processes

Stereogrammetric measurement of the shape of Eros using images obtained by NEAR's Multispectral Imager provides a survey of the major topographic features and slope processes on this asteroid. This curved asteroid has radii ranging from 3.1 to 17.7 km and a volume of 2535±20 km³. The center of figure is within 52 m of the center of mass provided by the Navigation team; this minimal difference suggests that there are only modest variations in density or porosity within the asteroid. Three large depressions 10, 8, and 5.3 km across represent different stages of degradation of large impact craters. Slopes on horizontal scales of ∼300 m are nearly all less than 35°, although locally scarps are much steeper. The area distribution of slopes is similar to those on Ida, Phobos, and Deimos. Regions that have slopes greater than 25° have distinct brighter markings and have fewer large ejecta blocks than do flatter areas. The albedo patterns that suggest downslope transport of regolith have sharper boundaries than those on Phobos, Deimos, and Gaspra. The morphology of the albedo patterns, their lack of discrete sources, and their concentration on steeper slopes suggest transport mechanisms different from those on the previously well-observed small bodies, perhaps due to a reduced relative effectiveness of impact gardening on Eros. Regolith is also transported in talus cones and in connected, sinuous paths extending as much as 2 km, with some evident as relatively darker material. Talus material in at least one area is a discrete superposed unit, a feature not resolved on other small bodies. Flat-floored craters that apparently contain ponded material also suggest discrete units that are not well mixed by impacts.     

Crater densities on the satellites of Mars

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.     

Life near the Roche limit: Behavior of ejecta from satellites close to planets

The distribution of ejecta from impact craters significantly affects the surface characters of satellites and asteroids. In order to understand better the distinctive features seen on Phobos, Deimos, and Amalthea, we study the dynamics of nearby debris but include several factors — planetary tides plus satellite rotation and nonspherical shape-that complicate the problem. We have taken several different approaches to investigate the behavior of ejecta from satellites near planets. For example, we have calculated numerically the usual pseudoenergy (Jacobi) integral. This is done in the framework of a restricted three-body problem where we model the satellites as triaxial ellipsoids rather than point masses as in past work. Iso-contours of this integral show that Deimos and Amalthea are entirely enclosed by their Roche lobes, and the surfaces of their model ellipsoids lie nearly along equipotentials. Presumably this was once also the case for Phobos, before tidal evolution brought it so close to Mars. Presently the surface of Phobos overflows its Roche lobe, except for the regions within a few kilometers of the sub- and anti-Mars points. Thus most surface material on Phobos is not energetically bound: nevertheless it is retained by the satellite because local gravity has an inward component everywhere. Similar situations probably prevail for the newly discovered satellite of Jupiter (J14) and for the several objects found just outside Saturn's rings. We have also examined the fate of crater ejecta from the satellites of Mars by numerical integration of trajectories for particles leaving their surfaces in the equatorial plane. The ejecta behavior depends dramatically on the longitude of the primary impact, as well as on the speed and direction of ejection. Material thrown farther than a few degrees of longitude remains in flight for an appreciable time. Over intervals of an hour or more, the satellites travel through substantial arcs of their orbits, so that the Coriolis effect then becomes important. For this reason the limit of debris deposition is elongated toward the west while debris thrown to the east escapes at lower ejection velocities. We display some typical trajectories, which include many interesting special effects, such as loops, cusps, “folded” ejecta blankets, and even a temporary satellite of Deimos. Besides being important for understanding the formation of surface features on satellites, our work is perhaps pertinent to regolith development on small satellites and asteroids, and also to the budgets of dust belts around planets.     

Radial variation of lunar crater rim topography

The variation of rim topography as a function of range from the crater rim has been determined for a group of morphologically fresh lunar craters (D = 10–140 km) using the recent series of Lunar Topographic Orthophotomaps. The rate at which exterior crater topography converges with the surrounding surface is highly variable along different radial directions at individual craters as well as between different craters. At several craters, oblique impact appears to have contributed to azimuthal elevation/range variations. The topographic expression of a crater above the surrounding surface typically decreases to one-tenth of the estimated rim height at a range of 1.3R–1.7R, well within the rough-textured ejecta deposit surrounding the crater. Comparisons with terrestrial craters suggest that the topographic crater rim is predominantly a structural feature. In most craters large portions of the hummocky facies and virtually all of the radial facies, in spite of their rough appearance and local topographic variations, provide no significant net topographic addition to the preexisting surface. The extreme variability of crater rim topography strongly suggests that ejecta thicknesses are highly variable and that a unique power-law expression cannot truly represent the radial variation of ejecta deposit thickness.     

A study of lunar impact crater size-distributions

Discrepancies in published crater frequency data prompted this study of lunar crater distributions. Effects modifying production size distributions of impact craters such as surface lava flows, blanketing by ejecta, superposition, infilling, and abrasion of craters, mass wasting, and the contribution of secondary and volcanic craters are discussed. The resulting criteria have been applied in the determination of the size distributions of unmodified impact crater populations in selected lunar regions of different ages. The measured cumulative crater frequencies are used to obtain a general calibration size distribution curve by a normalization procedure. It is found that the lunar impact crater size distribution is largely constant in the size range 0.3 km ?D ? 20 km for regions with formation ages between ≈ 3 × 10⁹ yr and ? 4 × 10⁹ yr. A polynomial of 4th degree, valid in the size range 0.8 km ?D ? 20 km, and a polynomial of 7th degree, valid in the size range 0.3 km ?D ? ? 20 km, have been approximated to the logarithm of the cumulative crater frequencyN as a function of the logarithm of crater diameterD. The resulting relationship can be expressed asN ～D ^α(D) where α is a function depending onD. This relationship allows the comparison of crater frequencies in different size ranges. Exponential relationships with constant α, commonly used in the literature, are shown to inadequately approximate the lunar impact crater size distribution. Deviations of measured size distributions from the calibration distribution are strongly suggestive of the existence of processes having modified the primary impact crater population.     

Dynamics of Mars-orbiting dust: Effects of light pressure and planetary oblateness

This paper deals with dynamics of impact ejecta from Phobos and Deimos initially on near-circular equatorial orbits around Mars. For particles emitted in a wide size regime of 1 micron and greater, and taking into account the typical particle lifetimes to be less than 100 years, the motion is governed by two perturbing forces: solar radiation pressure and influence of Mars' oblateness. The equations of motion of particles in Lagrangian non-singular elements are deduced and solved, both analytically and numerically, for different-sized ejecta. We state that the coupled effect of both forces above is essential so that on no account can the oblateness of Mars are be neglected. The dynamics of grains prove to be quite different for the ejecta of Phobos and Deimos. For Deimos, the qualitative results are relatively simple and imply oscillations of eccentricity and long-term variations of orbital inclination, with amplitudes and periods both depending on grain size. For Phobos, the dynamics are shown to be much more complicated, and we discuss it in detail. We have found an intriguous peculiar behavior of debris near 300 µm in size. Another finding is that almost all the Phobos ejecta with radii less than 30 µm (against the values of 5 to 20 µm adopted earlier by many authors) should be rapidly lost by collisions with martian surface. The results of the paper may be the base for constructing an improved model of dust belts that presumably exist around Mars.     

Pluto—The ninth planet's golden year: A meeting held at New Mexico State University,Las Cruces,New Mexico

The simple-to-complex transition for impact craters on Mars occurs at diameters between about 3 and 8 km. Ballistically emplaced ejecta surround primarily those craters that have a simple interior morphology, whereas ejecta displaying features attributable to fluid flow are mostly restricted to complex craters. Size-dependent characteristics of 73 relatively fresh Martian craters, emphasizing the new depth/diameter (d/D) data of D. W. G. Arthur (1980, to be submitted for publication), test two hypotheses for the mode of formation of central peaks in complex craters. In particular, five features appear sequentially with increasing crater size: first flat floors (3–4 km), then central peaks and shallower depths (4–5 km), next scalloped rims (? km), and lastly terraced walls (～8 km). This relative order indicates that a shallow depth of excavation and an unspecified rebound mechanism, not centripetal collapse and deep sliding, have produced central peaks and in turn have facilitated failure of the rim. The mechanism of formation of a shallow crater remains elusive, but probably operates only at the excavation stage of impact. This interpretation is consistent with two separate and complementary lines of evidence. First, field data have documented only shallow subsurface deformation and a shallow transient cavity in complex terrestrial meteorite craters and in certain surface-burst explosion craters; thus the shallow transient cavities of complex craters never were geometrically similar to the deep cavities of simple craters. Second, the average depths of complex craters and the diameters marking the transition from simple to complex craters on Mars and on three other terrestrial planets vary inversely with gravitational acceleration at the planetary surface, g, a variable more important in the excavation of a crater than in any subsequent modification of its geometry. The new interpretation is summarized diagrammatically for complex craters on all planets.     

Phobos: Surface density of impact craters

The preliminary conclusion of the Mariner 9 Television Team that the surfaces of Phobos and Deimos are saturated with craters larger than 0.2 km in diameter is reconsidered on the basis of more extensive and uniform crater counts. For Phobos, it is verified that the surface appears saturated with craters larger than 1 km in diameter. For craters smaller than 1 km, the data points fall below the saturation curve, and it is not clear that all the departure can be explained in terms of loss of resolution. For Deimos, because of the paucity of craters visible in the Mariner 9 images, a statistically meaningful crater density curve cannot be constructed. Definitive crater density curves for subkilometer craters can only be established once additional imagery at a resolution of better than 100 m is obtained. Such imagery will be provided by the 1976 Viking Orbiters.     

Generation and emplacement of fine-grained ejecta in planetary impacts

We report here on a survey of distal fine-grained ejecta deposits on the Moon, Mars, and Venus. On all three planets, fine-grained ejecta form circular haloes that extend beyond the continuous ejecta and other types of distal deposits such as run-out lobes or ramparts. Using Earth-based radar images, we find that lunar fine-grained ejecta haloes represent meters-thick deposits with abrupt margins, and are depleted in rocks ?1 cm in diameter. Martian haloes show low nighttime thermal IR temperatures and thermal inertia, indicating the presence of fine particles estimated to range from ∼10 μm to 10 mm. Using the large sample sizes afforded by global datasets for Venus and Mars, and a complete nearside radar map for the Moon, we establish statistically robust scaling relationships between crater radius R and fine-grained ejecta run-out r^* for all three planets. On the Moon, r^* ∼ R^−0.18 for craters 5-640 km in diameter. For Venus, radar-dark haloes are larger than those on the Moon, but scale as r^* ∼ R^−0.49, consistent with ejecta entrainment in Venus’ dense atmosphere. On Mars, fine-ejecta haloes are larger than lunar haloes for a given crater size, indicating entrainment of ejecta by the atmosphere or vaporized subsurface volatiles, but scale as R^−0.13, similar to the ballistic lunar scaling. Ejecta suspension in vortices generated by passage of the ejecta curtain is predicted to result in ejecta run-out that scales with crater size as R^1/2, and the wind speeds so generated may be insufficient to transport particles at the larger end of the calculated range. The observed scaling and morphology of the low-temperature haloes leads us rather to favor winds generated by early-stage vapor plume expansion as the emplacement mechanism for low-temperature halo materials.     

The photometric functions of Phobos and Deimos I. Disc-integrated photometry

We have used the integrated brightnesses from Mariner 9 high-resolution images to determine the large phase angle (20° to 80°) phase curves of Phobos and Deimos. The derived phase coefficients are β = 0.032 ± 0.001 mag/deg for Phobos and β = 0.030 ± 0.001 mag/deg for Deimos, while the corresponding phase integrals are q_Phobos = 0.52 and q_Deimos = 0.57. The predicted intrinsic phase coefficients of the surface material are β_i = 0.019 mag/deg and β_i = 0.017 mag/deg for Phobos and Deimos, respectively. The phase curves, phase coefficients and phase integrals are typical of objects whose surface layers are dark and intricate in texture, and are consistent with the presence of a regolith on both satellites. The relative reflectance of Deimos to Phobos is 1.15±0.10. The presence of several bright patches on Deimos could account for this slight difference in average reflectance.     

The photometric functions of Phobos and Deimos. III. Surface photometry of Phobos

At least three large areas on the surface of Phobos are covered by a dark material of complex texture which scatters light according to the Hapke-Irvine Law. The average 20° to 80° intrinsic and disc-integrated phase coefficients of this material are β_i = 0.020 ± 0.001 mag/deg and β = 0.033 mag/deg, respectively. These values are slightly greater than the values found for Deimos in Paper II (preceding article). On the largest scale the surface of Phobos is rougher than the surface of Deimos, perhaps accounting for the slightly greater phase coefficients. Contrary to the situation on Deimos, no definite regions of intrinsically brighter material are apparent on Phobos. This difference could account for the slightly lower average reflectance of Phobos relative to Deimos. No evidence for large exposures of solid rock has been found in the three areas studied.     

Comparison of the crater morphology-size relationship for Mars,Moon, and Mercury

New central peak-crater size data for Mars shows that a higher percentage of relatively unmodified Martian craters have central peaks than do fresh lunar craters below a diameter of 30 km. For example, in the diameter range 10 to 20 km, 60% of studied Martian craters have central peaks compared to 26% for the Moon. Gault et al. (1975, J. Geophys. Res.80, 2444–2460) have demonstrated that central peaks occur in smaller craters on Mercury than on the Moon, and that this effect is due to the different gravity fields in which the craters formed. Similar differences when comparing Mars and the Moon show that gravity has affected the diameter at which central peaks form on Mars. Erosion on Mars, therefore, does not completely mask differences in crater interior structure that are caused by differences in gravity. Effects of Mars' higher surface gravity when compared to the Moon are not detected when comparing terrace and crater shape data. The morphology-crater size statistics also show that a full range of crater shapes occur on Mars, and craters tend to become more morphologically complex with increasing diameter. Comparisons of Martian and Mercurian crater data show differences which may be related to the greater efficacy of erosion on Mars.     

Spectral heterogeneity on Phobos and Deimos: HiRISE observations and comparisons to Mars Pathfinder results

The High-Resolution Imaging Science Experiment (HiRISE) onboard Mars Reconnaissance Orbiter (MRO) has been used to observe Phobos and Deimos at spatial scales of around 6 and 20 m/px, respectively. HiRISE (McEwen et al., JGR, 112, CiteID E05S02, DOI: 10.1029/2005JE002605, 2007) has provided, for the first time, high-resolution colour images of the surfaces of the Martian moons. When processed, by the production of colour ratio images for example, the data show considerable small-scale heterogeneity, which might be attributable to fresh impacts exposing different materials otherwise largely hidden by a homogenous regolith. The bluer material that is draped over the south-eastern rim of the largest crater on Phobos, Stickney, has been perforated by an impact to reveal redder material and must therefore be relatively thin. A fresh impact with dark crater rays has been identified. Previously identified mass-wasting features in Stickney and Limtoc craters stand out strongly in colour. The interior deposits in Stickney appear more inhomogeneous than previously suspected. Several other local colour variations are also evident.Deimos is more uniform in colour but does show some small-scale inhomogeneity. The bright “streamers” (Thomas et al., Icarus, 123, 536–556,1996) are relatively blue. One crater to the south-west of Voltaire and its surroundings appear quite strongly reddened with respect to the rest of the surface. The reddening of the surroundings may be the result of ejecta from this impact.The spectral gradients at optical wavelengths observed for both Phobos and Deimos are quantitatively in good agreement with those found by unresolved photometric observations made by the Imager for Mars Pathfinder (IMP; Thomas et al., JGR, 104, 9055–9068, 1999). The spectral gradients of the blue and red units on Phobos bracket the results from IMP.     

A model of Phobos

A model of the Martian satellite Phobos was constructed at a scale of 1:60 000 using 25 Mariner 9 photorecords and a solar-simulation technique. Measurements of the crater diameters D, depths d, ratios $\text{d}\text{D}$, longitude and latitude locations of the centers, IAU designations, crater shapes, and rim class are given in a catalog of 260 depressions. An open-ended indexing of the craters is based on their locations by octant and diameter magnitude. Six craters were found with sharply defined rims. At least 28 craters have raised rims. The range of the $\text{d}\text{D}$ ratios is from 0.002 to 0.26, with a mean $\text{d}\text{D}$ of 0.10. The mean diameter of Stickney is interpreted to be 11.1 km, its minimum 9.6 km, and the diameter of Hall 5.9 km. A 100-m contour-interval topographic map has been drawn from measurements of the model. This is rendered on an elliptical form of a Lambert equal-area polar projection. The topographic map made it possible to estimate vector lengths from the center of Phobos to vertices on a 6-frequency octahedron that fits the sattelite. A mean radius of 11.0 km results from averaging the vector lengths to the 146 well-distributed vertices of the polyhedron. A volume of 5620 km³ is deduced.     

Interplanet variations in scale of crater morphology—Earth,Mars, Moon

Craters on the Earth, Mars, and the Moon show a spectrum of morphologies with diameter increasing from simple, bowl-shaped craters through craters with increasingly complex central peaks, to craters with “peak rings” and basins with multiple concentric scarps. In each category there is a range of diameters, centered around a characteristic diameter, D_c. It is found that D_c decreases as the size of the planet increases. Several possible explanations are considered. It is suggested that the effect results from a gravity scaling law derived here and having approximately from the D_c∂ 1/g^1.25, where g is the surface gravity. All geological structures in which gravity is the dominant parameter affecting the morphology should follow such a law.     

Impact and explosion crater ejecta,fragment size,and velocity

A model was developed for the mass distribution of fragments that are ejected at a given velocity for impact and explosion craters. The model is semiempirical in nature and is derived from (1) numerical calculations of cratering and the resultant mass versus ejection velocity, (2) observed ejecta blanket particle size distributions, (3) an empirical relationships between maximum ejecta fragment size and crater diameter, (4) measurements of maximum ejecta size versus ejecta velocity, and (5) an assumption on the functional form for the distribution of fragments ejected at a given velocity. This model implies that for planetary impacts into competent rock, the distribution of fragments ejected at a given velocity is broad; e.g., 68% of the mass of the ejecta at a given velocity contains fragments having a mass less than 0.1 times a mass of the largest fragment moving at that velocity. Using this model, we have calculated the largest fragment that can be ejected from asteroids, the Moon, Mars, and Earth as a function of crater diameter. The model is unfortunately dependent on the size-dependent ejection velocity limit for which only limited data are presently available from photography of high explosive-induced rock ejecta. Upon formation of a 50-km-diameter crater on an atmosphereless planet having the planetary gravity and radius of the Moon, Mars, and Earth, fragments having a maximum mean diameter of ≈30, 22, and 17 m could be launched to escape velocity in the ejecta cloud. In addition, we have calculated the internal energy of ejecta versus ejecta velocity. The internal energy of fragments having velocities exceeding the escape velocity of the moon (～2.4 km/sec) will exceed the energy required for incipient melting for solid silicates and thus, the fragments ejected from Mars and the Earth would be melted.     

The rayed crater Zunil and interpretations of small impact craters on Mars

A 10-km diameter crater named Zunil in the Cerberus Plains of Mars created ∼¹⁰⁷ secondary craters 10 to 200 m in diameter. Many of these secondary craters are concentrated in radial streaks that extend up to 1600 km from the primary crater, identical to lunar rays. Most of the larger Zunil secondaries are distinctive in both visible and thermal infrared imaging. MOC images of the secondary craters show sharp rims and bright ejecta and rays, but the craters are shallow and often noncircular, as expected for relatively low-velocity impacts. About 80% of the impact craters superimposed over the youngest surfaces in the Cerberus Plains, such as Athabasca Valles, have the distinctive characteristics of Zunil secondaries. We have not identified any other large (?10 km diameter) impact crater on Mars with such distinctive rays of young secondary craters, so the age of the crater may be less than a few Ma. Zunil formed in the apparently youngest (least cratered) large-scale lava plains on Mars, and may be an excellent example of how spallation of a competent surface layer can produce high-velocity ejecta (Melosh, 1984, Impact ejection, spallation, and the origin of meteorites, Icarus 59, 234-260). It could be the source crater for some of the basaltic shergottites, consistent with their crystallization and ejection ages, composition, and the fact that Zunil produced abundant high-velocity ejecta fragments. A 3D hydrodynamic simulation of the impact event produced 10¹⁰ rock fragments ?10 cm diameter, leading to up to 10⁹ secondary craters ?10 m diameter. Nearly all of the simulated secondary craters larger than 50 m are within 800 km of the impact site but the more abundant smaller (10-50 m) craters extend out to 3500 km. If Zunil is representative of large impact events on Mars, then secondaries should be more abundant than primaries at diameters a factor of ∼1000 smaller than that of the largest primary crater that contributed secondaries. As a result, most small craters on Mars could be secondaries. Depth/diameter ratios of 1300 small craters (10-500 m diameter) in Isidis Planitia and Gusev crater have a mean value of 0.08; the freshest of these craters give a ratio of 0.11, identical to that of fresh secondary craters on the Moon (Pike and Wilhelms, 1978, Secondary-impact craters on the Moon: topographic form and geologic process, Lunar Planet. Sci. IX, 907-909) and significantly less than the value of ∼0.2 or more expected for fresh primary craters of this size range. Several observations suggest that the production functions of Hartmann and Neukum (2001, Cratering chronology and the evolution of Mars, Space Sci. Rev. 96, 165-194) predict too many primary craters smaller than a few hundred meters in diameter. Fewer small, high-velocity impacts may explain why there appears to be little impact regolith over Amazonian terrains. Martian terrains dated by small craters could be older than reported in recent publications.     

The equilibrium figures of Phobos and other small bodies

The shape of a close planetary satellite is distorted from a self-gravitating sphere into a triaxial ellipsoid maintained by tidal and centrifugal forces. Using the family of Roche ellipsoids calculated by Chandrasekhar, it should be possible in some cases to determine the density of an inner satellite by an accurate measurement of its shape alone. The equilibrium figure of Phobos is expected to be the most extreme of any satellite. The shape of Phobos as observed by Mariner 9 approaches but appears not to be a Roche ellipsoid, although the uncertainties of measurement remain too large to exclude the possibility. In any case, Phobos is so small that even the low mechanical strength of an impact-compressed regolith is sufficient to maintain substantial departures from the equipotential figure. If larger close satellites, particularly Amalthea, are found to be Roche ellipsoids, their densities can be estimated immediately from the data presented.Asteroids of size comparable to Phobos and Deimos appear to have more irregular shapes than the Martian satellites. This may reflect the absence of a deep regolith on those asteroids due to the low effective escape velocity for impact ejecta. For Phobos and Deimos, on the other hand, ejecta will tend to remain in orbit about Mars until swept up again by the satellite, contributing to a deeper equilibrium layer of debris.     

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.     

Mariner 9 polarimetry of Phobos and Deimos

We have used the Mariner 9 A-camera system to measure the polarization (P) of Phobos and Deimos at large phase angles (α). For Deimos, P = +22 ± 4% at α = 74°; for Phobos P = +20.5 ± 4% at α = 77°, and P = +24.5 ± 4% at α = 81°. These measurements refer to orange light at about 0.57 μm. A comparison with laboratory measurements of powdered rock samples indicates that the observations are consistent with the presence of regoliths on the satellites.