Spatially-resolved high-resolution spectroscopy of Venus 1. Variations of CO2, CO, HF, and HCl at the cloud tops

Variations of the upper cloud boundary and the CO, HF, and HCl mixing ratios were observed using the CSHELL spectrograph at NASA IRTF. The observations were made in three sessions (October 2007, January 2009, and June 2009) at early morning and late afternoon on Venus in the latitude range of ±60°. CO₂ lines at 2.25 μm reveal variations of the cloud aerosol density (∼25%) and scale height near 65 km. The measured reflectivity of Venus at low latitudes is 0.7 at 2.25 μm and 0.028 at 3.66 μm, and the effective CO₂ column density is smaller at 3.66 μm than those at 2.25 μm by a factor of 4. This agrees with the almost conservative multiple scattering at 2.25 μm and single scattering in the almost black aerosol at 3.66 μm. The expected difference is just a factor of (1 − g)⁻¹ = 4, where g = 0.75 is the scattering asymmetry factor for Venus’ clouds. The observed CO mixing ratio is 52 ± 4 ppm near 08:00 and 40 ± 4 ppm near 16:30 at 68 km, and the higher ratio in the morning may be caused by extension of the CO morningside bulge to the cloud tops. The observed weak limb brightening in CO indicates an increase of the CO mixing ratio with altitude. HF is constant at 3.5 ± 0.2 ppb at 68 km in both morningside and afternoon observations and in the latitude range ±60°. Therefore the observations do not favor a bulge of HF, though HF is lighter than CO. Probably a source in the upper atmosphere facilitates the bulge formation. The recent measurements of HCl near 70 km are controversial (0.1 and 0.74 ppm) and require either a strong sink or a strong source of HCl in the clouds. The HCl lines of the (2-0) band are blended by the solar and telluric lines. Therefore we observed the P8 lines of the (1-0) band at 3.44 μm. These lines are spectrally clean and result in the HCl mixing ratio of 0.40 ± 0.03 ppm at 74 km. HCl does not vary with latitude within ±60°. Our observations support a uniformly mixed HCl throughout the Venus atmosphere.     

Latitudinal variations of CO and OCS in the lower atmosphere of Venus from near-infrared nightside spectro-imaging

Spectro-imaging of Venus' nightside in the 2.3-μm window provides a powerful means of probing the lower atmosphere in the 25-40 km altitude range. We present observations recorded at the NASA/IRTF in February 2003 and August 2004, using the SpeX spectro-imager in the 2.1-2.5-μm region. Abundances of CO and OCS have been derived as a function of latitude for different longitudes. The CO abundance increases by about 15% between the equatorial region and higher latitudes (±40°). No longitudinal or temporal variations are observed. The OCS abundance shows the opposite variation in observational sets with sufficient S/N. These variations and anticorrelation are consistent with upwelling motions in the equatorial region and downwelling at higher latitudes.     

A sensitive search for nitric oxide in the lower atmospheres of Venus and Mars: Detection on Venus and upper limit for Mars

High-resolution spectra of Venus and Mars at the NO fundamental band at 5.3 μm with resolving power ν/δν=76,000 were acquired using the TEXES spectrograph at NASA IRTF on Mauna Kea, Hawaii. The observed spectrum of Venus covered three NO lines of the P-branch. One of the lines is strongly contaminated, and the other two lines reveal NO in the lower atmosphere at a detection level of 9 sigma. A simple photochemical model for NO and N at 50-112 km was coupled with a radiative transfer code to simulate the observed equivalent widths of the NO and some CO₂ lines. The derived NO mixing ratio is 5.5±1.5 ppb below 60 km and its flux is . Predissociation of NO at the (0-0) 191 nm and (1-0) 183 nm bands of the δ-system and the reaction with N are the only important loss processes for NO in the lower atmosphere of Venus. The photochemical impact of the measured NO abundance is significant and should be taken into account in photochemical modeling of the Venus atmosphere. Lightning is the only known source of NO in the lower atmosphere of Venus, and the detection of NO is a convincing and independent proof of lightning on Venus. The required flux of NO is corrected for the production of NO and N by the cosmic ray ionization and corresponds to the lightning energy deposition of . For a flash energy on Venus similar to that on the Earth (∼¹⁰⁹ J), the global flashing rate is ∼90 s⁻¹ and ∼6 km⁻² y⁻¹ which is in reasonable agreement with the existing optical observations. The observed spectrum of Mars covered three NO lines of the R-branch. Two of these lines are contaminated by CO₂ lines, and the line at 1900.076 cm⁻¹ is clean and shows some excess over the continuum. Some photochemical reactions may result in a significant excitation of NO (v=1) in the lowest 20 km on Mars. However, quenching of NO (v=1) by CO₂ is very effective below 40 km. Excitation of NO (v=1) in the collisions with atomic oxygen is weak because of the low temperature in the martian atmosphere, and we do not see any explanation of a possible emission of NO at 5.3 μm. Therefore the data are treated as the lack of absorption with a 2 sigma upper limit of 1.7 ppb to the NO abundance in the lower atmosphere of Mars. This limit is above the predictions of photochemical models by a factor of 3.     

Spatially-resolved high-resolution spectroscopy of Venus 2. Variations of HDO, OCS, and SO2 at the cloud tops

While CO, HCl, and HF, that were considered in the first part of this work, have distinct absorption lines in high-resolution spectra and were detected four decades ago, the lines of HDO, OCS, and SO₂ are either very weak or blended by the telluric lines and have not been observed previously by ground-based infrared spectroscopy at the Venus cloud tops. The H₂O abundance above the Venus clouds is typically below the detection limit of ground-based IR spectroscopy. However, the large D/H ratio on Venus facilitates observations of HDO. Converted to H₂O with D/H ≈ 200, our observations at 2722 cm⁻¹ in the Venus afternoon show a H₂O mixing ratio of ∼1.2 ppm at latitudes between ±40° increasing to ±60° by a factor of 2. The observations in the early morning reveal the H₂O mixing ratio that is almost constant at 2.9 ppm within latitudes of ±75°. The measured H₂O mixing ratios refer to 74 km. The observed increase in H₂O is explained by the lack of photochemical production of sulfuric acid in the night time. The recent observations at the P-branch of OCS at 4094 cm⁻¹ confirm our detection of OCS. Four distributions of OCS along the disk of Venus at various latitudes and local times have been retrieved. Both regular and irregular components are present in the variations of OCS. The observed OCS mixing ratio at 65 km varies from ∼0.3 to 9 ppb with the mean value of ∼3 ppb. The OCS scale height is retrieved from the observed limb darkening and varies from 1 to 4 km with a mean value of half the atmospheric scale height. SO₂ at the cloud tops has been detected for the first time by means of ground-based infrared spectroscopy. The SO₂ lines look irregular in the observed spectra at 2476 cm⁻¹. The SO₂ abundances are retrieved by fitting by synthetic spectra, and two methods have been applied to determine uncertainties and detection limits in this fitting. The retrieved mean SO₂ mixing ratio of 350 ± 50 ppb at 72 km favors a significant increase in SO₂ above the clouds since the period of 1980-1995 that was observed by the SOIR occultations at Venus Express. Scale heights of OCS and SO₂ may be similar, and the SO₂/OCS ratio is ∼500 and may be rather stable at 65-70 km under varying conditions on Venus.     

Meridional variations of C2H2 and C2H6 in Jupiter's atmosphere from Cassini CIRS infrared spectra

Hydrocarbons such as acetylene (C₂H₂) and ethane (C₂H₆) are important tracers in Jupiter's atmosphere, constraining our models of the chemical and dynamical processes. However, our knowledge of the vertical and meridional variations of their abundances has remained sparse. During the flyby of the Cassini spacecraft in December 2000, the Composite Infrared Spectrometer (CIRS) instrument was used to map the spatial variation of emissions from 10 to 1400 cm⁻¹ (1000-7 μm). In this paper we analyze a zonally averaged set of CIRS spectra taken at the highest (0.48 cm⁻¹) resolution, firstly to infer atmospheric temperatures in the stratosphere at 0.5-20 mbar via the ν₄ band of CH₄, and in the troposphere at 150-400 mbar, via the H₂ absorption at 600-800 cm⁻¹. Stratospheric temperatures at 5 mbar are generally warmer in the north than the south by 7-8 K, while tropospheric temperatures show no such asymmetry. Both latitudinal temperature profiles however do show a pattern of maxima and minima which are largely anti-correlated between the two levels. We then use the derived temperature profiles to infer the vertical abundances of C₂H₂ and C₂H₆ by modeling tropospheric absorption (∼200 mbar) and stratospheric emission (∼5 mbar) in the C₂H₂ν₅ and C₂H₆ν₉ bands, and also emission of the acetylene (ν₄+ν₅)−ν₄ hotband (∼0.1 mbar). Acetylene shows a distinct north-south asymmetry in the stratosphere, with 5 mbar abundances greatest close to 20° N and decreasing from there towards both poles by a factor of ∼4. At 200 mbar in contrast, acetylene is nearly flat at a level of ∼3×¹⁰⁻⁹. Additionally, the abundance gradient of C₂H₂ between 10 and 0.1 mbar is derived, based on interpolated temperatures at 0.1 mbar, and is found to be positive and uniform with latitude to within errors. Ethane at both 5 and 200 mbar shows increasing VMR towards polar regions of ∼1.75 towards 70° N and ∼2.0 towards 70° S. An explanation for the meridional trends is proposed in terms of a combination of photochemistry and dynamics. Poleward, the decreasing UV flux is predicted to decrease the abundances of C₂H₂ and C₂H₆ by factors of 2.7 and 3.5, respectively, at latitude 70°. However, the lifetime of C₂H₆ in the stratosphere (3×¹⁰¹⁰ s at 5 mbar) is much longer than the dynamical timescale for meridional mixing inferred from Comet SL-9 debris (5-50×¹⁰⁸ s), and therefore the rising abundance towards high latitudes likely indicates that meridional mixing dominates over photochemical effects. For C₂H₂, the opposite occurs, with the relatively short photochemical lifetime (3×¹⁰⁷ s), compared to meridional mixing times, ensuring that the expected photochemical trends are visible.     

Variability of CO concentrations in the Venus troposphere from Venus Express/VIRTIS using a Band Ratio Technique

A fast method is presented for deriving the tropospheric CO concentrations in the Venus atmosphere from near-infrared spectra using the night side 2.3 μm window. This is validated using the spectral fitting techniques of Tsang et al. [Tsang, C.C.C., Irwin, P.G.J., Taylor, F.W., Wilson, C.F., Drossart, P., Piccioni, G., de Kok, R., Lee, C., Calcutt, S.B., and the Venus Express/VIRTIS Team, 2008a. Tropospheric carbon monoxide concentrations and variability on Venus with Venus Express/VIRTIS-M observations. J. Geophys. Res. 113, doi: 10.1029/2008JE003089. E00B08] to show that monitoring CO in the deep atmosphere can be done quickly using large numbers of observations, with minimal effect from cloud and temperature variations. The new method is applied to produce some 1450 zonal mean CO profiles using data from the first eighteen months of operation from the Visible and Infrared Thermal Imaging Spectrometer infrared mapping subsystem (VIRTIS-M-IR) on Venus Express. These results show many significant long- and short-term variations from the mean equator-to-pole increasing trend previously found from earlier Earth- and space-based observations, including a possible North-South dichotomy, with interesting implications for the dynamics and chemistry of the lower atmosphere of Venus.     

SWAS observations of water vapor in the Venus mesosphere

We present the first detections of the ground-state H₂¹⁶O (₁₁₀-₁₀₁) rotational transition (at 556.9 GHz) and the ¹³CO (5-4) rotational transition from the atmosphere of Venus, measured with the Submillimeter Wave Astronomy Satellite (SWAS). The observed spectral features of these submillimeter transitions originate primarily from the 70-100 km altitude range, within the Venus mesosphere. Observations were obtained in December 2002, and January, March, and July 2004, coarsely sampling one Venus diurnal period as seen from Earth. The measured water vapor absorption line depth shows large variability among the four observing periods, with strong detections of the line in December 2002 and July 2004, and no detections in January and March 2004. Retrieval of atmospheric parameters was performed using a multi-transition inversion algorithm, combining simultaneous retrievals of temperature, carbon monoxide, and water profiles under imposed constraints. Analysis of the SWAS spectra resulted in measurements or upper limits for the globally averaged mesospheric water vapor abundance for each of the four observation periods, finding variability over at least two orders of magnitude. The results are consistent with both temporal and diurnal variability, but with short-term fluctuations clearly dominating. These results are fully consistent with the long-term study of mesospheric water vapor from millimeter and submillimeter observations of HDO [Sandor, B.J., Clancy, R.T., 2005. Icarus 177, 129-143]. The December 2002 observations detected very rapid change in the mesospheric water abundance. Over five days, a deep water absorption feature consistent with a water vapor abundance of 4.5±1.5 parts per million suddenly gave way to a significantly shallower absorption, implying a decrease in the water vapor abundance by a factor of nearly 50 in less that 48 h. In 2004, similar changes in the water vapor abundance were measured between the March and July SWAS observing periods, but variability on time scales of less than a week was not detected. The mesospheric water vapor is expected to be in equilibrium with aerosol particles, primarily composed of concentrated sulfuric acid, in the upper haze layers of the Venus atmosphere. If true, moderate amplitude (10-15 K) variability in mesospheric temperature, previously noted in millimeter spectroscopy observations of Venus, can explain the rapid water vapor variability detected by SWAS.     

Photochemistry of the martian atmosphere: Seasonal, latitudinal, and diurnal variations

There is a significant progress in the observational data relevant to Mars photochemistry in the current decade. These data are not covered by and sometimes disagree with the published models. Therefore we consider three types of models for Mars photochemistry. A steady-state model for global-mean conditions is currently the only way to calculate the abundances of long living species (H₂, O₂, and CO). However, our model does not fit the observed CO abundance using gas-phase chemistry and reasonable values of heterogeneous loss of odd hydrogen on the water ice aerosol. The second type of the calculated models is steady-state models for local conditions. The MGS/TES data on temperature profiles, H₂O, and dust are input parameters for these models. The calculations have been made for nine seasonal points spread over the martian year and for twelve latitudes with a step of 10° for each season. The only adopted heterogeneous reaction is a weak loss of H₂O₂ on water ice with probability of 5×¹⁰⁻⁴. The results are in good agreement with the recent observations of the O₂ dayglow at 1.27 μm and the O₃ and H₂O₂ abundances. Global maps of the seasonal and latitudinal behavior of these species have been made. The third type of models is a time-dependent model for local conditions. These models show that odd hydrogen quickly converts to H₂O₂ at the nighttime and the chemistry is switched off while the association of O, the heterogeneous loss of H₂O₂, and eddy diffusion continue. This requires significant changes in the global-mean and local steady-state models discussed above, and these changes have been properly done. The calculated diurnal variations of Mars photochemistry are discussed. The martian photochemistry at low and middle latitudes is significantly different in the aphelion period at LS=10°-130° from that in the remaining part of the year.     

Simultaneous mapping of H2O and H2O2 on Mars from infrared high-resolution imaging spectroscopy

New maps of martian water vapor and hydrogen peroxide have been obtained in November-December 2005, using the Texas Echelon Cross Echelle Spectrograph (TEXES) at the NASA Infra Red Telescope facility (IRTF) at Mauna Kea Observatory. The solar longitude Ls was 332° (end of southern summer). Data have been obtained at 1235-1243 cm⁻¹, with a spectral resolution of 0.016 cm⁻¹ (R=8×¹⁰⁴). The mean water vapor mixing ratio in the region [0°-55° S; 345°-45° W], at the evening limb, is 150±50 ppm (corresponding to a column density of 8.3±2.8 pr-μm). The mean water vapor abundance derived from our measurements is in global overall agreement with the TES and Mars Express results, as well as the GCM models, however its spatial distribution looks different from the GCM predictions, with evidence for an enhancement at low latitudes toward the evening side. The inferred mean H₂O₂ abundance is 15±10 ppb, which is significantly lower than the June 2003 result [Encrenaz, T., Bézard, B., Greathouse, T.K., Richter, M.J., Lacy, J.H., Atreya, S.K., Wong, A.S., Lebonnois, S., Lefèvre, F., Forget, F., 2004. Icarus 170, 424-429] and lower than expected from the photochemical models, taking in account the change in season. Its spatial distribution shows some similarities with the map predicted by the GCM but the discrepancy in the H₂O₂ abundance remains to be understood and modeled.     

Infrared imaging spectroscopy of Mars: H2O mapping and determination of CO2 isotopic ratios

High-resolution infrared imaging spectroscopy of Mars has been achieved at the NASA Infrared Telescope Facility (IRTF) on June 19-21, 2003, using the Texas Echelon Cross Echelle Spectrograph (TEXES). The areocentric longitude was 206°. Following the detection and mapping of hydrogen peroxide H₂O₂ [Encrenaz et al., 2004. Icarus 170, 424-429], we have derived, using the same data set, a map of the water vapor abundance. The results appear in good overall agreement with the TES results and with the predictions of the Global Circulation Model (GCM) developed at the Laboratory of Dynamical Meteorology (LMD), with a maximum abundance of water vapor of 3±1.5×¹⁰⁻⁴(17±9 pr-μm). We have searched for CH₄ over the martian disk, but were unable to detect it. Our upper limits are consistent with earlier reports on the methane abundance on Mars. Finally, we have obtained new measurements of CO₂ isotopic ratios in Mars. As compared to the terrestrial values, these values are: (¹⁸O/¹⁷O)[M/E] = 1.03 ± 0.09; (¹³C/¹²C)[M/E] = 1.00 ± 0.11. In conclusion, in contrast with the analysis of Krasnopolsky et al. [1996. Icarus 124, 553-568], we conclude that the derived martian isotopic ratios do not show evidence for a departure from their terrestrial values.     

Atomic oxygen distribution in the Venus mesosphere from observations of O2 infrared airglow by VIRTIS-Venus Express

This VIRTIS instrument on board Venus Express has collected spectrally resolved images of the Venus nightside limb that show the presence of the (0,0) band of the infrared atmospheric system of O₂ at 1.27 μm. The emission is produced by three-body recombination of oxygen atoms created by photodissociation of CO₂ on the dayside. It is consistently bright so that emission limb profiles can be extracted from the images. The vertical distribution of O₂() may be derived following Abel inversion of the radiance limb profiles. Assuming photochemical equilibrium, it is combined with the CO₂ vertical distribution to determine the atomic oxygen density. The uncertainties on the O density caused by the Abel inversion reach a few percent at the peak, increasing to about 50% near 120 km. We first analyze a case when the CO₂ density was derived from a stellar occultation observed with the SPICAV spectrometer simultaneously with an image of the O₂ limb airglow. In other cases, an average CO₂ profile deduced from a series of ultraviolet stellar occultations is used to derive the O profile, leading to uncertainties on the O density less than 30%. It is found that the maximum O density is generally located between 94 and 115 km with a mean value of 104 km. It ranges from less than 1×¹⁰¹¹ to about 5×¹⁰¹¹ cm⁻³ with a global mean of 2.2×¹⁰¹¹ cm⁻³. These values are in reasonable agreement with the VIRA midnight oxygen profile. The vertical O distribution is generally in good agreement with the oxygen profile calculated with a one-dimensional chemical-diffusive model. No statistical latitudinal dependence of the altitude of the oxygen peak is observed, but the maximum O density tends to decrease with increasing northern latitudes. The latitudinal distribution at a given time exhibits large variations in the O density profile and its vertical structure. The vertical oxygen distribution frequently shows multiple peaks possibly caused by waves or variations in the structure of turbulent transport. It is concluded that the O₂ infrared night airglow is a powerful tool to map the distribution of atomic oxygen in the mesosphere between 90 and 115 km and improve future Venus reference atmosphere models.     

Chemical kinetic model for the lower atmosphere of Venus

A self-consistent chemical kinetic model of the Venus atmosphere at 0-47 km has been calculated for the first time. The model involves 82 reactions of 26 species. Chemical processes in the atmosphere below the clouds are initiated by photochemical products from the middle atmosphere (H₂SO₄, CO, S_x), thermochemistry in the lowest 10 km, and photolysis of S₃. The sulfur bonds in OCS and S_x are weaker than the bonds of other elements in the basic atmospheric species on Venus; therefore the chemistry is mostly sulfur-driven. Sulfur chemistry activates some H and Cl atoms and radicals, though their effect on the chemical composition is weak. The lack of kinetic data for many reactions presents a problem that has been solved using some similar reactions and thermodynamic calculations of inverse processes. Column rates of some reactions in the lower atmosphere exceed the highest rates in the middle atmosphere by two orders of magnitude. However, many reactions are balanced by the inverse processes, and their net rates are comparable to those in the middle atmosphere. The calculated profile of CO is in excellent agreement with the Pioneer Venus and Venera 12 gas chromatographic measurements and slightly above the values from the nightside spectroscopy at 2.3 μm. The OCS profile also agrees with the nightside spectroscopy which is the only source of data for this species. The abundance and vertical profile of gaseous H₂SO₄ are similar to those observed by the Mariner 10 and Magellan radio occultations and ground-based microwave telescopes. While the calculated mean S₃ abundance agrees with the Venera 11-14 observations, a steep decrease in S₃ from the surface to 20 km is not expected from the observations. The ClSO₂ and SO₂Cl₂ mixing ratios are ∼10⁻¹¹ in the lowest scale height. The existing concept of the atmospheric sulfur cycles is incompatible with the observations of the OCS profile. A scheme suggested in the current work involves the basic photochemical cycle, that transforms CO₂ and SO₂ into SO₃, CO, and S_x, and a minor photochemical cycle which forms CO and S_x from OCS. The net effect of thermochemistry in the lowest 10 km is formation of OCS from CO and S_x. Chemistry at 30-40 km removes the downward flux of SO₃ and the upward flux of OCS and increases the downward fluxes of CO and S_x. The geological cycle of sulfur remains unchanged.     

Helium on Mars and Venus: EUVE observations and modeling

Long-exposure spectroscopy of Mars and Venus with the Extreme Ultraviolet Explorer (EUVE) has revealed emissions of He 584 Å on both planets and He 537 Å/O⁺ 539 Å and He⁺ 304 Å on Venus. Our knowledge of the solar emission at 584 Å, eddy diffusion in Mars' upper atmosphere, electron energy distributions above Mars' ionopause, and hot oxygen densities in Mars' exosphere has been significantly improved since our analysis of the first EUVE observation of Mars [Krasnopolsky, Gladstone, 1996, Helium on Mars: EUVE and Phobos data and implications for Mars' evolution, J. Geophys. Res. 101, 15,765-15,772]. These new results and a more recent EUVE observation of Mars are the motivation for us to revisit the problem in this paper. We find that the abundance of helium in the upper atmosphere, where the main loss processes occur, is similar to that in the previous paper, though the mixing ratio in the lower and middle atmosphere is now better estimated at 10±6 ppm. Our estimate of the total loss of helium is almost unchanged at 8×¹⁰²³ s⁻¹, because a significant decrease in the loss by electron impact ionization above the ionopause is compensated by a higher loss in collisions with hot oxygen. We neglect the outgassing of helium produced by radioactive decay of U and Th because of the absence of current volcanism and a very low upper limit to the seepage of volcanic gases. The capture of solar wind α-particles is currently the only substantial source of helium on Mars, and its efficiency remains at 0.3. A similar analysis of EUV emissions from Venus results in a helium abundance in the upper atmosphere which is equal to the mean of the abundances measured previously with two optical and two mass spectrometers, and a derived helium mixing ratio in the middle and lower atmosphere of 9±6 ppm. Helium escape by ionization and sweeping out of helium ions by the solar wind above the ionopause is smaller than that calculated by Prather and McElroy [1983, Helium on Venus: implications for uranium and thorium, Science 220, 410-411] by a factor of 3. However, charge exchange of He⁺ ions with CO₂ and N₂ between the exobase and ionopause and collisions with hot oxygen ignored previously add to the total loss which appears to be at the level of 10⁶ cm⁻² s⁻¹ predicted by Prather and McElroy [1983, Science 220, 410-411]. The loss of helium is compensated by outgassing of helium produced by radioactive decay of U and Th and by the capture of the solar wind α-particles with an efficiency of 0.1. We also compare our derived α-particle capture efficiencies for Mars and Venus with observed X-ray emissions resulting from the charge exchange of solar wind heavy ions with the extended atmospheres on both planets [Dennerl et al., 2002, Discovery of X-rays from Venus with Chandra, Astron. Astrophys. 386, 319-330; Dennerl, 2002, Discovery of X-rays from Mars with Chandra, Astron. Astrophys. 394, 1119-1128]. The emissions from both disk and halo on Mars agree with our calculated values; however, we do not see a reasonable explanation for the X-ray halo emission on Venus. The ratio of the charge exchange efficiencies derived from the disk X-ray emissions of Mars and Venus is similar to the ratio of the capture efficiencies for these planets. The surprisingly bright emission of He⁺ at 304 Å observed by EUVE and Venera 11 and 12 suggests that charge exchange in the flow of the solar wind α-particles around the ionopause is much stronger than in the flow of α-particles into the ionosphere.     

Saturn's latitudinal C2H2 and C2H6 abundance profiles from Cassini/CIRS and ground-based observations

Hydrocarbons in the upper atmosphere of Saturn are known, from Voyager, ground-based, and early Cassini results, to vary in emission intensity with latitude. Of particular interest is the marked increase in hydrocarbon line intensity near the south pole during southern summer, as the increased line intensity cannot be simply explained by the increased temperatures observed in that region since the variations between C₂H₂ and C₂H₆ emission in the south pole region are different. In order to measure the latitudinal variations of hydrocarbons in Saturn's southern hemisphere we have used 3 cm⁻¹ resolution Cassini CIRS data from 2006 and combined this with measurements from the ground in October 2006 at NASA's IRTF using Celeste, an infrared high-resolution cryogenic grating spectrometer. These two data sets have been used to infer the molecular abundances of C₂H₂ and C₂H₆ across the southern hemisphere in the 1-10 mbar altitude region. We find that the latitudinal acetylene profile follows the yearly average mean daily insolation except at the southern pole where it peaks in abundance. Near the equator (5° S) the C₂H₂ abundance at the 1.2 mbar level is (1.6±0.19)×¹⁰−7 and it decreases by a factor of 2.7 from the equator toward the pole. However, at the pole (∼87° S) the C₂H₂ abundance jumps to (1.8±0.3)×¹⁰−7, approximately the equatorial value. The C₂H₆ abundance near the equator at the 2 mbar level is (0.7±0.1)×¹⁰−5 and stays approximately constant until mid-latitudes where it increases gradually toward the pole, attaining a value of (1.4±0.4)×¹⁰−5 there. The increase in ethane toward the pole with the corresponding decrease in acetylene is consistent with southern hemisphere meridional winds [Greathouse, T.K., Lacy, J.H., Bézard, B., Moses, J.I., Griffith, C.A., Richter, M.J., 2005. Icarus 177, 18-31]. The localized increase in acetylene at the pole provides evidence that there is dynamical transport of hydrocarbons from the equator to the southern pole.     

Half-widths, their temperature dependence, and line shifts for the HDO-CO2 collision system for applications to CO2-rich planetary atmospheres

Measurements of water vapor in the atmospheres of Venus or Mars by spectroscopic techniques in the infrared range are being made routinely by instruments onboard the Venus Express and the Mars Reconnaissance Orbiter. The interpretation of these measurements in the 2250-4450 cm⁻¹ region is being complicated by the presence of HDO lines absorbing radiation in this region. These spectra cannot be modeled properly because line shape parameters for CO₂ broadening (principal gas in these atmospheres) of HDO are not available. Here semi-classical line shape calculations for the HDO-CO₂ collision system are made using the Robert-Bonamy formalism for some 2300 rotational band transitions of HDO. From these calculations, the half-width, its temperature dependence, and the line shift are determined to aid in the reduction of the measured spectra. These data will greatly reduce the uncertainty of the reduced profiles from the Venus and Mars measurements and will also allow better estimates of the D/H ratio on these planets.     

Water vapor abundance in Venus' middle atmosphere from Pioneer Venus OIR and Venera 15 FTS measurements

The Pioneer Venus Orbiter Infrared Radiometer and Venera 15 Fourier Transform Spectrometer observations of thermal emission from Venus' middle atmosphere between 10° S and 50° N have been independently re-analyzed using a common method to determine global maps of temperature, cloud optical depth, and water vapor abundance. The spectral regions observed include the strong 15 μm carbon dioxide band and the 45 μm fundamental rotational water band. The different spatial and spectral resolutions of the two instruments have necessitated the development of flexible analysis tools. New radiative transfer and retrieval models have been developed for this purpose based on correlated-k absorption tables calculated with up-to-date spectral line data. The common analysis of these two sets of observations has hence been possible for the first time. From the PV OIR observations, the cloud-top unit optical depth pressure showed a minimum of ∼110±10 mbars in the evening equatorial region and a maximum of ∼160±12 mbars in the morning mid-latitude regions. From the Venera 15 FTS spectra, the cloud-top pressure was found to increase from morning values of ∼120±10 to 200±30 mbars in the late afternoon/early evening region. The cloud-top water vapor abundances observed by the PV OIR instrument were found to fluctuate from 10±5 ppm at night up to 90±15 ppm in the equatorial cloud-top region shortly after the sub-solar point. The mean Venera 15 FTS water vapor abundances were found to be 12±5 ppm with only a slight enhancement over the equatorial latitude bands and no clear day-night distinction. The common analysis of these two sets of observations broadly validates previously published individual findings. The differences in the retrieved atmospheric state can no longer be attributed to radiative transfer modeling bias and suggest significant temporal variability in the middle atmosphere of Venus.     

Vertical profile of H2SO4 vapor at 70-110 km on Venus and some related problems

The vertical profile of H₂SO₄ vapor is calculated using current atmospheric and thermodynamic data. The atmospheric data include the H₂O profiles observed at 70-112 km by the SOIR solar occultations, the SPICAV-UV profiles of the haze extinction at 220 nm, the VeRa temperature profiles, and a typical profile of eddy diffusion. The thermodynamic data are the saturated vapor pressures of H₂O and H₂SO₄ and chemical potentials of these species in sulfuric acid solutions. The calculated concentration of sulfuric acid in the cloud droplets varies from 85% at 70 km to a minimum of 70% at 90 km and then gradually increasing to 90-100% at 110 km. The H₂SO₄ vapor mixing ratio is ∼10⁻¹² at 70 and 110 km with a deep minimum of 3 × 10⁻¹⁸ at 88 km. The H₂O-H₂SO₄ system matches the local thermodynamic equilibrium conditions up to 87 km. The column photolysis rate of H₂SO₄ is 1.6 × 10⁵ cm⁻² s⁻¹ at 70 km and 23 cm⁻² s⁻¹ at 90 km. The calculated abundance of H₂SO₄ vapor at 90-110 km and its photolysis rate are smaller than those presented in the recent model by Zhang et al. (Zhang, X., Liang, M.C., Montmessin, F., Bertaux, J.L., Parkinson, C., Yung, Y.L. [2010]. Nat. Geosci. 3, 834-837) by factors of 10⁶ and 10⁹, respectively. Assumptions of 100% sulfuric acid, local thermodynamic equilibrium, too warm atmosphere, supersaturation of H₂SO₄ (impossible for a source of SO_X), and cross sections for H₂SO₄·H₂O (impossible above the pure H₂SO₄) are the main reasons of this huge difference. Significant differences and contradictions between the SPICAV-UV, SOIR, and ground-based submillimeter observations of SO_X at 70-110 km are briefly discussed and some weaknesses are outlined. The possible source of high altitude SO_X on Venus remains unclear and probably does not exist.     

Laboratory measurements of the microwave opacity of hydrochloric acid vapor in a carbon dioxide atmosphere

Laboratory measurements of the microwave opacity of HCl in a CO₂ atmosphere have been conducted in the S (13.3 cm), X (3.6 cm), and K (1.4 cm) microwave bands at a pressure of 7.2 bar and at two different mixing ratios. The results are consistent with an opacity model employing the Van Vleck-Weisskopf lineshape applied to the published submillimeter line intensities of HCl (JPL Catalog [Pickett et al., 1998, J. Quant. Spectrosc. Rad. Trans. 60, 883-890]) and empirically fitted with a modeled parameter for CO₂ broadening. Based on the deep atmospheric abundance of HCl inferred from near-infrared measurements [Dalton et al., 2000, Bull. Am. Astron. Soc. 32, 1120], the resulting modeled HCl opacity is constrained to have a small effect on the overall microwave absorption spectrum of Venus, but can be used in developing a more accurate radiative transfer model.     

Venus night airglow: Ground-based detection of OH, observations of O2 emissions, and photochemical model

Venus nightglow was observed at NASA IRTF using a high-resolution long-slit spectrograph CSHELL at LT = 21:30 and 4:00 on Venus. Variations of the O₂ airglow at 1.27 μm and its rotational temperature are extracted from the observed spectra. The mean O₂ nightglow is 0.57 MR at 21:30 at 35°S-35°N, and the temperature increases from 171 K near the equator to ∼200 K at ±35°. We have found a narrow window that covers the OH (1-0) P1(4.5) and (2-1) Q1(1.5) airglow lines. The detected line intensities are converted into the (1-0) and (2-1) band intensities of 7.2 ± 1.8 kR and <1.4 kR at 21:30 and 15.5 ± 2 kR and 4.7 ± 1 kR at 4:00. The f-component of the (1-0) P1(4.5) line has not been detected in either observation, possibly because of resonance quenching in CO₂. The observed Earth’s OH (1-0) and (2-1) bands were 400 and 90 kR at 19:30 and 250 and 65 kR at 9:40, respectively. A photochemical model for the nighttime atmosphere at 80-130 km has been made. The model involves 61 reactions of 24 species, including odd hydrogen and chlorine chemistries, with fluxes of O, N, and H at 130 km as input parameters. To fit the OH vibrational distribution observed by VEX, quenching of OH (v > 3) in CO₂ only to v ? 2 is assumed. According to the model, the nightside-mean O₂ emission of 0.52 MR from the VEX and our observations requires an O flux of 2.9 × 10¹² cm⁻² s⁻¹ which is 45% of the dayside production above 80 km. This makes questionable the nightside-mean O₂ intensities of ∼1 MR from some observations. Bright nightglow patches are not ruled out; however, the mean nightglow is ∼0.5 MR as observed by VEX and supported by the model. The NO nightglow of 425 R needs an N flux of 1.2 × 10⁹ cm⁻² s⁻¹, which is close to that from VTGCM at solar minimum. However, the dayside supply of N at solar maximum is half that required to explain the NO nightglow in the PV observations. The limited data on the OH nightglow variations from the VEX and our observations are in reasonable agreement with the model. The calculated intensities and peak altitudes of the O₂, NO, and OH nightglow agree with the observations. Relationships for the nightglow intensities as functions of the O, N, and H fluxes are derived.     

Detection of new hydrocarbons in Uranus' atmosphere by infrared spectroscopy

Ethane (C₂H₆), methylacetylene (CH₃C₂H or C₃H₄) and diacetylene (C₄H₂) have been discovered in Spitzer 10-20 μm spectra of Uranus, with 0.1-mbar volume mixing ratios of (1.0±0.1)×¹⁰⁻⁸, (2.5±0.3)×¹⁰⁻¹⁰, and (1.6±0.2)×¹⁰⁻¹⁰, respectively. These hydrocarbons complement previously detected methane (CH₄) and acetylene (C₂H₂). Carbon dioxide (CO₂) was also detected at the 7-σ level with a 0.1-mbar volume mixing ratio of (4±0.5)×¹⁰⁻¹¹. Although the reactions producing hydrocarbons in the atmospheres of giant planets start from radicals, the methyl radical (CH₃) was not found in the spectra, implying much lower abundances than in the atmospheres of Saturn or Neptune where it has been detected. This finding underlines the fact that Uranus' atmosphere occupies a special position among the giant planets, and our results shed light on the chemical reactions happening in the absence of a substantial internal energy source.