How much material do the radar-bright craters at the Mercurian poles contain?

The depth-to-diameter (d/D) ratios were determined for 12 craters located near the Mercurian north pole that were identified by Harmon et al. (2001, Icarus 149) as having strong depolarized radar echos. We find that the mean d/D value of these radar-bright craters is the mean d/D value of the general population of non-radar-bright craters in the surrounding north polar region. Previous studies, however, show no difference between d/D values of Mercurian polar and equatorial crater populations, suggesting that no terrain softening which could modify crater structure exists at the Mercurian poles (Barlow et al., 1999, 194, Icarus 141). Thus, the change in d/D is governed by a change in crater depth, probably due to deposition of material inside the crater. The volume of infilling material, including volatiles, in the radar-bright craters is significantly greater than predicted by proposed mechanisms for the emplacement of either water ice or sulfur.     

Can spatial statistics help decipher impact crater saturation?

Impact crater saturation equilibrium is a state where a surface is so densely cratered that a new crater cannot form without removing older craters and the observed crater density is in (quasi-)equilibrium. Whether densely cratered surfaces throughout the solar system are saturated for large, kilometer-sized craters has been debated for decades. This work explores if spatial statistics can provide insight if these crater distributions are in saturation equilibrium. The supposition is that crater distributions become more spatially uniform (more evenly spaced) as they reach saturation (Squyres et al. 1997 ). A numerical simulation of crater saturation is combined with observations of cratered terrains throughout the solar system to assess the utility of using spatial statistics. The numerical simulations examine spatial statistics and saturation equilibrium for crater distributions for various input population size-frequency distribution (SFD) slopes, along with a range in the effective crater erasure size, effectiveness of smaller craters erasing the rims of larger craters, and the amount of rim needed to recognize a crater. Simulations show that saturated terrains do become more spatially uniform, and that the degree of uniformity appears to be most dependent on the input SFD slope. When simulation results are compared to observed crater distributions, I find that large, kilometer-sized craters on densely cratered terrains throughout the solar system are likely in saturation equilibrium.     

Crater modification and geologic activity in Enceladus' heavily cratered plains: Evidence from the impact crater distribution

Although we can observe current activity on Saturn's satellite Enceladus with Cassini, insight into past activity is best achieved (for now) through studying the impact crater distributions. Furthermore, approximation of terrain ages can only be attained through calculations using crater densities and estimations of impact rates in the saturnian system. Here we focus on what the impact crater distribution in Enceladus' heavily cratered plains can tell us about Enceladus' geologic history. We use Cassini ISS images to count craters in the heavily cratered plains on Enceladus, along with Rhea, Dione, Tethys and Mimas as references, to develop and compare their size-frequency distributions. Comparisons of our counts show that Enceladus' cratered plains distribution is unique in that it appears to have a relative deficiency of craters for diameters ?2 km and ?6 km compared to the other satellites' heavily cratered plains. Our data also indicates that the impact crater density within the cratered plains changes with latitude. Specifically, both the north and south mid-latitude regions have approximately three times higher density than the equatorial region. We hypothesize that the “missing” small and large craters in Enceladus' cratered plains is due to a combination of viscous relaxation of the larger craters, and burial of the relaxed large craters and small craters by south polar plume and possibly E-ring material. We also conclude that the spatial density distribution is not consistent with recent polar wander.     

The effect of target properties on crater morphology: Comparison of central peak craters on the Moon and Ganymede

Abstract— We examine the morphology of central peak craters on the Moon and Ganymede in order to investigate differences in the near‐surface properties of these bodies. We have extracted topographic profiles across craters on Ganymede using Galileo images, and use these data to compile scaling trends. Comparisons between lunar and Ganymede craters show that crater depth, wall slope and amount of central uplift are all affected by material properties. We observe no major differences between similar‐sized craters in the dark and bright terrain of Ganymede, suggesting that dark terrain does not contain enough silicate material to significantly increase the strength of the surface ice. Below crater diameters of ?12 km, central peak craters on Ganymede and simple craters on the Moon have similar rim heights, indicating comparable amounts of rim collapse. This suggests that the formation of central peaks at smaller crater diameters on Ganymede than the Moon is dominated by enhanced central floor uplift rather than rim collapse. Crater wall slope trends are similar on the Moon and Ganymede, indicating that there is a similar trend in material weakening with increasing crater size, and possibly that the mechanism of weakening during impact is analogous in icy and rocky targets. We have run a suite of numerical models to simulate the formation of central peak craters on Ganymede and the Moon. Our modeling shows that the same styles of strength model can be applied to ice and rock, and that the strength model parameters do not differ significantly between materials.     

Investigating target versus impactor influences on Martian crater morphology at the simple‐complex transition

Comparing craters of identical diameter on a planet is an empirical method of studying the effects of different target and impactor properties while holding total impact energy nearly constant. We have analyzed the Martian crater population within a narrow diameter range (7 km < crater diameter < 9 km) at the simple‐complex crater transition using three approaches. We looked for correlations of morphology with surface geology using a global crater database and global geologic map. We examined selected regions in detail with high‐resolution images to further understand the relationship between crater morphology and bulk target properties. Finally, we examined craters in close proximity to each other in order to hold target properties constant, so that we could isolate impactor effects on crater morphology. We found a strong correlation between target properties and interior crater morphology, and we found little evidence that impactor properties (other than impact angle) affect crater appearance. Central uplift and wall slumping are enhanced for less consolidated targets. Layered targets affected both the excavation and modification stages of complex crater formation; the resulting craters have pseudoterraces, flat floors, and central pits.     

Degradation of mid-latitude craters on Mars

Recent geomorphic, remote sensing, and atmospheric modeling studies have shown evidence for abundant ground ice deposits in the martian mid-latitudes. Numerous potential water/ice-rich flow features have been identified in craters in these regions, including arcuate ridges, gullies, and small flow lobes. Previous studies (such as in Newton Basin) have shown that arcuate ridges and gullies are mainly found in small craters (∼2-30 km in diameter). These features are located on both pole-facing and equator-facing crater walls, and their orientations have been found to be dependent on latitude. We have conducted surveys of craters >20 km in diameter in two mid-latitude regions, one in the northern hemisphere in Arabia Terra, and one in the southern hemisphere east of Hellas basin. In these regions, prominent lobes, potentially ice-rich, are commonly found on the walls of craters with diameters between ∼20-100 km. Additional water/ice-rich features such as channels, valleys, alcoves, and debris aprons have also been found in association with crater walls. In the eastern Hellas study region, channels were found to be located primarily on pole-facing walls, whereas valleys and alcoves were found primarily on equator-facing walls. In the Arabia Terra study region, these preferences are less distinct. In both study regions, lobate flows, gullies, and arcuate ridges were found to have pole-facing orientation preferences at latitudes below 45° and equator-facing orientation preferences above 45°, similar to preferences previously found for gullies and arcuate ridges in smaller craters. Interrelations between the features suggest they all formed from the mobilization of accumulated ice-rich materials. The dependencies of orientations on latitude suggest a relationship to differences in total solar insolation along the crater walls. Differences in slope of the crater wall, differences in total solar insolation with respect to wall orientation, and variations in topography along the crater rim can explain the variability in morphology of the features studied. The formation and evolution of these landforms may best be explained by multiple cycles of deposition of ice-rich material during periods of high obliquity and subsequent modification and transport of these materials down crater walls.     

The geology of the Viking Lander 2 site revisited

Reevaluating the geologic history of the prior Mars landing sites provides important ground truth for recent and ongoing orbital missions. At the Viking 2 Lander (VL2) site, topographic measurements of relict landforms indicate that at least 100 m of sedimentary mantle material has been stripped away. The observed paucity of impact craters <100 m in diameter suggests that resurfacing processes (likely in the form of the recent deposition and removal of thin 1-10 m mantle layers) continue up to the present. A dearth of craters in the 100-500 m diameter range, however, also necessitates erosion of a thicker mantle layer. Partially inverted chains of secondary craters from nearby Mie Crater indicate that the mantle was already in place when the impact occurred. The density of craters superposed on Mie ejecta is consistent with a Late Hesperian age and provides a minimum age constraint for the mantle's emplacement. The thermophysical properties of the surface around VL2 as observed with Thermal Emission Imaging System (THEMIS) data indicate that the landing site occurs in an intracrater region that may typify mid to high northern latitude sites. Elevated thermal inertias of a pedestal crater superposed atop a larger pedestal crater suggest that rocky or indurated material can be created by impacts into sedimentary targets. Rock abundances at VL2 are consistent with the addition of impact-emplaced material from the missing small impact crater population documented in this study. Thus, the VL2 site may be a reasonable proxy for the landscape expected at the upcoming Phoenix Lander site.     

Crater geometry and ejecta thickness of the Martian impact crater Tooting

Abstract— We use Mars Orbiter Laser Altimeter (MOLA) topographic data and Thermal Emission Imaging System (THEMIS) visible (VIS) images to study the cavity and the ejecta blanket of a very fresh Martian impact crater ?29 km in diameter, with the provisional International Astronomical Union (IAU) name Tooting crater. This crater is very young, as demonstrated by the large depth/diameter ratio (0.065), impact melt preserved on the walls and floor, an extensive secondary crater field, and only 13 superposed impact craters (all 54 to 234 meters in diameter) on the ?8120 km² ejecta blanket. Because the pre‐impact terrain was essentially flat, we can measure the volume of the crater cavity and ejecta deposits. Tooting crater has a rim height that has >500 m variation around the rim crest and a very large central peak (1052 m high and >9 km wide). Crater cavity volume (i.e., volume below the pre‐impact terrain) is ?380 km³ the volume of materials above the pre‐impact terrain is ?425 km³. The ejecta thickness is often very thin (<20 m) throughout much of the ejecta blanket. There is a pronounced asymmetry in the ejecta blanket, suggestive of an oblique impact, which has resulted in up to ?100 m of additional ejecta thickness being deposited down‐range compared to the up‐range value at the same radial distance from the rim crest. Distal ramparts are 60 to 125 m high, comparable to the heights of ramparts measured at other multi‐layered ejecta craters. Tooting crater serves as a fresh end‐member for the large impact craters on Mars formed in volcanic materials, and as such may be useful for comparison to fresh craters in other target materials.     

Evidence for subsurface water ice in Korolev crater, Mars

Following the work of Kieffer and Titus (2001, Icarus 154, 162-180), we present results of thermal IR observations of Korolev crater, located at ∼73° latitude in the martian northern polar region. Similar to techniques employed by Titus et al. (2003, Science 299, 1048-1050), we use infrared images from the Thermal Emission Imaging System (THEMIS) aboard Mars Odyssey to identify several regions within the crater basin with distinct thermal properties that correlate with topography. The THEMIS results show these regions exhibit temperature variations, spatially within the crater and throughout the martian year. In addition to the variations identified in the THEMIS observations, Mars Global Surveyor Thermal Emission Spectrometer (TES) observations show differences in albedo and temperature of these regions on both daily and seasonal cycles. Modeling annual temperature variations of the surface, we use TES observations to examine the thermal properties of these regions. This analysis reveals the crater interior deposits are likely thick layers (several meters) of high thermal inertia material (water ice, or extremely ice-rich regolith). Spatial variations of the physical properties of these regions are likely due to topography and possibly variations in the subsurface material itself. The nature of these deposits may help constrain polar processes, as well as provide context for the polar lander mission, Phoenix.     

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 conceptual model for explanation of Albedo changes in Martian craters

We propose a conceptual model to interpret AM/PM high albedo events (HAEs) in crater interiors at the Martian seasonal polar caps. This model consists of two components: (1) a relatively permanent high-albedo water–ice body exposed in a crater interior and (2) a variable crater albedo in response to aerosol optical depth, dust contamination, and H₂O/CO₂ frost deposits or sublimes in four phases, based on temperature and solar longitude changes. Two craters (Korolev crater of fully exposed water–ice layer and ‘Louth’ crater of partially exposed water–ice layer) are used to demonstrate the model. This model explains the HAEs and their seasonal changes and suggests that many crater-like features formed in the last episodic advance of the polar ice cap in the last high obliquity period should have water–ice exposed or covered. For the AM-only HAEs craters, there seems no need of a water–ice layer to be fully exposed, but a subsurface water–ice layer (or ice-rich regolith) is a necessary condition.     

Martian cratering 11. Utilizing decameter scale crater populations to study Martian history

New information has been obtained in recent years regarding formation rates and the production size‐frequency distribution (PSFD) of decameter‐scale primary Martian craters formed during recent orbiter missions. Here we compare the PSFD of the currently forming small primaries (P) with new data on the PSFD of the total small crater population that includes primaries and field secondaries (P + fS), which represents an average over longer time periods. The two data sets, if used in a combined manner, have extraordinary potential for clarifying not only the evolutionary history and resurfacing episodes of small Martian geological formations (as small as one or few km²) but also possible episodes of recent climatic change. In response to recent discussions of statistical methodologies, we point out that crater counts do not produce idealized statistics, and that inherent uncertainties limit improvements that can be made by more sophisticated statistical analyses. We propose three mutually supportive procedures for interpreting crater counts of small craters in this context. Applications of these procedures support suggestions that topographic features in upper meters of mid‐latitude ice‐rich areas date only from the last few periods of extreme Martian obliquity, and associated predicted climate excursions.     

Viscous relaxation of impact craters on icy planetary surfaces: Determination of viscosity variation with depth

Spacecraft images show that the icy Galilean satellites have surfaces with very low topographic relief. Impact craters on Ganymede and Callisto are anomalously shallow and are characterized by sharp well-defined rims and domed floors. These morphological characteristics can be explained by viscous relaxation of topography on an icy crust in which the viscosity is uniform or decreases with depth. Under these conditions, large craters relax more rapidly than small craters, therefore explaining a possible underabundance of large craters. Viscous relaxation on an icy crust that is thin compared to the crater diameter or on a thick icy crust in which viscosity increases with depth could not produce this crater morphology and would result in the more rapid relaxation of small craters rather than large craters. The results of this study suggest that more detailed analysis of relaxing impact crater morphology may resolve the rate of viscosity decrease with depth and so provide evidence on the interior thermal evolution of icy planetary bodies.     

The geology of 433 Eros

Abstract— The global high‐resolution imaging of asteroid 433 Eros by the Near‐Earth Asteroid Rendezvous (NEAR) Shoemaker spacecraft has made it possible to develop the first comprehensive picture of the geology of a small S‐type asteroid. Eros displays a variety of surface features, and evidence of a substantial regolith. Large scale facets, grooves, and ridges indicate the presence of at least one global planar structure. Directional and superposition relations of smaller structural features suggest that fracturing has occurred throughout the object. As with other small objects, impact craters dominate the overall shape as well as the small‐scale topography of Eros. Depth/diameter ratios of craters on Eros average ～0.13, but the freshest craters approach lunar values of ～0.2. Ejecta block production from craters is highly variable; the majority of large blocks appear to have originated from one 7.6 km crater (Shoemaker). The interior morphology of craters does not reveal the influence of discrete mechanical boundaries at depth in the manner of craters formed on lunar mare regolith and on some parts of Phobos. This lack of mechanical boundaries, and the abundant evidence of regolith in nearly every high‐resolution image, suggests a gradation in the porosity and fracturing with depth. The density of small craters is deficient at sizes below ～200 m relative to predicted slopes of empirical saturation. This characteristic, which is also found on parts of Phobos and lunar highland areas, probably results from the efficient obliteration of small craters on a body with significant topographic slopes and a thick regolith. Eros displays a variety of regolith features, such as debris aprons, fine‐grained “ponded” deposits, talus cones, and bright and dark streamers on steep slopes indicative of efficient downslope movement of regolith. These processes serve to mix materials in the upper loose fragmental portion of the asteroid (regolith). In the instance of “ponded” materials and crater wall deposits, there is evidence of processes that segregate finer materials into discrete deposits. The NEAR observations have shown us that surface processes on small asteroids can be very complex and result in a wide variety of morphologic features and landforms that today seem exotic. Future missions to comets and asteroids will surely reveal still as yet unseen processes as well as give context to those discovered by the NEAR Shoemaker spacecraft.     

Geology of the lunar farside crater Necho

The lunar farside crater Necho (30 km diameter) displays intricate morphological and structural characteristics. The highland setting provides a complex impact site when compared with the relatively uniform setting of mare craters. Therefore, the effects of pre-impact topography and structure play a dominant role in Necho's formation and modification. Necho's bright ejecta, extensive rays, fresh morphology, and lack of superposed craters indicate that it is extremely young. The asymmetric distribution of ejecta materials may be due to substrate effects, topographic shalowing, or oblique impact.Necho's interior is divided into five physiographic units based on morphologic differences: three floor units (Necho does not display a true flat floor), one hilly central unit, and the wall unit which includes terraces and smooth walls. The interior of the crater also exhibits an unusual asymmetry in the prevalence of terraced units on the western wall. Interior morphology and terrace orientations are probably the result of pre-impact effects. Structural and topographic orientations associated with three large pre-existing degraded craters dominate the impact site.     

Impact craters on Titan

Five certain impact craters and 44 additional nearly certain and probable ones have been identified on the 22% of Titan’s surface imaged by Cassini’s high-resolution radar through December 2007. The certain craters have morphologies similar to impact craters on rocky planets, as well as two with radar bright, jagged rims. The less certain craters often appear to be eroded versions of the certain ones. Titan’s craters are modified by a variety of processes including fluvial erosion, mass wasting, burial by dunes and submergence in seas, but there is no compelling evidence of isostatic adjustments as on other icy moons, nor draping by thick atmospheric deposits. The paucity of craters implies that Titan’s surface is quite young, but the modeled age depends on which published crater production rate is assumed. Using the model of Artemieva and Lunine (2005) suggests that craters with diameters smaller than about 35 km are younger than 200 million years old, and larger craters are older. Craters are not distributed uniformly; Xanadu has a crater density 2-9 times greater than the rest of Titan, and the density on equatorial dune areas is much lower than average. There is a small excess of craters on the leading hemisphere, and craters are deficient in the north polar region compared to the rest of the world. The youthful age of Titan overall, and the various erosional states of its likely impact craters, demonstrate that dynamic processes have destroyed most of the early history of the moon, and that multiple processes continue to strongly modify its surface. The existence of 24 possible impact craters with diameters less than 20 km appears consistent with the Ivanov, Basilevsky and Neukum (1997) model of the effectiveness of Titan’s atmosphere in destroying most but not all small projectiles.     

Hydrothermal dynamics in a CM‐based model of Ceres

A 2‐D numerical study of the evolution of Ceres from a “frozen mudball” to the present era emphasizes the importance of hydrothermal processes. Particulates released as the “frozen mudball” thaws settle to form a roughly 290 km radius core. Hydrothermal flow is driven by radiogenic heating and serpentinization. Both salt‐free and brine fluids are considered. Our modeling suggests that Ceres’s core has been warm over most of its history and is still above freezing, and convective processes are active in core and mantle to the present. The addition of low eutectic solutes greatly expands the region of active convection. A global muddy ocean persists for the first 3 Gyr, and at present, there may be several regional mud seas buried under a frozen crust. Transport of interior material to the near surface occurs throughout our model's history. Eutectic brines drive convective flow to near the surface, even breaching the surface in isolated regions, on the order of 30 km in width, similar in size to some mounds detected using the Dawn visible imaging camera (Sizemore et al. 2015). Surface features such as the bright spot in Occator crater and Ahuna Mons could be the result of eutectic plumes. The CM‐based model density profile is within 10% of Ermakov et al.'s ( 2017 ) results. The model mud mantle has a roughly 42:58 volumetric partitioning of H₂O to rock. Our mud model is consistent with the absence of large craters (Marchi et al. 2016 ) and an internal viscosity decreasing with depth (Fu et al. 2017 ).     

The role of arcuate ridges and gullies in the degradation of craters in the Newton Basin region of Mars

A survey of craters in the vicinity of Newton Basin, using high-resolution images from Mars Global Surveyor and Mars Odyssey, was conducted to find and analyze examples of gullies and arcuate ridges and assess their implications for impact crater degradation processes. In the Phaethontis Quadrangle (MC-24), we identified 225 craters that contain these features. Of these, 188 had gullies on some portion of their walls, 118 had arcuate ridges at the bases of the crater walls, and 104 contained both features, typically on the same crater wall. A major result is that the pole-facing or equator-facing orientation of these features is latitude dependent. At latitudes >44° S, equator-facing orientations for both ridges and gullies are prevalent, but at latitudes <44° S, pole-facing orientations are prevalent. The gullies and arcuate ridges typically occupy craters between ∼2 and 30 km in diameter, at elevations between −1 and 3 km. Mars Orbiter Laser Altimeter (MOLA) elevation profiles indicate that most craters with pole-facing arcuate ridges have floors sloping downward from the pole-facing wall, and some of these craters show asymmetry in crater rim heights, with lower pole-facing rims. These patterns suggest viscous flow of ice-rich materials preferentially away from gullied crater walls. Clear associations exist between gullies and arcuate ridges, including (a) geometric congruence between alcoves and sinuous arcs of arcuate ridges and (b) backfilling of arcuate ridges by debris aprons associated with gully systems. Chronologic studies suggest that gullied walls and patterned crater floor deposits have ages corresponding to the last few high obliquity cycles. Our data appear consistent with the hypothesis that these features are associated with periods of ice deposition and subsequent erosion associated with obliquity excursions within the last few tens of millions of years. Arcuate ridges may form from cycles of activity that also involve gully formation, and the ridges may be in part due to mass-wasted, ice-rich material transported downslope from the alcoves, which then interacts with previously emplaced floor deposits. Most observed gullies may be late-stage features in a degradational cycle that may have occurred many times on a given crater wall.     

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.     

The surface of Io: Geologic units,morphology, and tectonics

A prelimanary geologic map, representing 26.5% of the surface of Io, has been compiled using best-resolution (0.5 to 5 km/line pair) Voyager 1 images and (as a base) a preliminary pictorial map of Io. Nine volcanic units are identified, including materials of mountains (1.9% of total area), plains (49.6%), flows (31.1%), cones (0.1%), and crater vents (4.0%), in addition to seven types of structural features. Photogeologic evidence indicates a dominantly silicate composition for the mountain material, which supports heights of at least 9 ± 1 km. Sulfur flows of diverse viscosity and sulfur-silicate mixtures are thought to compose the pervasive plains. Pit crater and shield crater vent wall scarps reach heights of 2 km and layered plains boundary scarps have estimated heights of 150 to 1700 m; such scarps indicate a material with considerable strenght. A cumulative, volcanic crater size-frequency distribution plot has been prepared using 170 mapped Ionian vents with diameters > 14 km; the shape and slope of the curve are like those for impact craters on other bodies in the solar system, attesting to a similar nonrandom distribution to crater diameters and a surplus of small craters. Io's equatorial zone has six times the number of vents per unit area as the south polar zone. No craters of unequivocal impact origin have been identified on Io to date. A total of 151 lineaments and grabens are recognized with four dominant azimuthal trends forming two nearly orthogonal sets spaces 110° apart (N 85° E, N 25° W and N 45° E, N 55°W). The mapped area lies within the longitudinal zone (250 to 323°) of least-abundant SO₂ frost, indicating that other sulfurous components dominate the upper surface layers in this area.