Titan’s aerosol and stratospheric ice opacities between 18 and 500 μm: Vertical and spectral characteristics from Cassini CIRS

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 Nχ_E 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 Nχ_E(1.583 μm)/Nχ_E(62.5 μm) of roughly 70 and 20, respectively, for Titan’s aerosol and stratospheric ices. The inferred particle cross-section ratios χ_E(1.583 μm)/χ_E(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.     

Spectral imaging of martian water ice clouds and their diurnal behavior during the 1999 aphelion season (Ls = 130°)

We report high-spectral-resolution (λ/δλ = 800-2300) near-infrared mapping observations of Mars at L_s = 130° (April 1999), which were obtained by drift-scanning the cryogenic long-slit spectrometer at the KPNO 2.2-m telescope across the disk. Data were reformatted into calibrated spectral image cubes (x,y,λ) spanning 2.19 to 4.12 μm, which distinguish atmospheric CO₂ features, solar lines, and surface and aerosol features. Maps of relative band depth between 3.0 and 3.5 μm trace water ice clouds and show the diurnal evolution of features in the persistent northern summer aphelion cloud belt, which was mapped contemporaneously but at fixed local time by the Mars Global Surveyor Thermal Emission Spectrometer (MGS/TES). Cloud optical depth, particle sizes, and ice aerosol content were estimated using a two-stream, single-layer scattering model, with Mie coefficients derived from recently published ice optical constants, followed by a linear spectral deconvolution process. A comparison of data and model spectra shows evaporating nighttime clouds in the morning followed by afternoon growth of a prominent orographic cloud feature on the west flank of Elysium Mons. Cloud optical depth at 3.2 μm evolved to 0.28 ± 0.13 and ice aerosol column abundance to 0.9 ± 0.3 pr μm in the afternoon. Column abundances as large as 0.17 pr μm were retrieved in nonorographic clouds within the aphelion cloud band around midday. These clouds exhibit a modest decline in optical depth during the afternoon. Results show that ice particle radii from <2 μm to >4 μm exist in both cloud types. However, large particles dominate the spectra, consistent with recent MGS/TES emission phase function measurements of aphelion cloud aerosol properties.     

Spatial variability,composition and thickness of the seasonal north polar cap of Mars in mid-spring

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.     

Subvisible CO2 ice clouds detected in the mesosphere of Mars

The formation of CO₂ ice clouds in the upper atmosphere of Mars has been suggested in the past on the basis of a few temperature profiles exhibiting portions colder than CO₂ frost point. However, the corresponding clouds were never observed. In this paper, we discuss the detection of the highest clouds ever observed on Mars by the SPICAM ultraviolet spectrometer on board Mars Express spacecraft. Analyzing stellar occultations, we detected several mesospheric detached layers at about 100 km in the southern winter subtropical latitudes, and found that clouds formed where simultaneous temperature measurements indicated that CO₂ was highly supersaturated and probably condensing. Further analysis of the spectra reveals a cloud opacity in the subvisible range and ice crystals smaller than 100 nm in radius. These layers are therefore similar in nature as the noctilucent clouds which appear on Earth in the polar mesosphere. We interpret these phenomena as CO₂ ice clouds forming inside supersaturated pockets of air created by upward propagating thermal waves. This detection of clouds in such an ultrararefied and supercold atmosphere raises important questions about the martian middle-atmosphere dynamics and microphysics. In particular, the presence of condensates at such high altitudes begs the question of the origin of the condensation nuclei.     

Search for methane and upper limits to ethane and SO2 on Mars

The IRTF/CSHELL observations in February 2006 at L_S = 10° and 63–93°W show ～10 ppb of methane at 45°S to 7°N and ～3 ppb outside this region that covers the deepest canyon Valles Marineris. Observations in December 2009 at L_S = 20° and 0–30°W included spectra of the Moon at a similar airmass as a telluric calibrator. A technique for extraction of the martian methane line from a combination of the Mars and Moon spectra has been developed. The observations reveal no methane with an upper limit of 8 ppb. The results of both sessions agree with the observations by Mumma et al. (Mumma, M.J. et al. [2009]. Science 323, 1041–1045) at the same season in February 2006 and are smaller than those in the PFS and TES maps. Production and removal of the biological methane on Mars do not significantly change the redox state of the atmosphere and the balance of hydrogen. A search for ethane at 2977 cm^?1 results in an upper limit of 0.2 ppb. However, this limit does not help to establish the origin of methane on Mars. Reanalysis of our search for SO₂ using TEXES confirms the recently established upper limit of 0.3 ppb. Along with the lack of hot spots and gas vents with endogenic heat sources in the THEMIS observations, the very low upper limit to SO₂ on Mars does not favor geological methane that is less abundant than SO₂ in the outgassing from the terrestrial planets.     

Mars mesospheric zonal wind around northern spring equinox from infrared heterodyne observations of CO2

We present direct observations of Mars zonal wind velocities around northern spring equinox (L_S = 336°, L_S = 355°, L_S = 42°) during martian year 27 and 29. Data was acquired by means of infrared heterodyne spectroscopy of CO₂ features at 959.3917 cm^?1 (10.4232 μm) and 957.8005 cm^?1 (10.4405 μm) using the Cologne Tuneable Heterodyne Infrared Spectrometer (THIS) at the McMath–Pierce telescope of the National Solar Observatory on Kitt Peak in Arizona and the NASA Infrared Telescope Facility on Mauna Kea, Hawaii between 2005 and 2008. Winds were measured on the dayside of Mars with an unprecedented spatial resolution allowing sampling of up to nine independent latitudes over the martian disk. Retrieved wind velocities depend strongly on latitude and season with values ranging from 180 m/s prograde to ?94 m/s retrograde. A comparison of the observational results to predicted values from the Mars Climate Database yield a reasonable agreement between modeling and observation.     

Locations of thin liquid water layers on present-day Mars

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.     

Interannual variability in TES atmospheric observations of Mars during 1999-2003

We use infrared spectra returned by the Mars Global Surveyor Thermal Emission Spectrometer (TES) to retrieve atmospheric and surface temperature, dust and water ice aerosol optical depth, and water vapor column abundance. The data presented here span more than two martian years (Mars Year 24, L_s=104°, 1 March 1999 to Mars Year 26, L_s=180°, 4 May 2003). We present an overview of the seasonal (L_s), latitudinal, and longitudinal dependence of atmospheric quantities during this period, as well as an initial assessment of the interannual variability in the current martian climate. We find that the perihelion season (L_s=180°-360°) is relatively warm, dusty, free of water ice clouds, and shows a relatively high degree of interannual variability in dust optical depth and atmospheric temperature. On the other hand, the aphelion season (L_s=0°-180°) is relatively cool, cloudy, free of dust, and shows a low degree of interannual variability. Water vapor abundance shows a moderate amount of interannual variability at all seasons, but the most in the perihelion season. Much of the small amount of interannual variability that is observed in the aphelion season appears to be caused by perihelion-season planet-encircling dust storms. These dust storms increase albedo through deposition of bright dust on the surface causing cooler daytime surface and atmospheric temperatures well after dust optical depth returns to prestorm values.     

Tracking the edge of the south seasonal polar cap of Mars

The advance and retreat of the polar caps were one of the first observations that indicated Mars had seasons. Because a large portion of the atmosphere is cycled in and out of the seasonal caps during the year, the frost deposits play a significant role in regional and global atmospheric circulation. Understanding the nature of the seasonal polar caps is imperative if we are to understand the current Martian climate. In this study, we track the seasonal cap edges as a function of season and longitude for the fall and winter seasons (MY27), using data from the Planetary Fourier Spectrometer (PFS) onboard the Mars Express (MEX) ESA mission. Making use of the rapid rise (decrease) in surface temperature that occurs when CO₂ ice is removed (deposited), in a first approach, we defined the advancing cap edge to be where the surface temperature drops below 150 K, and the retreating cap edge where the surface temperature rises above 160 K. In this case, starting from Ls∼50°, the edge progression speed start to be longitude dependent. In the hemisphere that extends form the eastern limit of the Hellas basin to the western limit of the Argyrae basin (and containing the two) the edges progression speed is about a half than that of the other hemisphere; the cap is thus asymmetric and, unexpectedly, no CO₂ ice seems to be present inside the basins. This is because the above mentioned surface temperatures used in this approach to detect the cap edges are not adequate (too low) for the high-pressure regions inside the basins where, following the Clausius–Clapeyron's law, the CO₂ condensation temperature can be several degrees higher than that of the adjacent lower-pressure regions. In the second, final approach, special attention has been given to this aspect and the advancing and retreating cap edges are defined where, respectively, the surface temperatures drop below and rise above the CO₂ condensation temperature for the actual surface pressure values. Now, the results show an opposite situation than the previous one, with the progression speed being higher and the cap more extended (up to −30° latitude) in the hemisphere containing the two major Martian basins. During the fall season, up to Ls∼50° the South Martian polar cap consists of CO₂ frost deposits that advance towards lower latitudes at a constant speed of 10° of latitude per 15 degrees of Ls. The maximum extension (−40° latitude) of the South polar cap occurs somewhere in the 80°–90° Ls range. At the winter solstice, when the edges of the polar night start moving poleward, the cap recession has already started, in response to seasonal changes in insolation. The CO₂ ice South polar cap will recede with a constant speed of ∼5° of latitude every 25° degrees of Ls during the whole winter. The longitudinal asymmetries reduce during the cap retreat and completely disappear around Ls=145°.     

Io’s atmosphere: Constraints on sublimation support from density variations on seasonal timescales using NASA IRTF/TEXES observations from 2001 to 2010

We present an analysis of 19 μm spectra of Io’s SO₂ atmosphere from the TEXES mid-infrared high spectral resolution spectrograph on NASA’s Infrared Telescope Facility, incorporating new data taken between January 2005 and June 2010 and a re-analysis of earlier data taken from November 2001 to January 2004. This is the longest set of contiguous observations of Io’s atmosphere using the same instrument and technique thus far. We have fitted all 16 detected blended absorption lines of the ν₂ SO₂ vibrational band to retrieve the subsolar values of SO₂ column abundance and the gas kinetic temperature. By incorporating an existing model of Io’s surface temperatures and atmosphere, we retrieve sub-solar column densities from the disk-integrated data. Spectra from all years are best fit by atmospheric temperatures <150 K. Best-fit gas kinetic temperatures on the anti-Jupiter hemisphere, where SO₂ gas abundance is highest, are low and stable, with a mean of 108 (±18) K. The sub-solar SO₂ column density between longitudes of 90–220° varies from a low of 0.61 (±0.145) × 10^?17 cm^?2, near aphelion in 2004, to a high of 1.51 (±0.215) × 10¹⁷ cm^?2 in 2010 when Jupiter was approaching its early 2011 perihelion. No correlation in the gas temperature was seen with the increasing SO₂ column densities outside the errors.Assuming that any volcanic component of the atmosphere is constant with time, the correlation of increasing SO₂ abundance with decreasing heliocentric distance provides good evidence that the atmosphere is at least partially supported by frost sublimation. The SO₂ frost thermal inertias and albedos that fit the variation in atmospheric density best are between 150–1250 W m^?2 s^?1/2 K^?1 and 0.613–0.425 respectively. Photometric evidence favors albedos near the upper end of this range, corresponding to thermal inertias near the lower end. This relatively low frost thermal inertia produces larger amplitude seasonal variations than are observed, which in turn implies a substantial additional volcanic atmospheric component to moderate the amplitude of the seasonal variations of the total atmosphere on the anti-Jupiter hemisphere. The seasonal thermal inertia we measure is unique both because it refers exclusively to the SO₂ frost surface component, and also because it refers to relatively deep subsurface layers (few meters) due to the timescales of many years, while previous studies have determined thermal inertias at shallower levels (few centimeters), relevant for timescales of ～2 h (eclipse) or ～2 days (diurnal curves).     

PFS-MEX observation of ices in the residual south polar cap of Mars

The Mars Express spacecraft has a highly inclined orbit around Mars and so has been able to observe the south pole of Mars in illuminated conditions at the end of the southern summer (L_s=330). Spectra from the planetary Fourier spectrometer (PFS) short wavelength (SW) channel were recorded over the permanent ice cap to study its composition in terms of CO₂ ice and H₂O ice. Models are fitted to the observed data, which include a spatial mixture of soil (not covered by ice) and CO₂ frost (with a specific grain size and a small amount of included dust and H₂O ice). Two different kinds of spectra were observed: those over the permanent polar cap with almost pure CO₂ ice, negligible water ice, no soil fraction required, and bright; and those over mixed terrain (at the edge of the cap or near troughs) containing a significant soil spatial fraction, more water ice and smaller CO₂ grain size. The amount of water ice given by fits to scaled albedo models is less than 10 ppm by weight. When using multi-stream reflectance models with the appropriate lighting geometry, the water amount must be 2-5 times greater than the albedo fit (less than 50 ppm). At the periphery of the residual polar cap, we found a region almost completely covered by water frost, modeled as a mixture of micron-sized and sub-mm sized grains. Our result using a granular mixture of micron-sized grains of water ice and dust with the CO₂ grains is different from the modeling of OMEGA polar cap observations using molecular mixtures.     

Water ice clouds in the Martian atmosphere: Two Martian years of SPICAM nadir UV measurements

The SPICAM instrument onboard Mars Express has successfully performed two Martian years (MY 27 and MY28) of observations. Water ice cloud optical depths spatial and temporal distribution was retrieved from nadir measurements in the wavelength range 300–320 nm. During the northern spring the cloud hazes complex distribution was monitored. The clouds in the southern hemisphere formed a zonal belt in the latitude range 30–60°S. The edge of the retreating north polar hood merged with the northern tropical clouds in the range 250–350°E. The development of the aphelion cloud belt (ACB) started with the weak hazes formation (cloud optical thickness 0.1–0.3) in the equatorial region. At the end of the northern spring, the ACB cloud optical thickness reached already values of 0.3–1. The ACB decay in the end of the northern summer was accompanied with a presence of clouds in the north mid-latitudes. The expanded north polar hood merged with the north mid-latitude clouds in the eastern hemisphere. The interannual comparison indicates a decrease in cloud activity immediately after a strong dust storm in southern summer of MY28. The strong dust storms of the MY28 may also be a reason of the observed north polar hood edge shifting northward by 5°.     

The O2 nightglow in the martian atmosphere by SPICAM onboard of Mars-Express

We present observations of the O₂(a¹Δ_g) nightglow at 1.27 μm on Mars using the SPICAM IR spectrometer onboard of the Mars Express orbiter. In contrast to the O₂(a¹Δ_g) dayglow that results from the ozone photodissociation, the O₂(a¹Δ_g) nightglow is a product of the recombination of O atoms formed by CO₂ photolysis on the dayside at altitudes higher than 80 km and transported downward above the winter pole by the Hadley circulation. The first detections of the O₂(a¹Δ_g) nightglow in 2010 indicate that it is about two order of magnitude less intense than the dayglow (Bertaux, J.-L., Gondet, B., Bibring, J.-P., Montmessin, F., Lefèvre, F. [2010]. Bull. Am. Astron. Soc. 42, 1040; Clancy et al. [2010]. Bull. Am. Astron. Soc. 42, 1041). SPICAM IR sounds the martian atmosphere in the near-IR range (1–1.7 μm) with the spectral resolution of 3.5 cm^?1 in nadir, limb and solar occultation modes. In 2010 the vertical profiles of the O₂(a¹Δ_g) nightside emission have been obtained near the South Pole at latitudes of 82–83°S for two sequences of observations: L_s = 111–120° and L_s = 152–165°. The altitude of the emission maximum varied from 45 km on L_s = 111–120° to 38–49 km on L_s = 152–165°. Averaged vertically integrated intensity of the emission at these latitudes has shown an increase from 0.22 to 0.35 MR. Those values of total vertical emission rate are consistent with the OMEGA observations on Mars-Express in 2010. The estimated density of oxygen atoms at altitudes from 50 to 65 km varies from 1.5 × 10¹¹ to 2.5 × 10¹¹ cm^?3. Comparison with the LMD general circulation model with photochemistry (Lefèvre, F., Lebonnois, S., Montmessin, F., Forget, F. [2004]. J. Geophys. Res. 109, E07004; Lefèvre et al. [2008]. Nature 454, 971–975) shows that the model reproduces fairly well the O₂(a¹Δ_g) emission layer observed by SPICAM when the large field of view (>20 km on the limb) of the instrument is taken into account.     

THEMIS-VIS observations of clouds in the martian mesosphere: Altitudes, wind speeds, and decameter-scale morphology

We present measurements of the altitude and eastward velocity component of mesospheric clouds in 35 imaging sequences acquired by the Mars Odyssey (ODY) spacecraft’s Thermal Emission Imaging System visible imaging subsystem (THEMIS-VIS). We measure altitude by using the parallax drift of high-altitude features, and the velocity by exploiting the time delay in the THEMIS-VIS imaging sequence.We observe two distinct classes of mesospheric clouds: equatorial mesospheric clouds observed between 0° and 180° L_s; and northern mid-latitude clouds observed only in twilight in the 200–300° L_s period. The equatorial mesospheric clouds are quite rare in the THEMIS-VIS data set. We have detected them in only five imaging sequences, out of a total of 2048 multi-band equatorial imaging sequences. All five fall between 20° south and 0° latitude, and between 260° and 295° east longitude. The mid-latitude mesospheric clouds are apparently much more common; for these we find 30 examples out of 210 northern winter mid-latitude twilight imaging sequences. The observed mid-latitude clouds are found, with only one exception, in the Acidalia region, but this is quite likely an artifact of the pattern of THEMIS-VIS image targeting. Comparing our THEMIS-VIS images with daily global maps generated from Mars Orbiter Camera Wide Angle (MOC-WA) images, we find some evidence that some mid-latitude mesospheric cloud features correspond to cloud features commonly observed by MOC-WA. Comparing the velocity of our mesospheric clouds with a GCM, we find good agreement for the northern mid-latitude class, but also find that the GCM fails to match the strong easterly winds measured for the equatorial clouds.Applying a simple radiative transfer model to some of the equatorial mesospheric clouds, we find good model fits in two different imaging sequences. By using the observed radiance contrast between cloud and cloud-free regions at multiple visible-band wavelengths, these fits simultaneously constrain the optical depths and particles sizes of the clouds. The particle sizes are constrained primarily by the relative contrasts at the available wavelengths, and are found to be quite different in the two imaging sequences: r_eff = 0.1 μm and r_eff = 1.5 μm. The optical depths (constrained by the absolute contrasts) are substantial: 0.22 and 0.5, respectively. These optical depths imply a mass density that greatly exceeds the saturated mass density of water vapor at mesospheric temperatures, and so the aerosol particles are probably composed mainly of CO₂ ice. Our simple radiative transfer model is not applicable to twilight, when the mid-latitude mesospheric clouds were observed, and so we leave the properties of these clouds as a question for further work.     

Scientific goals for the observation of Venus by VIRTIS on ESA/Venus express mission

The Visible and Infrared Thermal Imaging Spectrometer (VIRTIS) on board the ESA/Venus Express mission has technical specifications well suited for many science objectives of Venus exploration. VIRTIS will both comprehensively explore a plethora of atmospheric properties and processes and map optical properties of the surface through its three channels, VIRTIS-M-vis (imaging spectrometer in the 0.3–1 μm range), VIRTIS-M-IR (imaging spectrometer in the 1–5 μm range) and VIRTIS-H (aperture high-resolution spectrometer in the 2–5 μm range). The atmospheric composition below the clouds will be repeatedly measured in the night side infrared windows over a wide range of latitudes and longitudes, thereby providing information on Venus's chemical cycles. In particular, CO, H₂O, OCS and SO₂ can be studied. The cloud structure will be repeatedly mapped from the brightness contrasts in the near-infrared night side windows, providing new insights into Venusian meteorology. The global circulation and local dynamics of Venus will be extensively studied from infrared and visible spectral images. The thermal structure above the clouds will be retrieved in the night side using the 4.3 μm fundamental band of CO₂. The surface of Venus is detectable in the short-wave infrared windows on the night side at 1.01, 1.10 and 1.18 μm, providing constraints on surface properties and the extent of active volcanism. Many more tentative studies are also possible, such as lightning detection, the composition of volcanic emissions, and mesospheric wave propagation.     

Trapped Ar isotopes in meteorite ALH 84001 indicate Mars did not have a thick ancient atmosphere

Water is not currently stable in liquid form on the martian surface due to the present mean atmospheric pressure of ～7 mbar and mean global temperature of ～220 K. However, geomorphic features and hydrated mineral assemblages suggest that Mars’ climate was once warmer and liquid water flowed on the surface. These observations may indicate a substantially more massive atmosphere in the past, but there have been few observational constraints on paleoatmospheric pressures. Here we show how the ⁴⁰Ar/³⁶Ar ratios of trapped gases within martian meteorite ALH 84001 constrain paleoatmospheric pressure on Mars during the Noachian era [～4.56–3.8 billion years (Ga)]. Our model indicates that atmospheric pressures did not exceed ～1.5 bar during the first 400 million years (Ma) of the Noachian era, and were <400 mbar by 4.16 Ga. Such pressures of CO₂ are only sufficient to stabilize liquid water on Mars’ surface at low latitudes during seasonally warm periods. Other greenhouse gases like SO₂ and water vapor may have played an important role in intermittently stabilizing liquid water at higher latitudes following major volcanic eruptions or impact events.     

THEMIS observations of Mars aerosol optical depth from 2002-2008

We use infrared images obtained by the Thermal Emission Imaging System (THEMIS) instrument on-board Mars Odyssey to retrieve the optical depth of dust and water ice aerosols over more than 3.5 martian years between February 2002 (MY 25, Ls=330°) and December 2008 (MY 29, Ls=183°). These data provide an important bridge between earlier TES observations and recent observations from Mars Express and Mars Reconnaissance Orbiter. An improvement to our earlier retrieval [Smith, M.D., Bandfield, J.L., Christensen, P.R., Richardson, M.I., 2003. J. Geophys. Res. 108, doi:10.1029/2003JE002114] to include atmospheric temperature information from THEMIS Band 10 observations leads to much improved retrievals during the largest dust storms. The new retrievals show moderate dust storm activity during Mars Years 26 and 27, although details of the strength and timing of dust storms is different from year to year. A planet-encircling dust storm event was observed during Mars Year 28 near Southern Hemisphere Summer solstice. A belt of low-latitude water ice clouds was observed during the aphelion season during each year, Mars Years 26 through 29. The optical depth of water ice clouds is somewhat higher in the THEMIS retrievals at ∼5:00 PM local time than in the TES retrievals at ∼2:00 PM, suggestive of possible local time variation of clouds.     

SPICAV on Venus Express: Three spectrometers to study the global structure and composition of the Venus atmosphere

Spectroscopy for the investigation of the characteristics of the atmosphere of Venus (SPICAV) is a suite of three spectrometers in the UV and IR range with a total mass of 13.9 kg flying on the Venus Express (VEX) orbiter, dedicated to the study of the atmosphere of Venus from ground level to the outermost hydrogen corona at more than 40,000 km. It is derived from the SPICAM instrument already flying on board Mars Express (MEX) with great success, with the addition of a new IR high-resolution spectrometer, solar occultation IR (SOIR), working in the solar occultation mode. The instrument consists of three spectrometers and a simple data processing unit providing the interface of these channels with the spacecraft.A UV spectrometer (118–320 nm, resolution 1.5 nm) is identical to the MEX version. It is dedicated to nadir viewing, limb viewing and vertical profiling by stellar and solar occultation. In nadir orientation, SPICAV UV will analyse the albedo spectrum (solar light scattered back from the clouds) to retrieve SO₂, and the distribution of the UV-blue absorber (of still unknown origin) on the dayside with implications for cloud structure and atmospheric dynamics. On the nightside, γ and δ bands of NO will be studied, as well as emissions produced by electron precipitations. In the stellar occultation mode the UV sensor will measure the vertical profiles of CO₂, temperature, SO₂, SO, clouds and aerosols. The density/temperature profiles obtained with SPICAV will constrain and aid in the development of dynamical atmospheric models, from cloud top (∼60 km) to 160 km in the atmosphere. This is essential for future missions that would rely on aerocapture and aerobraking. UV observations of the upper atmosphere will allow studies of the ionosphere through the emissions of CO, CO⁺, and CO₂⁺, and its direct interaction with the solar wind. It will study the H corona, with its two different scale heights, and it will allow a better understanding of escape mechanisms and estimates of their magnitude, crucial for insight into the long-term evolution of the atmosphere.The SPICAV VIS-IR sensor (0.7–1.7 μm, resolution 0.5–1.2 nm) employs a pioneering technology: an acousto-optical tunable filter (AOTF). On the nightside, it will study the thermal emission peeping through the clouds, complementing the observations of both VIRTIS and Planetary Fourier Spectrometer (PFS) on VEX. In solar occultation mode this channel will study the vertical structure of H₂O, CO₂, and aerosols.The SOIR spectrometer is a new solar occultation IR spectrometer in the range λ=2.2–4.3 μm, with a spectral resolution λ/Δλ>15,000, the highest on board VEX. This new concept includes a combination of an echelle grating and an AOTF crystal to sort out one order at a time. The main objective is to measure HDO and H₂O in solar occultation, in order to characterize the escape of D atoms from the upper atmosphere and give more insight about the evolution of water on Venus. It will also study isotopes of CO₂ and minor species, and provides a sensitive search for new species in the upper atmosphere of Venus. It will attempt to measure also the nightside emission, which would allow a sensitive measurement of HDO in the lower atmosphere, to be compared to the ratio in the upper atmosphere, and possibly discover new minor atmospheric constituents.     

Enceladus: A hypothesis for bringing both heat and chemicals to the surface

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.     

Dynamical evolution of the Hungaria asteroids

The Hungarias are a stable asteroid group orbiting between Mars and the main asteroid belt, with high inclinations (16–30°), low eccentricities (e < 0.18), and a narrow range of semi-major axes (1.78–2.06 AU). In order to explore the significance of thermally-induced Yarkovsky drift on the population, we conducted three orbital simulations of a 1000-particle grid in Hungaria a–e–i space. The three simulations included asteroid radii of 0.2, 1.0, and 5.0 km, respectively, with run times of 200 Myr. The results show that mean motion resonances—martian ones in particular—play a significant role in the destabilization of asteroids in the region. We conclude that either the initial Hungaria population was enormous, or, more likely, Hungarias must be replenished through collisional or dynamical means. To test the latter possibility, we conducted three more simulations of the same radii, this time in nearby Mars-crossing space. We find that certain Mars crossers can be trapped in martian resonances, and by a combination of chaotic diffusion and the Yarkovsky effect, can be stabilized by them. Therefore, some Hungarias (around 5% of non-family members with absolute magnitudes H < 15.5 and 10% for H < 17) may represent previously transient Mars crossers that have been adopted in this manner.