The ultraviolet absorption band at 2175 Å: Correlations with other interstellar features

The hump in the ultraviolet part of the interstellar extinction curve is interpreted as a broad diffuse absorption band. Its equivalent width is estimated for 36 stars by means of OAO-2 data. The equivalent widths are correlated with the following parameters: colour excessE(B-V), colour excessE(B – V), depth of the band m _max, equivalent widths of the diffuse bands at 5780 and 6284 Å, and the column density of neutral hydrogenN _HI. The physical parameters half-width and oscillator strength of the band at 2175 Å are estimated.     

Measurements of the helium 584 Å airglow during the Cassini flyby of Venus

The helium resonance line at 584 Å has been observed with the UltraViolet Imaging Spectrograph (UVIS) Extreme Ultraviolet channel during the flyby of Venus by Cassini at a period of high solar activity. The brightness was measured along the disk from the morning terminator up to the bright limb near local noon. The mean disk intensity was ∼320 R, reaching ∼700 R at the bright limb. These values are slightly higher than those determined from previous observations. The sensitivity of the 584 Å intensity to the helium abundance is analyzed using recent cross-sections and solar irradiance measurements at 584 Å. The intensity distribution along the UVIS footprint on the disk is best reproduced using the EUVAC solar flux model and the helium density distribution from the VTS3 empirical model. It corresponds to a helium density of 8×10⁶ cm⁻³ at the level of where the CO₂ is 2×10¹⁰ cm⁻³.     

Spatially Resolved Far Ultraviolet Spectroscopy of the Jovian Aurora

Spatially resolved spectra in four 50-Å FUV spectral windows were obtained across the jovian aurora with the Space Telescope Imaging Spectrograph (STIS) on board the Hubble Space Telescope. Nearly simultaneous ultraviolet imaging makes it possible to correlate the intensity variations along the STIS slit with those observed in the images and to characterize the global auroral context prevailing at the time of the observations. Spectra at ∼1-Å resolution taken in pairs included an unabsorbed window and a spectral region affected by hydrocarbon absorption. Both sets of spectra correspond to an aurora with a main oval brightness of about 130 kilorayleighs of H₂ emission. The far ultraviolet color ratios I(1550-1620 Å)/I(1230-1300 Å) are 2.3 and 5.9 for the noon and morning sectors of the main oval, respectively. We use an interactive model coupling the energy degradation of incoming energetic electrons, auroral temperature and composition, and synthetic H₂ spectra to fit the intensity distribution of the H₂ lines. It is found that the model best fitting globally the spectra has a soft energy component in addition to a 10 erg cm⁻² s⁻¹ flux of 80 keV electrons. It provides an effective H₂ temperature of 540 K. The relative intensity of temperature-sensitive H₂ lines indicates differences between the auroral main oval and polar cap emissions. The amount of methane absorption across the polar region is shown to vary in a way consistent with temperature. For the second spectral pair, the polar cap shows a higher attenuation by CH₄, indicating a harder precipitation along high-latitude magnetic field lines.     

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.     

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.     

The Violet Opacity in the Red Peculiar Stars – III. Spectral analysis of the SC star GP Ori

Medium-resolution (1.5-2.5 Å) spectrum has been obtained for SC star GP Ori in a wide range of wavelengths from 3730 to 6250 Å. It is characterised mainly by strong atomic lines, as well as moderately strong bands of the CN violet system and CH (0,0) and (0,1) bands at 4315 and 4890 Å. Weak bands of C₂, C₃, and ZrO molecules are observed, however, no evidence is found for the presence of either SiC₂ (Merrill-Sanford bands) or YO. The most prominent atomic lines along with the NaID12 are those of CaII (K,H), CaI at 4227 Å, SrI at 4607 Å, SrII at 4077 and 4215 Å, and BaII at 4554 Å. H seems to be in emission.     

Radiative forcing of the Venus mesosphere: I. Solar fluxes and heating rates

Most of the solar energy absorbed by Venus is deposited in the atmosphere, at levels more than 60 km above the surface. This unusual flux distribution should have important consequences for the thermal structure and dynamical state of that atmosphere. Because there are few measurements of the solar flux at levels above 60 km, a radiative transfer model was used to derive the structure and amplitude of the solar fluxes and heating rates in the Venus mesosphere (60–100 km). This model accounts for all sources of extinction known to be important there, including absorption and scattering by CO₂, H₂O, SO₂, H₂SO₄ aerosols and an unidentified UV absorber. The distributions of these substances in our model atmosphere were constrained by a broad range of spacecraft and ground-based observations. Above the cloud tops, (71 km), near-infrared CO₂ bands absorb enough sunlight to produce globally averaged heating rates ranging from 4° K/day (24-hr period) at 71 km to more than 50° K/day at 100 km. The sulfuric acid aerosols that compose the Venus clouds are primarily scattering agents at solar wavelengths. These aerosols reflect about 75% of the incident solar flux before it can be absorbed by the atmosphere or surface. The unknown substance that causes the observed cloud-top ultraviolet contrasts is responsible for most of the absorption of sunlight within the upper cloud deck (57.5−71 km). This substance absorbs almost half of the sunlight deposited on Venus and contributes to solar heating rates as large as 6° K/day at levels near 65 km. With the exception of CO₂, all of the important sources of solar extinction have concentrations that vary with position, and, in general, these concentrations are not well known. To determine the sensitivity of the model results to these uncertainties, the concentrations of these opacity sources were varied in the model atmosphere and solar fluxes were computed for each case. These tests indicate that CO₂ dominates the solar absorption at levels above the cloud tops and that heating rates are relatively insensitive to the distribution of other sources of extinction there. Within the upper cloud deck, uncertainties in the distribution of the UV absorber and the H₂SO₄ aerosols can produce heating rate errors as large as 50% at some levels. Diurnally averaged solar heating rates for the nominal opacity distribution were computed as a function of latitude at altitudes between 55 and 100 km, where most of the solar flux is deposited. The zonal wavenumber 1 (diurnal) and zonal wavenumber 2 (semidiurnal) components of the diurnally varying solar heating rates were also computed in this domain. These results should be sufficiently reliable for use in numerical dynamical models of the Venus atmosphere.     

Ultra-violet nightglow spectrum from 1900 Å to 3400 Å

Two spectra of the ultra-violet nightglow from 1900 Å to 3400 Å have been recorded by a fast wide-field spectrograph during balloon flights from Aire sur Adour, France on 15 September, 1969 and on 5 October, 1970. These two spectra are composed of theOi line at 2972 Å, of the molecular oxygen Herzberg band systemA ³ _u ⁺ X ³ _g ^– and of a molecular band system that seems to belong to the NO SystemX ² C ² X ⁺, situated at 1991 ű4 Å, 2060 ű4 Å and 2136 ű4 Å. Around 2540 Å, there is absorption by ozone at the altitudes at which the spectra were recorded (35 km and 40 km). We present our calculated value of ozone absorption at 35 km, and the zenith-horizon variation of the nightglow emission.     

Evidence of N2-ice on the surface of the icy dwarf Planet 136472 (2005 FY9)

We present high signal precision optical reflectance spectra of 2005 FY9 taken with the Red Channel Spectrograph and the 6.5-m MMT telescope on 2006 March 4 UT (5000-9500 Å; 6.33 Å pixel⁻¹) and 2007 February 12 UT (6600-8500 Å; 1.93 Å pixel⁻¹). From cross-correlation experiments between the 2006 March 4 spectrum and a pure CH₄-ice Hapke model, we find the CH₄-ice bands in the MMT spectrum are blueshifted by 3 ± 4 Å relative to bands in the pure CH₄-ice Hapke spectrum. The higher resolution MMT spectrum of 2007 February 12 UT enabled us to measure shifts of individual CH₄-ice bands. We find the 7296, 7862, and 7993 Å CH₄-ice bands are blueshifted by 4 ± 2, 4 ± 4, and 6 ± 5 Å. From four measurements we report here and one of our previously published measurements, we find the CH₄-ice bands are shifted by 4 ± 1 Å. This small shift is important because it suggest the presence of another ice component on the surface of 2005 FY9. Laboratory experiments show that CH₄-ice bands in spectra of CH₄ mixed with other ices are blueshifted relative to bands in spectra of pure CH₄-ice. A likely candidate for the other component is N₂-ice because its weak 2.15 μm band and blueshifted CH₄ bands are seen in spectra of Triton and Pluto. Assuming the shift is due to the presence of N₂, spectra taken on two consecutive nights show no difference in CH₄/N₂. In addition, we find no measurable difference in CH₄/N₂ at different depths into the surface of 2005 FY9.     

A layer of ozone detected in the nightside upper atmosphere of Venus

To date, ozone has only been identified in the atmospheres of Earth and Mars. This study reports the first detection of ozone in the atmosphere of Venus by the SPICAV ultraviolet instrument onboard the Venus Express spacecraft. Venusian ozone is characterized by a vertically confined and horizontally variable layer residing in the thermosphere at a mean altitude of 100 km, with local concentrations of the order of 10⁷–10⁸ molecules cm⁻³. The observed ozone concentrations are consistent with values expected for a chlorine-catalyzed destruction scheme, indicating that the key chemical reactions operating in Earth’s upper stratosphere may also operate on Venus.     

A comparison of water vapor line parameters for modeling the Venus deep atmosphere

The discovery of the near infrared windows into the Venus deep atmosphere has enabled the use of remote sensing techniques to study the composition of the Venus atmosphere below the clouds. In particular, water vapor absorption lines can be observed in a number of the near-infrared windows allowing measurement of the H₂O abundance at several different levels in the lower atmosphere. Accurate determination of the abundance requires a good database of spectral line parameters for the H₂O absorption lines at the high temperatures (up to ∼700 K) encountered in the Venus deep atmosphere. This paper presents a comparison of a number of H₂O line lists that have been, or that could potentially be used, to analyze Venus deep atmosphere water abundances and shows that there are substantial discrepancies between them. For example, the early high-temperature list used by Meadows and Crisp [Meadows, V.S., Crisp, D., 1996. J. Geophys. Res. 101 (E2), 4595-4622] had large systematic errors in line intensities. When these are corrected for using the more recent high-temperature BT2 list of Barber et al. [Barber, R.J., Tennyson, J., Harris, G.J., Tolchenov, R.N., 2006. Mon. Not. R. Astron. Soc. 368, 1087-1094] their value of 45±10 ppm for the water vapor mixing ratio reduces to 27±6 ppm. The HITRAN and GEISA lists used for most other studies of Venus are deficient in “hot” lines that become important in the Venus deep atmosphere and also show evidence of systematic errors in line intensities, particularly for the 8000 to 9500 cm⁻¹ region that includes the 1.18 μm window. Water vapor mixing ratios derived from these lists may also be somewhat overestimated. The BT2 line list is recommended as being the most complete and accurate current representation of the H₂O spectrum at Venus temperatures.     

Characteristics of Saturn's polar atmosphere and auroral electrons derived from HST/STIS, FUSE and Cassini/UVIS spectra

Ultraviolet (UV) spectra of Saturn's aurora obtained with the Hubble Space Telescope Imaging Spectrograph (STIS), the Cassini Ultraviolet Imaging Spectrograph (UVIS) and the Far Ultraviolet Spectroscopic Explorer (FUSE) have been analyzed. Comparisons between the observed spectra and synthetic models of electron-excited H₂ have been used to determine various auroral characteristics. Far ultraviolet (FUV: 1200-1700 Å) STIS and UVIS spectra exhibit, below 1400 Å, weak absorption due to methane, with a vertical column ranging between 1.4×¹⁰¹⁵ and . Using the low-latitude Moses et al. [Moses, J.I., Bézard, B., Lellouch, E., Feuchtgruber, H., Gladstone, G.R., Allen, M., 2000. Icarus, 143, 244-298] atmospheric model of Saturn and an electron energy-H₂ column relationship, these methane columns are converted into the mean energy of the primary precipitating electrons, estimated to lie in the range 10-18 keV. This result is confirmed by the study of self-absorption with UVIS and FUSE extreme ultraviolet (EUV: 900-1200 Å) spectra. Below 1200 Å, it is seen that transitions connecting to the v^″<2 vibrational levels of the H₂ electronic ground state are partially self-absorbed by H₂ molecules overlying the auroral emission. Because of its low spectral resolution (∼5.5 Å), the UVIS EUV spectrum we analyzed does not allow us to unequivocally determine reasonable ranges of temperatures and H₂ columns. On the other hand, the high spectral resolution (∼0.2 Å) of the FUSE LiF1a and LiF2a EUV spectra we examined resolve the H₂ rotational lines and makes it possible to determine the H₂ temperature. The modeled spectrum best fitting the FUSE LiF1a observation reveals a temperature of 500 K and self-absorption by a H₂ vertical column of . When converted to energy of precipitating electrons, this H₂ column corresponds to primary electrons of ∼10 keV. The model that best fits the LiF2a spectrum is characterized by a temperature of 400 K and is not self-absorbed, making this segment ideal to determine the H₂ temperature at the altitude of the auroral emission. The latter value is in agreement with temperatures obtained from infrared polar spectra. Self-absorption is detectable in the LiF2a segment for H₂ columns exceeding , which sets the maximum mean energy determined from the FUSE observations to ∼15 keV. The total electron energy range of 10-18 keV deduced from FUV and EUV observations places the auroral emission peak between the 0.1 and 0.3 μbar pressure levels. These values should be seen as an upper limit, since most of the Voyager UVS spectra of Saturn's aurora examined by Sandel et al. [Sandel, B.R., Shemansky, D.E., Broadfoot, A.L., Holberg, J.B., Smith, G.R., 1982. Science 215, 548] do not exhibit methane absorption. The auroral H₂ emission is thus likely located above but close to the methane homopause. The H₂ auroral brightness in the 800-1700 Å bandwidth varies from 2.9 kR to 139 kR, comparable to values derived from FUV Faint Object Camera (FOC) and STIS images.     

An investigation of the SO2 content of the venusian mesosphere using SPICAV-UV in nadir mode

Using the SPICAV-UV spectrometer aboard Venus Express in nadir mode, we were able to derive spectral radiance factors in the middle atmosphere of Venus in the 170-320 nm range at a spectral resolution of R ? 200 during 2006 and 2007 in the northern hemisphere. By comparison with a radiative transfer model of the upper atmosphere of Venus, we could derive column abundance above the visible cloud top for SO₂ using its spectral absorption bands near 280 and 220 nm. SO₂ column densities show large temporal and spatial variations on a horizontal scale of a few hundred kilometers. Typical SO₂ column densities at low latitudes (up to 50°N) were found between 5 and 50 μm-atm, whereas in the northern polar region SO₂ content was usually below 5 μm-atm. The observed latitudinal variations follow closely the cloud top altitude derived by SPICAV-IR and are thought to be of dynamical origin. Also, a sudden increase of SO₂ column density in the whole northern hemisphere has been observed in early 2007, possibly related to a convective episode advecting some deep SO₂ into the upper atmosphere.     

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.     

Vertical profiling of SO2 and SO above Venus’ clouds by SPICAV/SOIR solar occultations

New measurements of sulfur dioxide (SO₂) and monoxide (SO) in the atmosphere of Venus by SPICAV/SOIR instrument onboard Venus Express orbiter provide ample statistics to study the behavior of these gases above Venus’ clouds. The instrument (a set of three spectrometers) is capable to sound atmospheric structure above the clouds in several observation modes (nadir, solar and stellar occultations) either in the UV or in the near IR spectral ranges. We present the results from solar occultations in the absorption ranges of SO₂ (190–230 nm, and at 4 μm) and SO (190–230 nm). The dioxide was detected by the SOIR spectrometer at the altitudes of 65–80 km in the IR and by the SPICAV spectrometer at 85–105 km in the UV. The monoxide’s absorption was measured only by SPICAV at 85–105 km. We analyzed 39 sessions of solar occultation, where boresights of both spectrometers are oriented identically, to provide complete vertical profiling of SO₂ of the Venus’ mesosphere (65–105 km). Here we report the first firm detection and measurements of two SO₂ layers. In the lower layer SO₂ mixing ratio is within 0.02–0.5 ppmv. The upper layer, also conceivable from microwave measurements by Sandor et al. (Sandor, B.J., Todd Clancy, R., Moriarty-Schieven, G., Mills, F.P. [2010]. Icarus 208, 49–60) is characterized by SO₂ increasing with the altitude from 0.05 to 2 ppmv, and the [SO₂]/[SO] ratio varying from 1 to 5. The presence of the high-altitude SO_x species could be explained by H₂SO₄ photodissociation under somewhat warmer temperature conditions in Venus mesosphere. At 90–100 km the content of the sulfur dioxide correlates with temperature increasing from 0.1 ppmv at 165–170 K to 0.5–1 ppmv at 190–192 K. It supports the hypothesis of SO₂ production by the evaporation of H₂SO₄ from droplets and its subsequent photolysis at around 100 km.     

Microwave Remote Sensing of the Temperature and Distribution of Sulfur Compounds in the Lower Atmosphere of Venus

A multi-wavelength radio frequency observation of Venus was performed on April 5, 1996, with the Very Large Array to investigate potential variations in the vertical and horizontal distribution of temperature and the sulfur compounds sulfur dioxide (SO₂) and sulfuric acid vapor (H₂SO₄(g)) in the atmosphere of the planet. Brightness temperature maps were produced which feature significantly darkened polar regions compared to the brighter low-latitude regions at both observed frequencies. This is the first time such polar features have been seen unambiguously in radio wavelength observations of Venus. The limb-darkening displayed in the maps helps to constrain the vertical profile of H₂SO₄(g), temperature, and to some degree SO₂. The maps were interpreted by applying a retrieval algorithm to produce vertical profiles of temperature and abundance of H₂SO₄(g) given an assumed sub-cloud abundance of SO₂. The results indicate a substantially higher abundance of H₂SO₄(g) at high latitudes (above 45°) than in the low-latitude regions. The retrieved temperature profiles are up to 25 K warmer than the profile obtained by the Pioneer Venus sounder probe at altitudes below 40 km (depending on location and assumed SO₂ abundance). For 150 ppm of SO₂, it is more consistent with the temperature profile obtained by Mariner 5, extrapolated to the surface via a dry adiabat. The profiles obtained for H₂SO₄(g) at high latitudes are consistent with those derived from the Magellan radio occultation experiments, peaking at around 8 ppm at an altitude of 46 km and decaying rapidly away from that altitude. At low latitudes, no significant H₂SO₄(g) is observed, regardless of the assumed SO₂ content. This is well below that measured by Mariner 10 (Lipa and Tyler 1979, Icarus39, 192-208), which peaked at ∼14 ppm near 47 km. Our results favor ≤100 ppm of SO₂ at low latitudes and ≤50 ppm in polar regions. The low-latitude value is statistically consistent with the results of Bézard et al. (1983, Geophs. Res. Lett.20, 1587-1590), who found that a sub-cloud SO₂ abundance of 130±40 ppm best matched their observations in the near-IR. The retrieved temperature profile and higher abundance of H₂SO₄(g) in polar regions are consistent with a strong equatorial-to-polar, cloud-level flow due to a Hadley cell in the atmosphere of Venus.     

The identification of new forbidden coronal lines in the solar EUV spectrum

Identifications are proposed for twenty of the twenty-eight coronal lines observed in the spectra obtained during a rocket flight into the path of the 7 March, 1970 solar eclipse. The methods by which the lines have been identified are discussed. Most of the lines identified are from forbidden transitions between levels in the ground 2p ⁿ and 3p ⁿ configurations in high ions of magnesium, silicon, sulphur, iron, and nickel. The temperature range represented is from 6.9 × 10⁵ K to 2.5 × 10⁶ K. The classification of three lines of Fexii and two of Nixiv has led to a revised identification for the near ultraviolet ² D _3/2-² P _1/2 transition in Fe xii. This transition can be identified with the line at 3072 Å rather than that at 3021 Å as previously suggested in the literature.     

Characteristics of Saturn’s FUV airglow from limb-viewing spectra obtained with Cassini-UVIS

This study reports the analysis of far ultraviolet (FUV) limb spectra of the airglow of Saturn in the 1150-1850 Å spectral window, obtained with the Ultraviolet Imaging Spectrograph (UVIS) onboard Cassini, spanning altitudes from −1200 to 4000 km. The FUV limb emission consists of three main contributions: (1) H Ly-α peaking at 1100 km with a brightness of 0.8 kilo-Rayleighs (kR), (2) reflected sunlight longward of 1550 Å which maximizes at −950 km with 16.5 kR and (3) H₂ bands in the 1150-1650 Å bandwidth, peaking at 1050 km reaching a maximum of 3.9 kR.A vertical profile of the local H₂ volume emission rate has been derived using the hydrocarbon density profiles from a model of the Saturn equatorial atmosphere. It is well matched by a Chapman function, characterized by a maximum value of 3.5 photons cm⁻³ s⁻¹ in the 800-1650 Å UV bandwidth, peaking at 1020 km.Comparisons between the observed spectra and a first-order synthetic airglow H₂ model in the 1150-1650 Å bandwidth show that the spectral shape of the H₂ bands is accounted for by solar fluorescence and photoelectron excitation. The best fits are obtained with a combination of H₂ fluorescence lines and 20 eV electron impact spectra, the latter contributing ∼68% of the total H₂ airglow emission.     

Observations of Venus water vapor over the disk of Venus: The 1972–1974 data using the H2O lines at 8197 Å and 8176 Å

The Venus water vapor line at 8197.71 Å has been monitored at several positions on the disk of Venus and at phase angles between 21° and 162°. Variations in the abundance have been found with spatial location, phase angle and time. During the 1972–1974 period, the total two-way absorption has varied from less than 1 to 77 μm of water vapor. The dependence on phase angle indicates 20 to 50 μm over the disk between 30° and 110° and small, but detectable amounts present during the rest of the observations. The spatial distribution with respect to the intensity equator is uniform with no location on the disk having systematically a higher or lower abundance. Comparisons made between the water vapor abundandances and the CO₂ abundances determined from near-simultaneous observations of CO₂ bands at the same positions on the disk of Venus show no correlation for the majority of the samples.     

The thermal structure of the dayside upper atmosphere of venus above 125 km

A one-dimensional model of the Venus thermosphere has been constructed which includes computation of the heating efficiency of solar ultraviolet radiation, heat loss by radiation to space of infrared-active species, thermal transport by molecular and eddy conduction, and viscous dissipation. By comparing model predictions with results obtained from the Pioneer Venus Orbiter space-craft, the results indicate that energy transport parameterized by eddy heat conduction plays a dominant role in determining thermospheric temperature T. It is suggested that there exists a feedback mechanism linking heating and thermospheric circulation such that eddy cooling maintains an asymptotic temperature T_∞～300°K for both solar-maximum and solar-minimum conditions. We also study the variation in thermospheric temperature with solar zenith angle, atomic oxygen-mixing ratio, rate of vibrational excitation of CO₂ by ground-state O atoms, and the assumed transfer of O(¹D) electronic energy to CO₂ vibrational energy.