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.     

An analysis of bending waves in Saturn's rings using voyager radio occultation data

The oblique geometry of the Voyager 1 radio occulation of Saturn's rings resulted in a strong coupling between the local slope of the ring midplane and the associated radio opacity (optical depth). We apply a model of this relationship to those regions of the rings where bending waves have been observed in the radio data. Using the Shu et al. linear model for a bending wave (F.H. Shu, J.N. Cuzzi, and J.J. Lissauer, 1983,Icarus53, 185–206), we obtain height profiles for the Mimas 5:3 and 7:4 bending waves. The first oscillation of the Mimas 5:3 bending wave has an amplitude of about 800 m, in agreement with the prediction of the Shu et al. model. However, the rest of the wave may be explained only by either a greatly decreased amplitude in the region beyond the second cycle, or by a significant enhancement in radio optical depth in the region of the bending wave. The shape of the enhancement necessary is similar to that of the enhancement at photopolarimetry wavelengths (L.W. Esposito, M. O'Callaghan, and R.A. West, 1983,Icarus56, 439–452), but differs in the region of the first cycle. Our solution gives 131,901±6 km as the resonance location, and a surface mass density of 35±6g cm⁻². The error bars on the resonance location do not include the uncertainty in the radial scale of the radio occultation data, which is approximately 10 km (R.A. Simpson, G.L. Tyler, and J.B. Holberg, 1983,Astron. J.88, 1531–1536). The Mimas 7:4 bending wave conforms more closely to the linear model, and requires no reduction in amplitude or enhancement in optical depth. We find a surface mass density of 30.5±9 g cm⁻², and resonance location at 127,765±7km.     

Statistical and computational uncertainties in atmospheric profiles from radio occultation: Mariner 10 at Venus

Radio occultation studies of planetary atmospheres and ionospheres are based on measurements of the frequency and amplitude of the received radio signal. These measurements have random errors due to noise in the receiving system and linearly mapped into atmospheric profiles to give uncertainties can be estimated from the data and linearly mapped into atmospheric profiles to give uncertainties in temperature, T, pressure, p, and absorption profiles. For Mariner 10 occultation immersion at Venus, the standard deviations of T and p due to receiver noise are less than 2° K and 2 mbar over the range of radii from 6087 to 6140 km, based on our reduction from analog, “ open-loop” data. The temperature has a systematic error due to boundary uncertainty, estimated to be 50°K at 6140 km, that decays rapidly with depth; below 6117 km, it is less than 0.5°K. For the attenuation profile, systematic errors incurred during our calculations are more important than statistical errors. We estimate an upper bound to the uncertainty which is 32% at the peak value of absorption, which is about 0.01 db/km and occurs at a radius of 6096 km. A calculation of the 95% confidence limits for T profiles indicates that the local deviations are statistically significant to about 1°K or less. We have also analyzed “closed-loop” data to give temperature profiles which deviate from the open-loop results by less than 0.2°K below 6110 km but by as much as 2°K in the upper atmosphere. For the same occultation and the same boundary conditions, our closed-loop T-p profile is within 2°K of that of P. D. Nicholson and D. O. Muhleman but differs from those derived by A. J. Kliore by as much as 10°K. We cannot account for deviations as large as the latter by minor differences in trajectory information or computational methods.     

Turbulence in planetary occultations.: IV. Power spectra of phase and intensity fluctuations

Power spectra of phase and intensity scintillations during occultation by turbulent planetary atmospheres are significantly affected by the inhomogeneous background upon which the turbulence is superimposed. Such coupling is particularly pronounced in the intensity, where there is also a marked difference in spectral shape between a central and a grazing occultation. While the former has its structural features smoothed by coupling to the inhomogeneous background, such features are enhanced in the latter. Indeed, the latter power spectrum peaks around the characteristic frequency that is determined by the size of the free-space Fresnel zone and the ray velocity in the atmosphere; at higher frequencies strong fringes develop in the power spectrum. A confrontation between the theoretical scintillation spectra computed here and those calculated from the Mariner 5 Venus mission by R. Woo, A. Ishimaru, and W. B. Kendall (1974, J. Atmos. Sci.31, 1698–1706) is inconclusive, mainly because of insufficient statistical resolution. Phase and/or intensity power spectra computed from occultation data may be used to deduce characteristics of the turbulence and to distinguish turbulence from other perturbations in the refractive index. Such determinations are facilitated if observations are made at two or more frequencies (radio occultation) or in two or more colors (stellar occultation).     

Infrared limb darkening of the Venus atmosphere

An analysis of the limb darkening component obtained by Ingersoll and Orton [Icarus21 (1974), 121–128] from the thermal infrared maps of Venus published by Murray, Wildey, and Westphal [J. Geophys. Res.68 (1963), 4813–4818] and Westphal, Wildey, and Murray [Astrophys. J.142 (1965), 799–802] shows that the Cytherean cloud tops were close to radiative equilibrium in 1962. A method for obtaining the optical depth, the extinction coefficient, and the extinction scale height from such data is derived and values are extracted from Marov's [Icarus16 (1972), 415–461] standard model of the Venus atmosphere.     

An electrodynamic model of the solar wind interaction with the ionospheres of Mars and Venus

The electrodynamic model for the solar wind interaction with non-magnetic planets. (Cloutier and Daniell, Planet. Space Sci.21, 463, 1973; Daniell and Cloutier, Planet. Space Sci.25, 621, 1977) is modified to include the effects of non-ohmic currents in the upper ionosphere. The model is then used to calculate convection patterns induced by the solar wind in the ionospheres of Mars and Venus. For Mars the observations of the neutral mass spectrometer or Vikings 1 and 2 provided the neutral atmosphere. Model calculations reproduced the retarding potential analyzer data and indicate that the ionosphere above about 200 km is probably controlled by convection rather than chemistry or diffusion. For Venus a model atmosphere based on Dickenson and Ridley, J. Atmos. Sci.32, 1219 (1975) and Mayr et al., J. geophys. Res.83, 4411 (1978) was used. The resulting model calculations were compared to radio occultation data from Mariners 5 and 10 and Venera 9 which represent extremes in the variability of the upper Cytherean ionosphere. The model calculations are shown to fall within this variation. These results represent the state of the theory immediately prior to the Pioneer-Venus encounter.     

Independent radio-occultation studies of venus' atmosphere

Two independent analyses of the dual-frequency radio-occultation experiment performed by Mariner 10 at Venus are presented. Using closed-loop frequency data obtained at NASA's Goldstone facility, we have computed S- and X-band pressure-temperature profiles for Venus' neutral atmosphere, and an S-band profile of the nightside ionosphere. Neutral atmosphere dispersion between the two frequencies is negligible (less than 0.1% in refractivity), as expected for a CO₂ atmosphere. The results confirm those obtained by Howard et al. (1974) from the same S-band data with an accuracy of ±5°K at a given pressure level, though there is a discrepancy of 1 km in the radial scale between the two analyses. These two Mariner 10 profiles are compared with the Mariner 5 occultation profile and in situ measurements by Veneras 8, 9, and 10. The occultation was also monitored at the Owens Valley Radio Observatory, though only at X-band. Despite the much lower quality of these data, a reasonable neutral atmosphere refractivity profile above 65 km was obtained from the occultation entry. Uncertainties in the calculated temperatures, however, are too large to permit useful comparison with previous results. The existence of real anomalies in both the amplitude and frequency of the signal during exit from occultation is confirmed.     

Turbulence in planetary occultations: II. Effects on atmospheric profiles derived from Doppler measurements

Turbulence in planetary atmospheres leads to both fluctuating and systematic errors in atmospheric profiles derived from Doppler measurements during radio occultation. If the upper atmospheres of Venus and Jupiter are about as turbulent as the earth's troposphere, we deduce rms fractional errors in temperature and pressure of less than ～ 10^?2 for the Mariner 10 and Pioneer 10/11 occultation experiments. Fractional systematic errors are typically of the order of 10^?6. These estimates depende rather weakly on quantities characterizing the atmosphere and the occultation, and it is conjectured that turbulence-induced errors in atmospheric profiles derived from Doppler measurements are always very small in the weak scattering limit     

A comment on the insolation at Mercury

E. Van Hemelrijck and J. Vercheval [Icarus48, 167–179 (1981)] presented calculations of the insolation at Mercury and Venus which neglect the finite angular size of the Sun. To determine the temperature structure in the subsurface a more accurate calculation is needed, especially at longitudes ±90° on Mercury, where the Sun takes 18 days to rise or set. These calculations are presented here.     

Stability of Pluto's atmosphere

L. Trafton (1980, Icarus44, 53–61) has pointed out that a substantial methane atmosphere, observed on Pluto by U. Fink, B.A. Smith, D.C. Benner, J.R. Johnson, and H.J. Reitsema (1980, Icarus44, 62–71), appears to be unstable against blowoff. The difficulty is shown to disappear if the actual heat balance and thermal structure are considered, instead of the classic assumption that the upper atmosphere is isothermal. An energy-limited flux (referred to the surface area) of 3.9 × 10¹⁰ cm^?2 sec^?1 is found. The loss of methane ice over the age of the solar system is an acceptable 3 km.     

Light scattering in a spherical,exponential atmosphere,with applications to Venus

The problem of the reflection of light from an optically thick, spherical atmosphere in which the scatterers are distributed exponentially with a scale height small compared to the radius of the planet is discussed. Exact formal solutions are obtained for the single scattered component. Useful approximate analytic solutions, which also include multiply scattered light, are given. The results are applied to the analysis of the Mariner 10 limb and terminator images of Venus. The altitude of the “detached” haze layer discovered by Mariner 10 is at 79–85 km, but in places the haze exists above 100 km. This layer apparently is a stable, planetwide feature which forms at the top of the Pioneer Venus upper haze layer. It was similar in location, scale height, and thickness at the times of the two missions, in contrast to the lower, high-altitude haze which changed dramatically. We discuss two possibilities for the nature of the limb hazes. (1) The lower haze is probably the sulfuric acid cloud and the “detached” layer may be a separate water-ice haze. (2)The “detached” haze layer may not be separate at all, but part of the sulfuric acid haze, and the apparent “gap” at 75–80 km may be the source region of a broadband absorber. The spatial distribution of the strong near-UV absorber, which may be elemental sulfur as first suggested by B. Hapke and R. Nelson (1975, J. Atmos. Sci.32, 1212–1218), is examined in light of our results. Several arguments indicate that there is no nonabsorbing, overlying haze and that the UV absorber extends to the top of the haze 8layer.     

Liquid content of the lower clouds of venus as determined from mariner 10 radio occultation

S-band (13.06-cm) and X-band (3.56-cm) radio occulation data obtained during the flyby of Venus by Mariner 10 on February 5, 1974 were analyzed to obtain the effects of dispersive microwave absorption by the clouds of Venus. The received power profiles were first corrected for the effects of refraction in the atmosphere of Venus, programmed changes in the pointing direction of the high-gain antenna, and limit-cycle motion of the spacecraft attitude control system. The resulting excess attenuation profiles presumbaly due to cloud absorption have been inverted discretely to obtain profiles of absorption coefficient at the two wavelenghts. The ratios of the absorptivities are consistent with a sulfuric acid-water mixture as the constituent of the absorbing clouds, having a sulfuric acid concentration of 75 ± 25%. Three absorption peaks are evident in the profiles at altitudes of 68, 60, and 48 km. With a sulfuric acid concentration of 75%, the upper cloud has a peak liquid content of 0.08 g/m³, and an integrated content of 0.024 g/cm², which corresponds roughly to terrestrial stratus or altostratus clouds. The major absorption layer has a peak of 1.1 g/m³ at an altitude of 48 km, with an integrated content of 0.5 g/cm², similar to that of terrestrial cumulus and cumulonimbus clouds. The absorption ratios for the middle cloud at 60 km are not consistent with a sulfuric acid-water mixture.     

Polarization studies of the Venus uv contrasts: Cloud height and haze variability

Polarimetry is able to show direct evidence for compositional differences in the Venus clouds. We present observations (collected during $\text{2}\text{1}\text{2}$ Venus years by the Pioneer Venus Orbiter) of the polarization in four colors of the bright and dark ultraviolet features. We find that the polarization is significantly different between the bright and dark areas. The data show that the “null” model of L. W. Esposito (1980, J. Geophys. Res.85, 8151–8157) and the “overlying haze” model of J. B. Pollack et al. (1980, J. Geophys. Res.85, 8223–8231) are insufficient. Exact calculations of the polarization, including multiple scattering and vertical inhomogeneity near the Venus cloud tops, are able to match the observations. Our results give a straightforward interpretation of the polarization differences in terms of known constituents of the Venus atmosphere. The submicron haze and uv absorbers are anticorrelated: for haze properties as given by K. Kawabata et al. (1980, J. Geophys. Res.85, 8129–8140) the excess haze depth at 9350 Å over the bright regions is Δτ_h = 0.03 ± 0.02. The cloud top is slightly lower in the dark features: the extra optical depth at 2700 Å in Rayleigh scattering above the darker areas is Δτ_R = 0.010 ± 0.005. This corresponds to a height difference of 1.2 ± 0.6 km at the cloud tops. The calculated polarization which matches our data also explains the relative polarization of bright and dark features observed by Mariner 10. The observed differential polarization cannot be explained by differential distribution of haze, if the haze aerosols have an effective size of 0.49 μm, as determined by K. Kawabata et al. (1982, submitted) for the aerosols overlying the Venus equator. We propose two models for the uv contrasts consistent with our results. In a physical model, the dark uv regions are locations of vertical convergence and horizontal divergence. In a chemical model, we propose that the photochemistry is limited by local variations in water vapor and molecular oxygen. The portions of the atmosphere where these constituents are depleted at the cloud tops are the dark uv features. Strong support for this chemical explanation is the observation that the number of sulfur atoms above the cloud tops is equal over both the bright and dark areas. The mass budget of sulfur at these altitudes is balanced between excess sulfuric acid haze over the bright regions and excess SO₂ in the dark regions.     

Theoretical interpretation of the 0.7820 μm CO2 band and 0.8226 μm H2O line on Venus

We have analyzed the P6, P8, and P10 lines in the 0.7820 μm CO₂ band of Venus using a scattering model. Our new results compare favorably with previous results from the 1.05 μm CO₂ band. We considered nonabsorbing and absorbing clouds. We found that the anisotropic scattering mean free path for both models at the 0.2atm level is between 0.55 and 0.73km, a range close to the value of 1 km for terrestrial hazes. We used our scattering models to synthesize the 0.8226 μm H₂O line, assuming that the clouds are composed of sulfuric acid drops, and found our nonabsorbing cloud required a sulfuric acid concentration of 82% by weight, while our thicker absorbing cloud required a concentration of 89%. A comparison of the variation of optical depth with height for our cloud models with the variation reported by Prinn (1973, Science182, 1132–1134) showed that, within a factor of 2, the variation for Prinn's thinnest cloud agreed with ours. Whitehill and Hansen (1973, Icarus20, 146–152) have recently confirmed the work of Regas et al. (1973a, J. Quant. Spectry. Radiative Transfer13, 461–463) which showed that two cloud layers are not required to explain the CO₂ phase variation of Venus. Prinn's recent photochemical study of sulfuric acid clouds further supports a single, continuous cloud layer in the line formation region instead of two cloud layers with an extensive clear region between. The single layer model appears more likely because the maximum particle density in Prinn's cloud occurs in the clear region between the two layers in the models of Hunt (1972, J. Quant. Spectry. Radiative Transfer12, 405–419) and Carleton and Traub (1972, Bull. Amer. Astron. Soc.4, 362.).     

Optical signature of Venus lightning as seen from space

The calculations of M. A. Williams, L. W. Thomason, and D. M. Hunten (Icarus52, 166–170, 1982) for the light transmitted to space by a Venus lightning flash by dropping the use of similarity relations have been improved. This revised model increases their escape fractions and image sizes by about a factor of 2; however, their conclusions remain valid.     

Joule heating of Io's ionosphere by unipolar induction currents

The unexpectedly large scale height of Io's ionosphere (Kliore, A., et al., 1975, Icarus24, 407–410) together with the relatively large molecular weight of the likely principal constituent, SO₂ (Pearl, J., et al., 1979, Nature280, 755–758), suggest a high ionospheric temperature. Electrical induction in Io's ionosphere due to the corotating plasma bound to the Jovian magnetosphere is one possible source for attainment of such high temperatures. Accordingly, unipolar induction models were constructed to calculate ionospheric joule heating numerically. Heating rates produced by highly simplified models lie in the range 10^?9 to 10^?8 W/m³. These heating rates are lower than those determined from uv photodissociative heating models (Kumar, S., 1980, Geophys. Res. Lett.7, 9–12) at low levels in the ionosphere but are comparable in the upper ionosphere. The low electrical heating rate throughout most of the ionosphere is due to the power limitation imposed by the Alfvén wings which complete the electrical circuit (Neubauer, F.M., 1980, J. Geophys. Res.85, 1171–1178). Contrary to the pre-Voyager calculations of Cloutier, P. A., et al. (1978, Astrophys. Space Sci.55, 93–112), our numerical results show that the J × B force density due to unipolar induction currents in the ionosphere is much less than the gravitational force density when the combined mass of the neutral species is included. The binding and coupling of the ionosphere is principally due to the relatively dense (possibly localized) neutral SO₂ atmosphere. In regions where the ions and neutrals are collisionally coupled the ionosphere will not be stripped off by the J × B forces. However at a level above that (to which the ions move by diffusion only) the charged species would be removed. Thus there appears to be no need to postulate the existence of an intrinsic Ionian magnetic field as suggested by Kivelson, M. G., et al. (79, Science 205, 491–493) and Southwood, S. J., et al. (1980, J. Geophys. Res., in press) in order to retain the observed ionosphere.     

Venus mesosphere and thermosphere: II. Global circulation,temperature, and density variations

Recent Pioneer Venus observations have prompted a return to comprehensive hydrodynamical modeling of the thermosphere of Venus. Our approach has been to reexamine the circulation and structure of the thermosphere using the framework of the R. E. Dickinson an E. C. Ridley (1977, Icarus 30, 163–178), symmetric two-dimensional model. Sensitivity tests were conducted to see how large-scale winds, eddy diffusion and conduction, and strong 15-μm cooling affect day-night contrasts of densities and temperatures. The calculated densities and temperatures are compared to symmetric empirical model fields constructed from the Pioneer Venus data base. We find that the observed day-to-night variation of composition and temperatures can be derived largely by a wave-drag parameterization that gives a circulation system weaker than predicted prior to Pioneer Venus. The calculated mesospheric winds are consistent with Earth-based observations near 115 km. Our studies also suggest that eddy diffusion is only a minor contributor to the maintenance of observed day and nightside densities, and that eddy coefficients are smaller than values used by previous one-dimensional composition models. The mixing that occurs in the Venus thermosphere results from small-scale and large-scale motions. Strong CO₂ 15-μm cooling buffers solar perturbation such that the response by the general circulation to solar cycle variation is relatively weak.     

Saturn's rings: 3-mm low-inclination observations and derived properties

Recent 3-mm observations of Saturn at low ring inclinations are combined with previous observations of E. E. Epstein, M. A. Janssen, J. N. Cuzzi, W. G. Fogarty, and J. Mottmann (Icarus41, 103–118) to determine a much more precise brightness temperature for Saturn's rings. Allowing for uncertainties in the optical depth and uniformity of the A and B rings and for ambiguities due to the C ring, but assuming the ring brightness to remain approximately constant with inclination, a mean brightness temperature for the A and B rings of 17 ± 4°K was determined. The portion of this brightness attributed to ring particle thermal emission is 11 ± 5°K. The disk temperature of Saturn without the rings would be 156 ± 6°K, relative to B. L. Ulich, J. H. Davis, P. J. Rhodes, and J. M. Hollis' (1980, IEEE Trans. Antennas Propag.AP-28, 367–376) absolutely calibrated disk temperature for Jupiter. Assuming that the ring particles are pure water ice, a simple slab emission model leads to an estimate of typical particle sizes of ≈0.3 m. A multiple-scattering model gives a ring particle effective isotropic single-scattering albedo of 0.85 ± 0.05. This albedo has been compared with theoretical Mie calculations of average albedo for various combinations of particle size distribution and refractive indices. If the maximum particle radius (≈5 m) deduced from Voyager bistatic radar observations (E. A. Marouf, G. L. Tyler, H. A. Zebker, V. R. Eshleman, 1983, Icarus54, 189–211) is correct, our results indicate either (a) a particle distribution between 1 cm and several meters radius of the form r^?s with 3.3 ? s ? 3.6, or (b) a material absorption coefficient between 3 and 10 times lower than that of pure water ice Ih at 85°K, or both. Merely decreasing the density of the ice Ih particles by increasing their porosity will not produce the observed particle albedo. The low ring brightness temperature allows an upper limit on the ring particle silicate content of ≈10% by mass if the rocky material is uniformly distributed; however, there could be considerably more silicate material if it is segregated from the icy material.     

Bending waves in Saturn's rings

We investigate certain brightness variations seen in Saturn's A ring and find them to be due to vertical corrugations of the local ring plane caused by a spiral bending wave. This wave is resonantly excited by Mimas and propagates inward via the collective gravity of the ring particles. B. A. Smith et al. [Science212, 163–191 (1981)] had previously associated vertical relief with this feature due to its observed azimuthal variations and its proximity to an inclination resonance with Mimas. We develop the theory of forced bending waves, some aspects of which have been treated in the galactic context by C. Hunter and A. Toomre [Astrophys. J.155, 747–776 (1969)] and by G. Bertin and J.W.-K. Mark [Astron. Astrophys.88, 289–297 (1980)]. Our theory is in good agreement with the observations. In particular, the presence of these bending waves may resolve the conflict between ground-based estimates of 1–2 km for the global ring thickness [e.g., A. Brahic and B. Sicardy, Nature289, 447–450 (1981)] and Voyager stellar occultation measurements of <200 m for the local ring thickness [A. L. Lane et al., Science215, 537–543 (1982); E. A. Marouf and G. L. Tyler, Science217, 243–245 (1982)].     

The composition and vertical structure of the lower cloud deck on Venus

The opportunity to determine the planetwide temperature and cloud structure of Venus using radio occultation techniques arose with Pioneer Venus. Amplitude and Doppler data provided by the radio occultation experiment offered a unique and powerful means of examining the atmospheric properties in the lower cloud region.Absorption due to gaseous components of the atmosphere was subtracted from the measured absorption coefficient profiles before they were used to compute cloud mass contents. This absorption was found to represent a small part of the total absorption, depending on the latitude. In the main cloud deck, gaseous absorption contributes 10 to 20%, however, at the bottom of the detected absorption layer the sulfuric acid vapor contributes up to 100% due to increased vapor pressures. The clouds are the primary contributing absorbers in the 1- to 3-bar level of the Venus atmosphere. Below about 3 bars, depending on the latitude, absorption due to sulfuric acid vapor dominates.If a cloud particle model consisting of a solid nonabsorbing dielectric sphere with a concentric liquid sulfuric acid coating is invoked, the absorptivity of the particles increases from that of a pure sulfuric acid liquid sphere, and the mass content derived from the absorption coefficient profiles decreases. As the ratio of the core radius to the total radius (q) increases, absorption increases by more than a factor of 10 for high values of q. In the case of pure sulfuric acid droplets, the conductivity is sufficiently high that some of the field is excluded from the interior of the droplet thereby reducing the absorption. When a dielectric core of nonabsorbing material is introduced, the surface charge density is reduced and the absorption increases.The mass contents for all orbits in the equatorial region of Venus were calculated using values of q from 0 to 1. The resulting profiles match the probe mass content profiles at similar locations when a q of 0.97 is chosen.The wavelength dependence of the absorption for the spherical shell model varies with q from 1/λ² for pure liquid to λ^0.2 for a large core. A q of from 0.96 to 0.98 results in a wavelength dependence of 1/λ^1.0 to 1/λ^1.4 which matches the radio occultation absorption wavelength dependence and the microwave opacity wavelength dependence.Mass content profiles using a q of 0.97 were determined for occultations in the polar, collar, midlatitudinal, and equatorial regions assuming q remains constant over the planet. The results show considerable variability in both the level and the magnitude of the lower cloud deck. The cloud layer is lowest in altitude in the polar region. This might be expected as the temperature profile is cooler in the polar region than over the rest of the planet. The mass content is greatest in the polar and collar regions; however, many of the collar profiles were cut off due to fluctuations resulting from increased turbulence in the collar region. The mass contents are least dense in the midlatitude regions. There is a sharp lower boundary at about 1.5 bars in the equatorial and midlatitude regions and at about 2.5 bars in the polar region. Measurements made by the Particle Size Spectrometer and nephelometers also showed sharp lower cloud boundaries at this level.