Photochemistry of the stratosphere of Venus: Implications for atmospheric evolution

The photochemistry of the stratosphere of Venus was modeled using an updated and expanded chemical scheme, combined with the results of recent observations and laboratory studies. We examined three models, with H₂ mixing ratio equal to 2 × 10^?5, 5 × 10^?7, and 1 × 10^?13, respectively. All models satisfactorily account for the observations of CO, O₂, O₂(¹Δ), and SO₂ in the stratosphere, but only the last one may be able to account for the diurnal behavior of mesospheric CO and the uv albedo. Oxygen, derived from CO₂ photolysis, is primarily consumed by CO₂ recombination and oxidation of SO₂ to H₂SO₄. Photolysis of HCl in the upper stratosphere provides a major source of odd hydrogen and free chlorine radicals, essential for the catalytic oxidation of CO. Oxidation of SO₂ by O occurs in the lower stratosphere. In the high-H₂ model (model A) the OO bond is broken mainly by S + O₂ and SO + HO₂. In the low-H₂ models additional reactions for breaking the OO bond must be invoked: NO + HO₂ in model B and ClCO + O₂ in model C. It is shown that lightning in the lower atmosphere could provide as much as 30 ppb of NO_x in the stratosphere. Our modeling reveals a number of intriguing similarities, previously unsuspected, between the chemistry of the stratosphere of Venus and that of the Earth. Photochemistry may have played a major role in the evolution of the atmosphere. The current atmosphere, as described by our preferred model, is characterized by an extreme deficiency of hydrogen species, having probably lost the equivalent of 10²–10³ times the present hydrogen content.     

Diurnal CO variations in the Venus mesophere from CO microwave spectra

Microwave spectra of carbon monoxide (¹²CO) in the mesosphere of Venus were measured in December 1978, May and December 1980, and January, September, and November 1982. These spectra are analyzed to provide mixing profiles of CO in the Venus mesosphere and best constrain the mixing profile of CO between ～ 100 and 80 km altitude. From the January 1982 measurement (which, of all our spectra, best constrains the abundance of CO below 80 km altitude) we find an upper limit for the CO mixing ratio below 80 km altitude that is two to three times smaller than the stratospheric (～65 km) value of 4.5 ± 1.0 × 10^?5 determined by P. Connes, J. Connes, L.D. Kaplan, and W. S. Benedict (1968, Astrophys. J.152, 731–743) in 1967, indicating a possible long-term change in the lower atmospheric concentration of CO. Intercomparison among the individual CO profiles derived from our spectra indicates considerable short-term temporal and/or spatial variation in the profile of CO mixing in the Venus mesosphere above 80 km. A more complete comparison with previously published CO microwave spectra from a number of authors specifies the basic diurnal nature of mesospheric CO variability. CO abundance above ～ 95 km in the Venus atmosphere shows approximately a factor of 2–4 enhancement on the nightside relative to the dayside of Venus. Peak nightside CO abundance above ～95 km occurs very near to the antisolar point on Venus (local time of peak CO abundance above ～95 km occurs at 0.6_?0.6^+0.7 hr after midnight on Venus), strongly suggesting that retrograde zonal flow is substantially reduced at an altitude of 100 km in the Venus mesosphere. In contrast, CO abundances between 80 and 90 km altitude show a maximum that is shifted from the antisolar point toward the morningside of Venus (local time of peak CO abundance between 80 and 90 km occurs at 8.5 ± 1.0 hr past midnight on Venus). The magnitude of the diurnal variation of CO abundance between 80 and 90 km is again, approximately a factor of 2–4. Disk-averaged spectra of Venus do not determine the exact form for the diurnal distribution of CO in the Venus mesosphere as indicated by comparison of synthetic spectra, based upon model distributions, and the measured spectra. However, the offset in phase for the diurnal variation for the >95 km and 80–90-km-altitude regions requires an asymmetric (in solar zenith angle) distribution.     

Ozone and photochemistry of the Martian lower atmosphere

The calculation of number densities of CO₂, H₂O and N₂ photolysis products was carried out for the Martian atmosphere at heights up to 60 km. The ozone distributed in the atmosphere as a layer of 10 km width with [O₃] max = 2.5 × 10⁹ cm³ at height of 35 km which agree well with the results of u.v. observations on the evening terminator from the Mars-5 satellite. The calculated densities of O₂, CO and H₂O are also in good agreement with the measured data. The eddy diffusion coefficient is equal to 3 × 10⁶ in the troposphere (h ? 30 km) and 10⁸ cm² s^?1 above 40 km. The dependence of the total ozone content on water vapour amount in the atmosphere is considered; the hypothesis about the influence of water ice aerosol on the ozone formation is proposed to explain the high concentrations of ozone in the morning.     

Lightnings and nitric oxide on Venus

In an updating of energy characteristics of lightnings on Venus obtained from Venera-9 and -10 optical observations, the flash energy is given as 8 × 10⁸ J and the mean energy release of lightnings is 1 erg cm^?2 s which is 25 times as high as that on the Earth. Lightnings were observed in the cloud layer. The stroke rate in the near-surface atmosphere is less than 5 s^?1 over the entire planet if the light energy of the stroke exceeds 4 × 10⁵ J and less than 15 s^?1 for (1–4) × 10⁵ J.The average NO production due to lightnings equals 5 × 10⁸ cm^?2 s^?1, the atomic nitrogen production is equal to 7 × 10⁹ cm^?2s^?1,the N flux toward the nightside is 3.2 × 10⁹ cm^?2s^?1, the number densities [N] = 3 × 10⁷cm^?3 and [NO] = 1.8 × 10⁶cm^?3 at 135 km. Almost all NO molecules in the upper atmosphere vanish interacting with N and the resulting NO flux at 90-80 km equals 5 × 10⁵cm^?2s^?1, which is negligibly small as compared with lightning production. If the predissociation at 80–90 km is regarded as the single sink of NO, its mixing ratio, f_NO, is 4 × 10^?8, for the case of a surface sink f_NO = 0.8 × 10^?9 at 50 km. Excess amounts, $\text{f}{}_{\mathrm{<mtext>NO</mtext>}}\text{? 4 × 10}{}^{\mathrm{?8}}$, may exist in the thunderstorm region.     

Photochemistry of nitrogen in the Martian atmosphere

Models are developed for the photochemistry of a CO₂H₂ON₂ atmosphere on Mars and estimates are given for the concentrations of N, NO, NO₂, NO₃, N₂O₅, HNO₂, HNO₃, and N₂O as a function of altitude. Nitric oxide is the most abundant form of odd nitrogen, present with a mixing ratio relative to CO₂ of order 10^?8. Deposition rates for nitrite and nitrate minerals could be as large as 3× 10⁵ N equivalent atoms cm^?2 sec^?1 under present conditions and may have been higher in the past.     

Progressive Build Up Of Co 2 In The AtmosphereOf Venus Through Multiple Volcanic Resurfacing Events

In absence of other mechanisms, the main input of CO₂into the Venusian atmosphere is via volcanic out gassing. Since Venus can be regarded as a planet-wide large igneous province, we can expect large quantities of CO₂ being transferred into its atmosphere via volcanic out gassing. We have quantified the maximum possible amount of CO₂ that can be out gassed via a single massive episode of resurfacing of the planet. This figure (5.6 × 10¹⁹ kg of CO₂) is about 8 times smaller than the total CO₂ present in the Venusian atmosphere (4.55 × 10²⁰ kg CO₂). The lack of planet-wide, efficient mechanisms for the recycling of CO₂ on Venus indicates that CO₂ has progressively accumulated in the atmosphere. Based on these considerations we suggest that the “equivalent” to eight global resurfacing episodes would be required to account for the present values of CO_{2 atmosphere}. This revised version was published online in July 2006 with corrections to the Cover Date.     

Variability of carbon monoxide in the mars atmosphere

In January of 1982 we measured a microwave spectrum of CO in the Martian atmosphere utilizing the rotational J = 1 → 2 transition of CO. We have analyzed data and reanalyzed the microwave spectra of R. K. Kakar, J. W. Waters, and W. J. Wilson, (Science196, 1090–1091, 1977, measured in 1975) and J. C. Good and F. P. Schloerb, (Icarus47, 166–172, 1981 measured in 1980) in order to constrain estimates of the temporal variability of CO abundance in the Martian atmosphere. Our values of CO column density from the data of Karar et al., Good and Schloerb, and our own are 1.7 ± 0.9 × 10²⁰, 3.0 ± 1.0 × 10²⁰, and 4.6 ± 2.0 × 10²⁰cm^?2, respectively. The most recent estimate of CO column density from the 1967 infrared spectra of J. Connes, P. Connes, and J.P. Maillard, (Atlas de Spectres Infarouges de Venus, Mars, Jupiter, et Saturne, Editions due Centre National de la Recherche Scientifique, Paris, 1969), is 2.0 ± 0.8 × 10²⁰ cm^?2 (L.D.G. Young and A.T. Young, Icarus30, 75–79, 1977). The large uncertainties given for the microwave measurements are due primarily to uncertainty in the difference between the continuum brightness temperature and atmospheric temperatures of Mars. We have accurately calculated the variation among the observations of the continuum (surface) brightness temperature of Mars, which is primaroly a function of the observed aspect of Mars. A more difficult problem to consider is variability of global atmospheric temperatures among the observations, particularly the effects of global dust storms and the ellipticity of the orbit of Mars. The large bars accompanying our estimates of CO column density from the three sets of microwave measurements are primarily caused by an assumed uncertainty of ±10°K in our atmospheric temperature model due to possible dust in the atmosphere. A qualitative consideration of seasonal variability of global atmospheric temperatures among the measurements suggests that there is not strong evidence for variability of the column abundance of CO on Mars, although variability of 0–100% over a time scale of several years is allowed by the data set. The implication for the variability of Mars O₂ is, crudely, a factor of two less. We found that the altitude distribution of CO in the atmosphere of Mars was not well constrained by any of the spectra, although our spectrum was marginally better fitted by an altitude increasing profile of CO mixing ratios.     

Spectroscopy of molecular oxygen in the Atmospheres of Venus and Mars

The abundances of molecular oxygen in the atmospheres of Venus and Mars are sensitive to fundamental photochemical processes. A new upper limit is reported for the molecular oxygen mixing ratio (O₂/CO₂ < × 10^?7) in the integrated column above the visible cloud tops of Venus, based on spectroscopic observations carried out in early spring, 1982. During the same observing period, an O₂ column abundance of 8.5 cm-am for the atmosphere of Mars was measured, slightly below the O₂ abundances measured a decade earlier.     

Nonthermal escape of hydrogen and deuterium from Venus and implications for loss of water

Nonthermal escape processes responsible for the escape of hydrogen and deuterium from Venus are examined for present and past atmospheres. Three mechanisms are important for the escape of hydrogen from the present atmosphere: (a) charge exchange of plasmaspheric H⁺ with exospheric H, (b) impact of exospheric hot O atoms on H, and (c) ion molecule reactions involving O⁺ and H₂. However, in the past when the H abundance was higher, the charge-exchange mechanism would be the strongest. The H escape flux increases rapidly with increasing hydrogen abundance in the upper atmosphere and saturates at a value of 1 × 10¹⁰ cm^?2 sec^?1 emerging primarily from the day side when the H mixing ratio at the homopause is 2 × 10^?3. This corresponds to an H₂O mixing ratio of 1 × 10^?3 at the cold trap and ～15% at the surface. Deuterium would also escape by the charge-exchange mechanism and a D/H enrichment by a factor of ～1000 over the nonthermal escape regime is expected, which could have lasted over the last 3 billion years. Coincidentally, the onset of hydrodynamic flow leading to efficient H escape occurs just at the H₂O mixing ratio at which the charge-exchange escape flux saturates. Thus it is possible that Venus has lost an Earth-equivalent ocean of water over geologic time. If so, either the D/H enrichment has been kept low by modest outgassing of juvenile water or Venus started out with a D/H ratio of ～4.0 × 10^?6.     

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.     

Some problems in Venus' aeronomy

Models are presented for the height distribution of various photochemically active gases in Venus' upper atmosphere. Attention is directed to the chemistry and vertical transport of odd hydrogen (H, OH, HO₂, H₂O₂), odd oxygen (O, O₃), free chlorine (Cl, ClO, ClOO, Cl₂), CO, O₂, H₂ and H₂O. Supply of O₂ may play a limiting role in the formation of a possible H₂SO₄ cloud on Venus. The supply rate is influenced by both chemical and dynamical processes in the stratosphere, and an analysis of recent spectroscopic data for O₂ implies a lower limit to the appropriate eddy coefficient of about 3 × 10⁵ cm²/sec. The abundances of thermospheric O and CO are determined largely by vertical mixing, and an analysis of Mariner 10 measurements of Venus' Lyman α airglow suggests that the eddy coefficient in the lower thermosphere may be as large as 5 × 10⁷cm²sec. The corresponding values for the mixing ratios of O and CO at the ionospheric peak are approximately 1 per cent. The Lyman α data could be reconciled with larger values for thermospheric O, and smaller values for the vertical eddy coefficient, if non-thermal loss processes were to play a dominant role in hydrogen escape, and if the corresponding flux were to exceed 10⁷ atoms/cm²/sec. A sink of this magnitude would imply major depletion of Venus' atmospheric water over geologic time, and would appear to require mixing ratios of H₂O in the lower atmosphere in excess of 4 × 10^?4. The extensive component to the Lyman α emission measured by Mariner 5 may be due to resonance scattering of sunlight by hot atoms formed by charge transfer with O⁺. The H scale height, therefore, may reflect the temperature of positive ions in Venus' topside ionosphere.     

An improved Venus cloud model

A simple radiative-transfer theory that allows for the change in the absorptions of sulfur and carbon dioxide with depth in the atmosphere of Venus can account simultaneously for (1) the spectral reflectance of Venus; (2) the wavelength dependence of contrast in uv cloud features; (3) the CO₂ line profile; (4) the change in slope of the curve of growth from the 7820- to the 10488-Å CO₂ bands; and (5) the rotational temperature near 246°K found for all CO₂ bands. The model cloud consists of 1-μm sulfuric-acid particles, which are well mixed between about 64 km and the 49-km cloud base found by Veneras 9 and 10, plus an overlapping cloud of much larger sulfur particles that extends down to the 35-km cloud base found by Venera 8. The mixing ratios (by number of molecules) below about 64 km are: H₂O, 2 × 10^?4; H₂SO₄, 10^?5; and sulfur, 10^?4. Although the cloud contains an order of magnitude more sulfur than sulfuric acid, the sulfur particles are an order of magnitude larger, and so have only about 1% of the number density of the acid droplets. The “black-white” radiative-transfer model assumes perfectly conservative scattering above the level (which depends on wavelength) where an absorber becomes “black” due to the local temperature and pressure. So-called homogeneous scattering models are inherently self-contradictory, and are inapplicable to planetary atmospheres; the vertical inhomogeneity is an essential feature that must be modeled correctly. The pressure of CO₂ line formation is about half the pressure in the region where uv markings occur.     

Possible production by lightning of aerosols and trace gases in Titan's atmosphere

Although lightning has not yet been observed in Titan's atmosphere, the presence of condensable vapors and the deposition of a significant amount of solar energy at the surface suggest the possibility of lightning activity. Based on an understanding of the relationship of lightning activity to the amount of convective energy available on Titan, a lightning energy dissipation rate of 4 × 10^?6, W/m² can be expected. This value is much lower than that for Earth or Jupiter, and is a result of both the reduced solar flux at Titan and the absorption of sunlight by the aerosols that lie above the convective layer. For this dissipation rate, the amount of HCN and C₂N₂ produced by lightning should be greater than that by solar UV, but could be less than that produced by electron precipitation and galactic cosmic rays. Equilibrium calculations indicate that large mole fractions of elemental solid phase carbon will also be produced. Using a simplified model of aerosol formation, coagulation, and settling, it is estimated that a lightning-produced aerosol could have a typical optical depth of 10^?2, with values as high as 0.1. The accumulation of soot over geological time might reach a meter or more in depth.     

Oxidation of CO on surface hematite in high CO2 atmospheres

We propose a mechanism for the oxidation of gaseous CO into CO₂ occurring on the surface mineral hematite (Fe₂O₃(s)) in hot, CO₂-rich planetary atmospheres, such as Venus. This mechanism is likely to constitute an important source of tropospheric CO₂ on Venus and could at least partly address the CO₂ stability problem in Venus’ stratosphere, since our results suggest that atmospheric CO₂ is produced from CO oxidation via surface hematite at a rate of 0.4 petagrammes (Pg) CO₂ per (Earth) year on Venus which is about 45% of the mass loss of CO₂ via photolysis in the Venusian stratosphere. We also investigated CO oxidation via the hematite mechanism for a range of planetary scenarios and found that modern Earth and Mars are probably too cold for the mechanism to be important because the rate-limiting step, involving CO(g) reacting onto the hematite surface, proceeds much slower at lower temperatures. The mechanism may feature on extrasolar planets such as Gliese 581c or CoRoT-7b assuming they can maintain solid surface hematite which, e.g. starts to melt above about 1200 K. The mechanism may also be important for hot Hadean-type environments and for the emerging class of hot Super-Earths with planetary surface temperatures between about 600 and 900 K.     

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.     

Mars: Photodesorption from mineral surfaces and its effects on atmospheric stability

A mechanism has been proposed for uv-accelerated desorption from Fe²⁺ sites on mineral surfaces that satisfies kinetic constraints determined in the laboratory by Huguenin. The process is an integral step of the photochemical weathering mechanism for producing dust on Mars, and it now appears that it may play primary roles in stabilizing CO₂ against dissociation by sunlight and in controlling the oxidation state of the atmosphere. We propose that adsorption occurs at octahedrally coordinated Fe²⁺ surface sites to form seven-coordinate transition-state complexes. These complexes acquire 16–18 kcal mole^?1 of ligand field stabilization energy. During illumination (λ ≤ 0.35 μm), electrons are photoemitted from the surfaced Fe²⁺, temporarily oxidizing them to Fe³⁺. Fe³⁺ has no ligand field stabilization energy, and the complexes lose 16–18 kcal mole^?1 of stabilization energy. This is a large fraction of the 19- to 28-kcal mole^?1 activation energy for dissociating the complexes, and desorption should proceed spontaneously. The gases that were observed to undergo adsorption-photodesorption include O₂, CO₂, CO, H₂O, N₂, and Ar. Photodesorption can drive several catalytic reactions, one of which is the oxidation of CO to CO₂. The rate of this reaction should be limited by the supply of CO and O₂ to the surface to ～2 × 10¹² cm^?2 sec^?1 (column photodissociation rate of CO₂). By including this surface reaction in models of Martian atmospheric CO₂ chemistry, CO₂ can be stabilized against photodissociation with eddy diffusion coefficients of only 3 × 10⁵?1 × 10⁷ cm² sec^?1 below 40 km, raising to ～ 10⁹ cm² sec^?1 at 140 km. Odd hydrogen is not needed to catalyze the oxidation of CO below 40 km, and odd hydrogen mixing ratios need only to be $\text{f}{}_{\mathrm{<mtext>H</mtext>}}\text{? 10}{}^{\mathrm{?10}}$ to depress ozone concentrations below the observed upper limit in equatorial regions. Another catalytic reaction that should be driven by photodesorption on Mars is $\text{20}\text{H}{}^{\mathrm{?}}{}_{\mathrm{<mtext>(</mtext><mtext>ads</mtext><mtext>)</mtext>}}\text{→}\text{H}{}_2\text{O}\text{+}\text{1}\text{2}\text{O}{}_{\mathrm{2(g)}}\text{+ 2e}{}^{\mathrm{?}}{}_{\mathrm{<mtext>crystal</mtext>}}$. This is an important source of atmospheric O₂, amounting to 7 × 10¹³?2 × 10¹⁷ O₂ molecules cm^?2 yr^?1, and it could have a significant effect on atmospheric oxidation state.     

Atomic oxygen on the Venus nightside: Global distribution deduced from airglow mapping

The Visible and Infra-Red Thermal Imaging Spectrometer (VIRTIS) instrument on board the Venus Express spacecraft has measured the O₂(a¹Δ) nightglow distribution at 1.27 μm in the Venus mesosphere for more than two years. Nadir observations have been used to create a statistical map of the emission on Venus nightside. It appears that the statistical 1.6 MR maximum of the emission is located around the antisolar point. Limb observations provide information on the altitude and on the shape of the emission layer. We combine nadir observations essentially covering the southern hemisphere, corrected for the thermal emission of the lower atmosphere, with limb profiles of the northern hemisphere to generate a global map of the Venus nightside emission at 1.27 μm. Given all the O₂(a¹Δ) intensity profiles, O₂(a¹Δ) and O density profiles have been calculated and three-dimensional maps of metastable molecular and atomic oxygen densities have been generated. This global O density nightside distribution improves that available from the VTS3 model, which was based on measurements made above 145 km. The O₂(a¹Δ) hemispheric average density is 2.1 × 10⁹ cm^?3, with a maximum value of 6.5 × 10⁹ cm^?3 at 99.2 km. The O density profiles have been derived from the nightglow data using CO₂ profiles from the empirical VTS3 model or from SPICAV stellar occultations. The O hemispheric average density is 1.9 × 10¹¹ cm^?3 in both cases, with a mean altitude of the peak located at 106.1 km and 103.4 km, respectively. These results tend to confirm the modeled values of 2.8 × 10¹¹ cm^?3 at 104 km and 2.0 × 10¹¹ cm^?3 at 110 km obtained by Brecht et al. [Brecht, A., Bougher, S.W., Gérard, J.-C., Parkinson, C.D., Rafkin, S., Foster, B., 2011a. J. Geophys. Res., in press] and Krasnopolsky [Krasnopolsky, V.A., 2010. Icarus 207, 17–27], respectively. Comparing the oxygen density map derived from the O₂(a¹Δ) nightglow observations, it appears that the morphology is very different and that the densities obtained in this study are about three times higher than those predicted by the VTS3 model.     

Observation of DCl and upper limit to NH3 on Venus

To search for DCl in the Venus atmosphere, a spectrum near the D³⁵Cl (1–0) R4 line at 2141.54 cm^?1 was observed using the CSHELL spectrograph at NASA IRTF. Least square fitting to the spectrum by a synthetic spectrum results in a DCl mixing ratio of 17.8 ± 6.8 ppb. Comparing to the HCl abundance of 400 ± 30 ppb (Krasnopolsky [2010a] Icarus, 208, 314–322), the DCl/HCl ratio is equal to 280 ± 110 times the terrestrial D/H = 1.56 × 10^?4. This ratio is similar to that of HDO/H₂O = 240 ± 25 times the terrestrial HDO/H₂O from the VEX/SOIR occultations at 70–110 km. Photochemistry in the Venus mesosphere converts H from HCl to that in H₂O with a rate of 1.9 × 10⁹ cm^?2 s^?1 (Krasnopolsky [2012] Icarus, 218, 230–246). The conversion involves photolysis of HCl; therefore, the photochemistry tends to enrich D/H in HCl and deplete in H₂O. Formation of the sulfuric acid clouds may affect HDO/H₂O as well. The enriched HCl moves down by mixing to the lower atmosphere where thermodynamic equilibriums for H₂ and HCl near the surface correspond to D/H = 0.71 and 0.74 times that in H₂O, respectively. Time to establish these equilibriums is estimated at ～3 years and comparable to the mixing time in the lower atmosphere. Therefore, the enriched HCl from the mesosphere gives D back to H₂O near the surface. Comparison of chemical and mixing times favors a constant HDO/H₂O up to ～100 km and DCl/HCl equal to D/H in H₂O times 0.74.Ammonia is an abundant form of nitrogen in the reducing environments. Thermodynamic equilibriums with N₂ and NO near the surface of Venus give its mixing ratio of 10^?14 and 6 × 10^?7, respectively. A spectrum of Venus near the NH₃ line at 4481.11 cm^?1 was observed at NASA IRTF and resulted in a two-sigma upper limit of 6 ppb for NH₃ above the Venus clouds. This is an improvement of the previous upper limit by a factor of 5. If ammonia exists at the ppb level or less in the lower atmosphere, it quickly dissociates in the mesosphere and weakly affects its photochemistry.     

Martian CO abundance from the J = 1 → 0 rotational transition: Evidence for temporal variations

Mars was observed in the CO (J = 1 → 0) 2.6-mm wavelength line between 29 March and 1 April, 1980. The data were analyzed using a model atmosphere based on Viking measurements. A least-squares fit of the model to the observed line profile yielded an average CO mixing ratio of (3.2 ± 1.1) × 10^?3. This value is four times larger than that obtained by L. D. Kaplan, J. Connes, and P. Connes, 1969 (Astrophys. J.157, L187-L195) from analysis of an infrared spectrum obtained in 1967 by J. Connes, P. Connes, and J. P. Maillard, 1969 (Atlas of Near Infrared Spectra of Venus, Mars, Jupiter, and Saturn, Centre National de la Recherche Scientifique, Paris). Models of the Martian atmospheric chemistry indicate that this implied temporal variation could easily exist and that it would be due primarily to variations in the abundance of H₂O.