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1.
Building on previous studies of volcanoes around the Hellas basin with new studies of imaging (High-Resolution Stereo Camera (HRSC), Thermal Emission Imaging System (THEMIS), Mars Orbiter Camera (MOC), High-Resolution Imaging Science Experiment (HiRISE), Context Imager (CTX)), multispectral (HRSC, Observatoire pour la Minéralogie, l’Eau, les Glaces et l’Activité (OMEGA)), topographic (Mars Orbiter Laser Altimeter (MOLA)) and gravity data, we define a new Martian volcanic province as the Circum-Hellas Volcanic Province (CHVP). With an area of >2.1 million km2, it contains the six oldest central vent volcanoes on Mars, which formed after the Hellas impact basin, between 4.0 and 3.6 Ga. These volcanoes mark a transition from the flood volcanism that formed Malea Planum ～3.8 Ga, to localized edifice-building eruptions. The CHVP volcanoes have two general morphologies: (1) shield-like edifices (Tyrrhena, Hadriaca, and Amphitrites Paterae), and (2) caldera-like depressions surrounded by ridged plains (Peneus, Malea, and Pityusa Paterae). Positive gravity anomalies are found at Tyrrhena, Hadriaca, and Amphitrites, perhaps indicative of dense magma bodies below the surface. The lack of positive-relief edifices and weak gravity anomalies at Peneus, Malea, and Pityusa suggest a fundamental difference in their formation, styles of eruption, and/or compositions. The northernmost volcanoes, the ～3.7–3.9 Ga Tyrrhena and Hadriaca Paterae, have low slopes, well-channeled flanks, and smooth caldera floors (at tens of meters/pixel scale), indicative of volcanoes formed from poorly consolidated pyroclastic deposits that have been modified by fluvial and aeolian erosion and deposition. The ～3.6 Ga Amphitrites Patera also has a well-channeled flank, but it and the ～3.8 Ga Peneus Patera are dominated by scalloped and pitted terrain, pedestal and ejecta flow craters, and a general ‘softened’ appearance. This morphology is indicative not only of surface materials subjected to periglacial processes involving water ice, but also of a surface composed of easily eroded materials such as ash and dust. The southernmost volcanoes, the ～3.8 Ga Malea and Pityusa Paterae, have no channeled flanks, no scalloped and pitted terrain, and lack the ‘softened’ appearance of their surfaces, but they do contain pedestal and ejecta flow craters and large, smooth, bright plateaus in their central depressions. This morphology is indicative of a surface with not only a high water ice content, but also a more consolidated material that is less susceptible to degradation (relative to the other four volcanoes). We suggest that Malea and Pityusa (and possibly Peneus) Paterae are Martian equivalents to Earth's giant calderas (e.g., Yellowstone, Long Valley) that erupted large volumes of volcanic materials, and that Malea and Pityusa are probably composed of either lava flows or ignimbrites. HRSC and OMEGA spectral data indicate that dark gray to slightly red materials (often represented as blue or black pixels in HRSC color images), found in the patera floors and topographic lows throughout the CHVP, have a basaltic composition. A key issue is whether this dark material represents concentrations of underlying basaltic material eroded by various processes and exposed by aeolian winnowing, or if the material was transported from elsewhere on Mars by regional winds. Understanding the provenance of these dark materials may be the key to understanding the volcanic diversity of the Circum-Hellas Volcanic Province.  相似文献   

2.
B.J. Travis  J. Palguta  G. Schubert 《Icarus》2012,218(2):1006-1019
A whole-moon numerical model of Europa is developed to simulate its thermal history. The thermal evolution covers three phases: (i) an initial, roughly 0.5 Gyr-long period of radiogenic heating and differentiation, (ii) a long period from 0.5 Gyr to 4 Gyr with continuing radiogenic heating but no tidal dissipative heating (TDH), and (iii) a final period covering the last 0.5 Gyr until the present, during which TDH is active. Hydrothermal plumes develop after the initial period of heating and differentiation and transport heat and salt from Europa’s silicate mantle to its ice shell. We find that, even without TDH, vigorous hydrothermal convection in the rocky mantle can sustain flow in an ocean layer throughout Europa’s history. When TDH becomes active, the ice shell melts quickly to a thickness of about 20 km, leaving an ocean 80 km or more deep. Parameterized convection in the ice shell is non-uniform spatially, changes over time, and is tied to the deeper ocean–mantle dynamics. We also find that the dynamics are affected by salt concentrations. An initially non-uniform salt distribution retards plume penetration, but is homogenized over time by turbulent diffusion and time-dependent flow driven by initial thermal gradients. After homogenization, the uniformly distributed salt concentrations are no longer a major factor in controlling plume transport. Salt transport leads to the formation of a heterogeneous brine layer and salt inclusions at the bottom of the ice shell; the presence of salt in the ice shell could strongly influence convection in that layer.  相似文献   

3.
We have mapped the area of Isidis Planitia (1–27°N, 75–103°E) in order to assess the geologic history of this region using modern data sets such as MOLA topography and the high-resolution images provided by the HRSC, CTX, and HiRISE cameras. Results of our mapping show that the geologic history of Isidis Planitia consists of three principal episodes. (1) Impact dominated episode (Noachian, until ～3.8 Ga): During this time, the oldest materials in the study area were formed mostly by impact reworking and mass-wasting. Other processes (e.g., volcanism and fluvial/glacial activity) likely operated at this time but played a subordinate role. (2) An episode related to volcanic and fluvial/glacial activities (late Noachian–early Amazonian, ～3.8–2.8 Ga): Volcanism appears as the most important process at the beginning of this episode (～3.8–3.5 Ga) and was responsible for the formation of a large circum-Isidis volcanic province by the early Hesperian epoch. Volcanic materials covered large portions of the Isidis rim, almost completely buried the previous crater record on the floor of the Isidis basin, and probably were the major contributors to the filling of the basin. Fluvial/glacial processes prevailed closer to the end of the episode (early Hesperian–early Amazonian, ～3.5–2.8 Ga) and were responsible for widespread resurfacing in the Isidis Planitia region, mostly at ～3.1–3.4 Ga. Glaciers and/or ice sheets probably resulted in a massive glaciation of the rim and the floor of the Isidis basin. The total volume of material eroded from the Isidis rim by glacial and fluvial activity is estimated to be about 35,000–50,000 km3, which is equivalent to a composite layer about 40–60 m thick on entire floor of the basin. More important, however, is that the eroded materials were likely saturated with ice/water and could form wet deposits on the floor. (3) Wind-dominated episode (since early Amazonian, ～2.8 Ga): Wind activity dominated the later geologic history of Isidis Planitia but resulted only in minor modification of the surface.  相似文献   

4.
《Planetary and Space Science》2007,55(10):1328-1345
The planetary fourier spectrometer (PFS) for the Mars express mission (MEX) is an infrared spectrometer operating in the wavelength range from 1.2 to 45 μm by means of two spectral channels, called SWC (short wavelength channel) and LWC (long wavelength channel), covering, respectively, 1.2–5.5 and 5.5–45 μm.The middle-spring Martian north polar cap (Ls∼40°) has been observed by PFS/MEX in illuminated conditions during orbit 452. The SWC spectra are here used to study the cap composition in terms of CO2 ice, H2O ice and dust content. Significant spectral variation is noted in the cap interior, and regions of varying CO2 ice grain sizes, water frost abundance, CO2 ice cover and dust contamination can be distinguished. In addition, we correlate the infrared spectra with an image acquired during the same orbit by the OMEGA imaging spectrometer and with the altimetry from MOLA data. Many of the spectra variations correlate with heterogeneities noted in the image, although significant spectral variations are not discernible in the visible. The data have been divided into five regions with different latitude ranges and strong similarities in the spectra, and then averaged. Bi-directional reflectance models have been run with the appropriate lighting geometry and used to fit the observed data, allowing for CO2 ice and H2O ice grain sizes, dust and H2O ice contaminations in the form of intimate granular mixtures and spatial mixtures.A wide annulus of dusty water ice surrounds the recessing CO2 seasonal cap. The inner cap exhibits a layered structure with a thin CO2 layer with varying concentrations of dark dust, on top of an H2O ice underneath ground. In the best-fits, the ices beneath the top layer have been considered as spatial mixtures. The results are still very good everywhere in the spectral range, except where the CO2 ice absorption coefficients are such that even a thin layer is enough to totally absorb the incoming radiation (i.e. the band is saturated). This only happens around 3800 cm−1, inside the strong 2.7-μm CO2 ice absorption band. The effect of finite snow depth has been investigated through a layered albedo model. The thickness of the CO2 ice deposits increases with latitude, ranging from 0.5–1 g cm−2 within region II to 60–80 g cm−2 within the highest-latitude (up to 84°N) region V.Region I is at the cap edge and extends from 65°N to 72°N latitude. No CO2 ice is present in this region, which consists of relatively large grains of water ice (20 μm), highly contaminated by dust (0.15 wt%). The adjacent region II is a narrow region [76–79°N] right at the edge of the north residual polar cap. This region is very distinct in the OMEGA image, where it appears to surround the whole residual cap. The CO2 ice features are barely visible in these spectra, except for the strong saturated 2.7 μm band. It basically consists of a thin layer of 5-mm CO2 ice on top of an H2O ice layer with the same composition as region I. A third interesting region III is found all along the shoulder of the residual cap [79–81°N]. It extends over 1.5 km in altitude and over only 2° of latitude and consists of CO2 ice with a large dust content. It is an admixture of CO2 ice (3–4 mm), with several tens of ppm by mass of water ice and more than 2 ppt by mass of dust. The surface temperatures have been retrieved from the LWC spectra for each observation. We found an increase in the surface temperature in this region, indicating a spatial mixture of cold CO2 ice and warmer dust/H2O ice. Region IV is close to the top of the residual cap [81–84°N]; it is much brighter than region III, with a dust content 10 times lower than the latter. The CO2 grain size is 3 mm and strong CO2 ice features are present in the data, indicating a thicker CO2 ice layer than in region II (1–2 g cm−2). The final region V is right at the top of the residual cap (⩾84°N). It is “pure” CO2 ice (no dust) of 5 mm grain sizes, with 30 ppm by weight of water ice. The CO2 ice features are very pronounced and the 2.7 μm band is saturated. The optical thickness is close to the semi-infinite limit (30–40 g cm−2). Assuming a snowpack density of 0.5 g cm−3, we get a minimum thickness of 1–2 cm for the top-layer of regions II and III, 4–10 cm for region IV, and ⩾60–80 cm thickness for region V. These values are in close agreement with several recent results for the south seasonal polar cap.These results should provide new, useful constraints in models of the Martian climate system and volatile cycles.  相似文献   

5.
The nature of cometary volatile materials is subject to debate. Theoretical models of cometary nuclei and laboratory studies suggest that these objects could be made of amorphous water ice in addition to other volatile molecules and refractory grains. This water ice structure has the ability to encapsulate the gases of surrounding environment, reflecting the physical and chemical conditions during their deposition. Therefore, the knowledge of the chemical composition of volatile molecules trapped in amorphous water ice provides a tool for probing the formation environment of cometary ice grains. Experimental studies of gas trapping efficiency in amorphous water ice have been previously conducted mostly under kinetic conditions, where dynamic pumping and temperature gradients prevented rigorous calibrations. In this work, we investigated the trapping efficiencies of Ar, CO, CH4, Kr and N2 by depositing water vapor as ice in the presence of trace gases in a volume submerged in liquid nitrogen at 77 K. The gas trapping efficiencies were determined simply by monitoring the pressure difference of the trace gases before and after the deposition of a known amount of water molecules as amorphous ice.Our results show that the trapped gas to water molecule ratio in amorphous ice is controlled primarily by the partial pressure of the gas during water ice deposition, and is independent of the ice deposition rate as well as the gas to water ratio in the vapor phase. The trapping efficiencies of gases decrease in the order of Kr > CH4 > CO > Ar > N2 in accordance with previous studies. Assuming that the water ice structure of comets is at least partially amorphous water ice at the time of their formation, these results suggest that the total pressure and composition of the surrounding environment of amorphous ice formation are significant controlling factors of trace gas concentrations in cometary ice. This further indicates that the evolution of the solar nebula and timing of cometary ice condensation can also be important parameters in linking the volatile contents of comets and their formation process.  相似文献   

6.
Spectral observations have detected methane within the martian atmosphere (Formisano, V., Atreya, S., Encrenaz, T., Ignatiev, N., Giuranna, M. [2004]. Science 306, 1758–1761; Mumma, M.J. et al. [2009]. Science 323, 1041–1045), however, the origin of the methane has not been determined. Methane clathrate (also referred to as methane hydrate) has been suggested as a potential subsurface reservoir, storing and releasing biologic and/or abiogenic methane. In this study, rates of methane hydrate formation and dissociation were measured experimentally at 234–264 K and 1.4–4.7 MPa to test the clathrate reservoir hypothesis. Initial formation rates range from 4.3 × 10?6 to 8.1 × 10?5 mol m?2 s?1. Results show decreasing rates of formation over time in individual experiments, indicating initial rapid clathration, followed by diffusion-limited transport of methane into the ice through the previously formed hydrate. These experiments indicate increased pressure results in increased formation rates, likely the result of higher concentration gradients, enhancing the methane diffusion flux into the solid phase. Experiments conducted at elevated temperatures produced faster initial rates of formation, resulting from increased kinetic energy of methane molecules and/or thickening of the Quasi-Liquid Layer. Based on this temperature dependence, the activation energy for methane hydrate formation from ice was determined to be 35.9 kJ/mol. Hydrate dissociation experiments initiated by depressurization or warming at conditions between 222 K and 265 K and 0.1–2.0 MPa were conducted following each formation experiment, yielding methane hydrate dissociation rates from 3.01 × 10?6 to 9.92 × 10?5 mol m?2 s?1. While both hydrate dissociation and formation showed decreasing instantaneous rates over the course of each experiment, the transition between the initial rate of dissociation and the interpreted diffusion-limited period of continued dissociation was more abrupt than that observed in formation experiments, supporting an ice shielding effect. The initial concentration of methane in the solid phase had a significant effect on hydrate dissociation rates. Higher methane concentrations in the solid phase produce faster initial rates, likely due to increased concentration gradients, thus increasing the diffusion component of dissociation. Increased temperatures also produced faster dissociation rates, yielding an activation energy for dissociation of 32.7 kJ/mol. The rates determined within this study suggest that small near-surface methane hydrate reservoirs are a feasible source for recent methane plumes detected on Mars. Rates of methane release from gas hydrates also indicate that gas hydrate dissociation may have played a role in forming ancient chaos terrain and associated outflow channels.  相似文献   

7.
An extensive region of low, sinuous ridges occupies the Hesperian plateau above Echus Chasma in the upper Kasei Valles, Mars. The ridges have lengths of up to 270 km, heights of 100 m and widths of 10 km. The total volume of the ridge material is 6×1011 m3. In this paper, volcanic flows, depositional and erosional features are discussed using Mars Observer Laser Altimeter (MOLA), THEMIS and Mars Orbiter Camera (MOC) imagery and a chronology that places the ridge formation in the Late Hesperian is developed.The plateau is bounded to the north and west by more recent Late Hesperian and Amazonian lava flows. The plateau floor suddenly changes from being relatively smooth, to elevated, rough, hummocky terrain that extends eastwards to Echus Chasma. This rough terrain is penetrated by 2 km broad, shallow entrant channels that join with the canyons of Echus Chasma. The sinuous ridges appear to control the surface drainage associated with the entrant channels.The sinuous ridges’ size and morphology are similar to those associated with volcanic ridge eruptions. Their degraded structure is reminiscent of Moberg ridges. The rough, hummocky terrain is interpreted as glacial outwash, subsequently eroded by short-lived floods associated with ridge eruptions. The presence of both volcanic and glacial structures on the Echus Plateau raises the possibility that the ridge system arose from subglacial, volcanic events. The resulting jokulhlaups eroded the broad, entrant channels. As surface flow declined, groundwater flows dominated and canyon heads eroded back along the entrant channels, by sapping.  相似文献   

8.
The eruptive plumes and large heat flow (～15 GW) observed by Cassini in the South Polar Region of Enceladus may be expressions of hydrothermal activity inside Enceladus. We hypothesize that a subsurface ocean is the heat reservoir for thermal anomalies on the surface and the source of heat and chemicals necessary for the plumes. The ocean is believed to contain dissolved gases, mostly CO2 and is found to be relatively warm (～0 °C). Regular tidal forces open cracks in the icy crust above the ocean. Ocean water fills these fissures. There, the conditions are met for the upward movement of water and the dissolved gases to exsolve and form bubbles, lowering the bulk density of the water column and making the pressure at its bottom less than that at the top of the ocean. This pressure difference drives ocean water into and up the conduits toward the surface. This transportation mechanism supports the thermal anomalies and delivers heat and chemicals to the chambers from which the plumes erupt. Water enters these chambers and there its bubbles pop and loft an aerosol mist into the ullage. The exiting plume gas entrains some of these small droplets. Thus, nonvolatile chemical species in ocean water can be present in the plume particles. A CO2 equivalent-gas molar fraction of ～4 × 10?4 for the ocean is sufficient to support the circulation. A source of heat is needed to keep the ocean warm at ～0 °C (about two degrees above its freezing point). The source of heat is unknown, but our hypothesis is not dependent on any particular mechanism for producing the heat.  相似文献   

9.
Laboratory experiments show that dusty bodies in a gaseous environment eject dust particles if they are illuminated. We find that even more intense dust eruptions occur when the light source is turned off. We attribute this to a compression of gas by thermal creep in response to the changing temperature gradients in the top dust layers. The effect is studied at a light flux of 13 kW/m2 and 1 mbar ambient pressure. The effect is applicable to protoplanetary disks and Mars. In the inner part of protoplanetary disks, planetesimals can be eroded especially at the terminator of a rotating body. This leads to the production of dust which can then be transported towards the disk edge or the outer disk regions. The generated dust might constitute a significant fraction of the warm dust observed in extrasolar protoplanetary disks. We estimate erosion rates of about 1 kg s?1 for 100 m parent bodies. The dust might also contribute to subsequent planetary growth in different locations or on existing protoplanets which are large enough not to be susceptible to particle loss by light induced ejection. Due to the ejections, planetesimals and smaller bodies will be accelerated or decelerated and drift outward or inward, respectively. The effect might also explain the entrainment of dust in dust devils on Mars, especially at high altitudes where gas drag alone might not be sufficient.  相似文献   

10.
The study is a detailed look on one of the several fluvial systems located on the eastern rim region of the Hellas basin on Mars. We analyzed the morphologic and morphometric characteristics of an extensive channel system, which extends for over 650 km from 35.8°S, 106.4°E in Hesperia Planum to Reull Vallis at 39.5°S, 98.1°E, and has a drainage area of 35,000–40,000 km2. During its traverse the channel changes its characteristics many times, indicating variations in the surface properties. Based on cross-cutting relations, the fluvial system post-dates the emplacement of the early Hesperian lava plains in Hesperia Planum but predates the Amazonian deposits. We describe the geomorphology and evolution of the system and provide evidence of both surface flow and groundwater sapping processes. A chain of channeled paleolake basins in the central parts of the system (38°S, 102°E) provides a rough estimate for the water volume (250–300 km3) which was required to form the system. The minimum volume of surface materials eroded by the channel system is ～74 km3. Although this study presents the detailed analysis of only one fluvial system, the presence of many similar channel systems along the margin of Hellas suggests that late-stage surface runoff has played a significant role in the degradation of the rim of the basin and also in the transportation of materials towards Hellas floor.  相似文献   

11.
CRISM indicates the presence of water ice patches in Richardson crater, located on Mars’ southern polar region at the area of the seasonal ice cap. Numerical simulations suggest that the maximum daytime temperature of the ice at these locations is between 195 and 220 K during local spring. Previous studies suggest that at these temperatures liquid interfacial water could be present. Here, for the first time, we provide an example where the environmental conditions allow for the formation of such liquid films on present day Mars at the southern hemisphere. The upper bound estimated H2O loss during the presence of these water ice patches is approximately 30 μm between Ls = 200 and 220, though it may be as low as 0.1 μm depending on the ambient water vapor. The upper bound value is larger than the expected condensation thickness in autumn; however, it may still be realistic due to CO2 gas jet generated deposition and possible subsequent accumulation on mineral grains. The presence of this interfacial water may have impact on local chemical processes along with astrobiological importance.  相似文献   

12.
In this work we investigated changes of the water emission from a model comet of the size and orbital elements of Comet P/2008 R1 (Garradd). We performed simulations for model cometary nuclei of different compositions and two different orientations in space. Our simulations indicate, that the emission of water decreases from one orbital period to another one, but in some cases slowly. When the rotation axis of the nucleus lies in the orbital plane the seasonal maximum of water production during the first two orbital periods can be as high as about 1026 mol s?1, but decreases by two orders of magnitude during only 50 orbital periods. The highest rate of water production after many orbital periods is expected when the rotation axis is perpendicular to the orbital plane – the seasonal maximum of water production can be about 5 × 1025 mol s?1 during the first two orbital periods after activation of the comet and no more than 0.8 × 1025 mol s?1 500 orbital periods later. The upper estimate for the production of water derived from observations of P/2008 R1 (Garradd) by Jewitt et al. (Jewitt, D., Yang, B., Haghighipour, N. [2009]. Astron. J. 137, 4313–4321) is 5 × 1025.  相似文献   

13.
In the history of Mars exploration its atmosphere and planetary climatology aroused particular interest. In the study of the minor gases abundance in the Martian atmosphere, water vapour became especially important, both because it is the most variable trace gas, and because it is involved in several processes characterizing the planetary atmosphere. The water vapour photolysis regulates the Martian atmosphere photochemistry, and so it is strictly related to carbon monoxide. The CO study is very important for the so-called “atmosphere stability problem”, solved by the theoretical modelling involving photochemical reactions in which the H2O and the CO gases are main characters.The Planetary Fourier Spectrometer (PFS) on board the ESA Mars Express (MEX) mission can probe the Mars atmosphere in the infrared spectral range between 200 and 2000 cm?1 (5–50 μm) with the Long Wavelength Channel (LWC) and between 1700 and 8000 cm?1 (1.2–5.8 μm) with the Short Wavelength Channel (SWC). Although there are several H2O and CO absorption bands in the spectral range covered by PFS, we used the 3845 cm?1 (2.6 μm) and the 4235 cm?1 (2.36 μm) bands for the analysis of water vapour and carbon monoxide, respectively, because these ranges are less affected by instrumental problems than the other ones. The gaseous concentrations are retrieved by using an algorithm developed for this purpose.The PFS/SW dataset used in this work covers more than two and a half Martian years from Ls=62° of MY 27 (orbit 634) to Ls=203° of MY 29 (orbit 6537). We measured a mean column density of water vapour of about 9.6 pr. μm and a mean mixing ratio of carbon monoxide of about 990 ppm, but with strong seasonal variations at high latitudes. The seasonal water vapour map reproduces very well the known seasonal water cycle. In the northern summer, water vapour and CO show a good anticorrelation most of the time. This behaviour is due to the carbon dioxide and water sublimation from the north polar ice cap, which dilutes non-condensable species including carbon monoxide. An analogous process takes place during the winter polar cap, but in this case the condensation of carbon dioxide and water vapour causes an increase of the concentration of non-condensable species. Finally, the results show the seasonal variation of the carbon monoxide mixing ratio with the surface pressure.  相似文献   

14.
Vertical distributions and spectral characteristics of Titan’s photochemical aerosol and stratospheric ices are determined between 20 and 560 cm?1 (500–18 μm) from the Cassini Composite Infrared Spectrometer (CIRS). Results are obtained for latitudes of 15°N, 15°S, and 58°S, where accurate temperature profiles can be independently determined.In addition, estimates of aerosol and ice abundances at 62°N relative to those at 15°S are derived. Aerosol abundances are comparable at the two latitudes, but stratospheric ices are ～3 times more abundant at 62°N than at 15°S. Generally, nitrile ice clouds (probably HCN and HC3N), as inferred from a composite emission feature at ～160 cm?1, appear to be located over a narrow altitude range in the stratosphere centered at ～90 km. Although most abundant at high northern latitudes, these nitrile ice clouds extend down through low latitudes and into mid southern latitudes, at least as far as 58°S.There is some evidence of a second ice cloud layer at ～60 km altitude at 58°S associated with an emission feature at ～80 cm?1. We speculate that the identify of this cloud may be due to C2H6 ice, which in the vapor phase is the most abundant hydrocarbon (next to CH4) in the stratosphere of Titan.Unlike the highly restricted range of altitudes (50–100 km) associated with organic condensate clouds, Titan’s photochemical aerosol appears to be well-mixed from the surface to the top of the stratosphere near an altitude of 300 km, and the spectral shape does not appear to change between 15°N and 58°S latitude. The ratio of aerosol-to-gas scale heights range from 1.3–2.4 at about 160 km to 1.1–1.4 at 300 km, although there is considerable variability with latitude. The aerosol exhibits a very broad emission feature peaking at ～140 cm?1. Due to its extreme breadth and low wavenumber, we speculate that this feature may be caused by low-energy vibrations of two-dimensional lattice structures of large molecules. Examples of such molecules include polycyclic aromatic hydrocarbons (PAHs) and nitrogenated aromatics.Finally, volume extinction coefficients NχE derived from 15°S CIRS data at a wavelength of λ = 62.5 μm are compared with those derived from the 10°S Huygens Descent Imager/Spectral Radiometer (DISR) data at 1.583 μm. This comparison yields volume extinction coefficient ratios NχE(1.583 μm)/NχE(62.5 μm) of roughly 70 and 20, respectively, for Titan’s aerosol and stratospheric ices. The inferred particle cross-section ratios χE(1.583 μm)/χE(62.5 μm) appear to be consistent with sub-micron size aerosol particles, and effective radii of only a few microns for stratospheric ice cloud particles.  相似文献   

15.
We have examined thermal emission from 240 active or recently-active volcanic features on Io and quantified the magnitude and distribution of their volcanic heat flow during the Galileo epoch. We use spacecraft data and a geological map of Io to derive an estimate of the maximum possible contribution from small dark areas not detected as thermally active but which nevertheless appear to be sites of recent volcanic activity. We utilize a trend analysis to extrapolate from the smallest detectable volcanic heat sources to these smallest mapped dark areas. Including the additional heat from estimates for “outburst” eruptions and for a multitude of very small (“myriad”) hot spots, we account for ～62 × 1012 W (～59 ± 7% of Io’s total thermal emission). Loki Patera contributes, on average, 9.6 × 1012 W (～9.1 ± 1%). All dark paterae contribute 45.3 × 1012 W (～43 ± 5%). Although dark flow fields cover a much larger area than dark paterae, they contribute only 5.6 × 1012 W (～5.3 ± 0.6%). Bright paterae contribute ～2.6 × 1012 W (～2.5 ± 0.3%). Outburst eruption phases and very small hot spots contribute no more than ～4% of Io’s total thermal emission: this is probably a maximum value. About 50% of Io’s volcanic heat flow emanates from only 1.2% of Io’s surface. Of Io’s heat flow, 41 ± 7.0% remains unaccounted for in terms of identified sources. Globally, volcanic heat flow is not uniformly distributed. Power output per unit surface area is slightly biased towards mid-latitudes, although there is a stronger bias toward the northern hemisphere when Loki Patera is included. There is a slight favoring of the northern hemisphere for outbursts where locations were well constrained. Globally, we find peaks in thermal emission at ～315°W and ～105°W (using 30° bins). There is a minimum in thermal emission at around 200°W (almost at the anti-jovian longitude) which is a significant regional difference. These peaks and troughs suggest a shift to the east from predicted global heat flow patterns resulting from tidal heating in an asthenosphere. Global volcanic heat flow is dominated by thermal emission from paterae, especially from Loki Patera (312°W, 12°N). Thermal emission from dark flows maximises between 165°W and 225°W. Finally, it is possible that a multitude of very small hot spots, smaller than the present angular resolution detection limits, and/or cooler, secondary volcanic processes involving sulphurous compounds, may be responsible for at least part of the heat flow that is not associated with known sources. Such activity should be sought out during the next mission to Io.  相似文献   

16.
The exosphere of an atmosphereless icy moon is the result of different surface release processes and subsequent modification of the released particles. At Europa icy moon, water molecules are directly released, but photolysis and radiolysis due to solar UV and Jupiter’s magnetospheric plasma, respectively, can result in OH, H, O and (possibly) H2 production. These molecules can recombine to reform water and/or new chemical species. As a consequence, Europa’s neutral environment becomes a mixture of different molecules, among which, H2O dominates in the highest altitudes and O2, formed mainly by radiolysis of ice and subsequent release of the produced molecules, prevails at lower altitudes. In this work, starting from a previously developed Monte Carlo model for the generation of Europa’s exosphere, where the only considered species was water, we make a first attempt to simulate also the H2 and O2 components of the neutral environment around Europa, already observed by the Hubble Space Telescope and the Ultraviolet Imaging Spectrograph on board Cassini, during its flyby of Jupiter. Considering a specific configuration where the leading hemisphere coincides with the sunlit hemisphere, we estimate along the Europa–Sun line an O2 column density of about 1.5 × 1019 m?2 at the dayside and 3 × 1018 m?2 at the nightside. In this work we also improve our previous estimation of the sputtered H2O exosphere of this moon, taking into consideration the trailing–leading asymmetry in the magnetospheric ion bombardment and the energy and temperature dependences of the process yields. We find that a density of 1.5 × 1012 H2O/m3 is expected at altitudes ～0.1RE above the surface of the trailing hemisphere. Additionally, we calculate the escape of H2O, O2 and H2. The total number of neutral atoms in Europa’s neutral torus, is estimated to be in the range 7.8 × 1032–3.3 × 1033.  相似文献   

17.
Global recharge of the martian hydrologic system has traditionally been viewed as occurring through basal melting of the south polar cap. We conclude that regional recharge of a groundwater system at the large volcanic provinces, Elysium and Tharsis, is also very plausible and has several advantages over a south polar recharge source in providing a more direct, efficient supply of water to the outflow channel source regions surrounding these areas. This recharge scenario is proposed to have operated concurrently with and within the context of a global cryosphere–hydrosphere system of the subsurface characteristic of post-Noachian periods. To complement existing groundwater flow modeling studies, we examine geologic evidence and possible mechanisms for accumulation of water at high elevations on the volcanic rises, such as melting snow, infiltration, and increased effective permeability of the subsurface between the recharge zone and outflow source. Evidence for the presence of large Amazonian-aged cold-based piedmont glaciers on the Tharsis Montes has been well documented. Climate modeling predicts snow accumulation on high volcanic rises at obliquities thought to be typical over much of martian history. Thermal gradients causing basal melting of snowpack over 1 km thick could provide several kg m−2 yr−1 of water, charging a volume equivalent to the pore space in a square meter column of subsurface in less than 1.5×105 yr. In order to account for estimated outflow channel volumes, the subsurface volume above the elevation of the outflow channels must be charged several times over the area of Tharsis. Complete aquifer recharge can be accomplished in ∼0.3–2 My through the snowpack melting mechanism at Tharsis and in ∼5×104 years for channel requirements at Elysium. Abundant radial dikes emanating from large martian volcanic rises can crack and/or melt the cryosphere, initiating water outflow and creating anisotropies that can channel subsurface water from a high-elevation groundwater reservoir to outflow sources. In this model, snow accumulation, infiltration of meltwater, and increased effective permeabilities are a consequence of the geologic, thermal, and climatic environment at Elysium and Tharsis, and may have had a genetic influence on the preferential distribution of outflow channels around volcanic rises on Mars.  相似文献   

18.
Currently, and throughout much of the Amazonian, the mean annual surface temperatures of Mars are so cold that basal melting does not occur in ice sheets and glaciers and they are cold-based. The documented evidence for extensive and well-developed eskers (sediment-filled former sub-glacial meltwater channels) in the south circumpolar Dorsa Argentea Formation is an indication that basal melting and wet-based glaciation occurred at the South Pole near the Noachian–Hesperian boundary. We employ glacial accumulation and ice-flow models to distinguish between basal melting from bottom-up heat sources (elevated geothermal fluxes) and top-down induced basal melting (elevated atmospheric temperatures warming the ice). We show that under mean annual south polar atmospheric temperatures (?100 °C) simulated in typical Amazonian climate experiments and typical Noachian–Hesperian geothermal heat fluxes (45–65 mW/m2), south polar ice accumulations remain cold-based. In order to produce significant basal melting with these typical geothermal heat fluxes, the mean annual south polar atmospheric temperatures must be raised from today’s temperature at the surface (?100 °C) to the range of ?50 to ?75 °C. This mean annual polar surface atmospheric temperature range implies lower latitude mean annual temperatures that are likely to be below the melting point of water, and thus does not favor a “warm and wet” early Mars. Seasonal temperatures at lower latitudes, however, could range above the melting point of water, perhaps explaining the concurrent development of valley networks and open basin lakes in these areas. This treatment provides an independent estimate of the polar (and non-polar) surface temperatures near the Noachian–Hesperian boundary of Mars history and implies a cold and relatively dry Mars climate, similar to the Antarctic Dry Valleys, where seasonal melting forms transient streams and permanent ice-covered lakes in an otherwise hyperarid, hypothermal climate.  相似文献   

19.
《Planetary and Space Science》2007,55(12):1741-1756
The dynamics of Venus’ mesosphere (70–110 km) is characterized by the superposition of two different wind regimes: (1) Venus’ retrograde superrotation; (2) a sub-solar to anti-solar (SS–AS) flow pattern, driven by solar EUV heating on the sunlit hemisphere. Here, we report on new ground-based velocity measurements in the lower part of the mesosphere. We took advantage of two essentially symmetric Venus elongations in 2001 and 2002 to perform high-resolution Doppler spectroscopy (R=120,000) in 12C16O2 visible lines of the 5ν3 band and in a few solar Fraunhofer lines near 8700 Å. These measurements, mapped over several points on Venus’ illuminated hemisphere, probe the region of cloud tops. More precisely, the solar Fraunhofer lines sample levels a few kilometers below the UV features (i.e. near ∼67 km), while the CO2 lines probe an altitude higher by about 7 km. The wind field over Venus’ disk is retrieved with an rms uncertainty of 15–25 m s−1 on individual measurements. Kinematical fit to a one- or two-component circulation model indicates the dominance of the zonal retrograde flow with a mean equatorial velocity of ∼75 m s−1, exhibiting very strong day-to-day variations (±65 m s−1). Results are very consistent for the two kinds of lines, suggesting a negligible vertical wind shear over 67–74 km. The SS–AS flow is not detected in single-day observations, but combining the results from all data suggests that this component may invade the lower mesosphere with a ∼40 m s−1 velocity.  相似文献   

20.
We suggest that planets, brown dwarfs, and even low mass stars can be formed by fragmentation of protoplanetary disks around very massive stars (M ? 100 M⊙). We discuss how fragmentation conditions make the formation of very massive planetary systems around very massive stars favorable. Such planetary systems are likely to be composed of brown dwarfs and low mass stars of ～0.1–0.3 M⊙, at orbital separations of ～ few × 100–104 AU. In particular, scaling from solar-like stars suggests that hundreds of Mercury-like planets might orbit very massive stars at ～103 AU where conditions might favor liquid water. Such fragmentation objects can be excellent targets for the James Webb Space Telescope and other large telescopes working in the IR bands. We predict that deep observations of very massive stars would reveal these fragmentation objects, orbiting in the same orbital plane in cases where there are more than one object.  相似文献   

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