Using the Gamma-Ray Spectrometer on the Mars Odyssey, we have identified two regions near the poles that are enriched in hydrogen. The data indicate the presence of a subsurface layer enriched in hydrogen overlain by a hydrogen-poor layer. The thickness of the upper layer decreases with decreasing distance to the pole, ranging from a column density of about 150 grams per square centimeter at -42 degrees latitude to about 40 grams per square centimeter at -77 degrees. The hydrogen-rich regions correlate with regions of predicted ice stability. We suggest that the host of the hydrogen in the subsurface layer is ice, which constitutes 35 +/- 15% of the layer by weight.

Coronae on Venus range from 60 to over 2000 km across and are characterized by a complex range of morphologies. The annuli around coronae range from about 10 to 150 km across and have tectonic features ranging from extensional to compressional to a combination of both. Topographically, coronae are domes, plateaus, plateaus with interior lows, and rimmed depressions. A subset of features classified here as coronae corresponds to depressions and is interpreted to consist of large‐scale calderas. A number of features have been identified with many of the basic characteristics of coronae (similar interior deformation, associations with volcanism, high topography) but lacking a distinct tectonic annulus. These features tend to be somewhat smaller than coronae and may represent “failed” coronae or coronae in an early stage of evolution. The size distribution of coronae and coronalike features with maximum widths greater than about 250 km is well represented by a power law of the form N(D) = kD−α, where N is the number of coronae with maximum widths greater than D (km) and α = 3.05. The spatial distribution of coronae is not random; the features are concentrated in a few groups and along several chains. Coronae are similar in many morphologic characteristics to major volcanic shield structures and volcanic rises such as Western Eistla Regio. The largest corona, Artemis, is actually larger than several volcanic rises on Venus. Coronae and volcanic rises appear to be surface manifestations of mantle plumes. There is no evidence of any systematic variation in age along chains of coronae as occurs in hotspot chains on Earth. Instead, a number of multiple and overlapping coronae may indicate limited movement of the surface above a hotspot or mantle plume. The morphology and size distribution of coronae, highlands, and major shields suggest that mantle upwelling on Venus operates either on several spatial scales, with coronae representing smaller‐scale upflows and major volcanic rises representing larger convective upwellings, or on several temporal scales, with coronae representing shorter duration upflows and major volcanic rises representing long‐term upwellings.

We report maps of the concentrations of H, Si, Cl, K, Fe, and Th as determined by the Gamma Ray Spectrometer (GRS) on board the 2001 Mars Odyssey Mission for ±∼45° latitudes. The procedures by which the spectra are processed to yield quantitative concentrations are described in detail. The concentrations of elements determined over the locations of the various Mars landers generally agree well with the lander values except for Fe, although the mean of the GRS Fe data agrees well with that of Martian meteorites. The water‐equivalent concentration of hydrogen by mass varies from about 1.5% to 7.5% (by mass) with the most enriched areas being near Apollinaris Patera and Arabia Terra. Cl shows a distribution similar to H over the surface except that the Cl content over Medusae Fossae is much greater than elsewhere. The map of Fe shows enrichment in the northern lowlands versus the southern highlands. Silicon shows only very modest variation over the surface with mass fractions ranging from 19% to 22% over most of the planet, though a significant depletion in Si is noted in a region west of Tharsis Montes and Olympus Mons where the Si content is as low as 18%. K and Th show a very similar pattern with depletions associated with young volcanic deposits and enrichments associated with the TES Surface Type‐2 material. It is noted that there appears to be no evidence of significant globally distributed thick dust deposits of uniform composition.

We report the concentrations of K, Th, and Fe on the Martian surface, as determined by the gamma ray spectrometer onboard the 2001 Mars Odyssey spacecraft. K and Th are not uniformly distributed on Mars. K ranges from 2000 to 6000 ppm; Th ranges from 0.2 to 1 ppm. The K/Th ratio varies from 3000 to 9000, but over 95% of the surface has K/Th between 4000 and 7000. Concentrations of K and Th are generally higher than those in basaltic Martian meteorites (K = 200–2600 ppm; Th = 0.1–0.7 ppm), indicating that Martian meteorites are not representative of the bulk crust. The average K/Th in the crust is 5300, consistent with the Wänke‐Dreibus model composition for bulk silicate Mars. Fe concentrations support the idea that bulk Mars is enriched in FeO compared to Earth. The differences in K/Th and FeO between Earth and Mars are consistent with the planets accreting from narrow feeding zones. The concentration of Th on Mars does not vary as much as it does on the Moon (where it ranges from 0.1 to 12 ppm), suggesting that the primary differentiation of Mars differed from that of the Moon. If the average Th concentration (0.6 ppm) of the surface is equal to the average of the entire crust, the crust cannot be thicker than about 118 km. If the crust is about 57 km thick, as suggested by geophysical studies, then about half the Th is concentrated in the crust.

K/Th determined by the Mars Odyssey Gamma Ray Spectrometer varies by a factor of 3 on Mars (3000 to 9000), but over 95% of the surface area has K/Th between 4000 and 7000. K/Th is distinctly lower than average in some areas, including west of Olympus Mons in the Amazonis Planitia, the region around Memnonia Fossae, Chryse Planitia, southeastern Arabia Terra, Syrtis Major Planum, and northwest of Apollinaris Patera. On the other hand, K/Th is distinctly higher than average in other areas, including the central part of Valles Marineris and the surrounding highlands, and in the northern part of Hellas. The generally modest variation in K/Th may be explained by inherent variations in igneous rocks and by variations in the extent of aqueous alteration.