Martian isotopic ratios and upper limits for possible minor constituents as derived from Mariner 9 infrared spectrometer data

The Mariner 9 infrared spectrometer obtained data over a large part of Mars for almost a year beginning late in 1971. Mars' infrared emission spectrum was measured from 200 to 2000 cm^?1 with an apodized resolution of 2.4 cm^?1. No significant deviation from terrestrial ratios of carbon (¹²C/¹³C) or oxygen (¹⁶O/¹⁸O; ¹⁶O/¹⁷O) isotopes was observed on Mars. The ¹²C/¹³C isotopic ratio was found to be terrestrial with an uncertainty of 15%. Upper limits have been calculated for several minor constituents. With an effective noise equivalent radiance of 1.2 × 10^?9 W cm^?2 sr^?1/cm^?1, new upper limits in centimeter-atmospheres of 2 × 10^?5 for C₂H₂, 4 × 10^?3 for C₂H₄, 3 × 10^?3 for C₂H₆, 2 × 10^?4 for CH₄, 1 × 10^?3 for N₂O, 1 × 10^?4 for NO₂, 4 × 10^?5 for NH₃, 1 × 10^?3 for PH₃, 7 × 10^?4 for SO₂, and 1 × 10^?4 for OCS have been derived.     

For infrared spectrophotometry of Jupiter and Saturn

Infrared spectral observations of Mars, Jupiter, and Saturn were made from 100 to 470 cm^?1 using NASA's G. P. Kuiper Airborne Observatory. Taking Mars as a calibration source, we determined brightness temperatures of Jupiter and Saturn with approximately 5 cm^?1 resolution. The data are used to determine the internal luminosities of the giant planets, for which more than 75% of the thermally emitted power is estimated to be in the measured bandpass: for Jupiter L_J = (8.0 ± 2.0) × 10^?10L_⊙ and for Saturn L_S = (3.6 ± 0.9) × 10^?10. The ratio R of thermally emitted power to solar power absorbed was estimated to be R_J = 1.6 ± 0.2, and R_S = 2.7 ± 0.8 from the observations when both planets were near perihelion. The Jupiter spectrum clearly shows the presence of the rotational ammonia transitions which strongly influence the opacity at frequencies ?250 cm^?1. Comparison of the data with spectra predicted from current models of Jupiter and Saturn permits inferences regarding the structure of the planetary atmospheres below the temperature inversion. In particular, an opacity source in addition to gaseous hydrogen and ammonia, such as ammonia ice crystals as suggested by Orton, may be necessary to explain the observed Jupiter spectrum in the vicinity of 250 cm^?1.     

The mass,gravity field,and ephemeris of Mercury

This paper represents a final report on the gravity analysis of radio Doppler and range data generated by the Deep Space Network (DSN) with Mariner 10 during two of its encounters with Mercury in March 1974 and March 1975. A combined least-squares fit to Doppler data from both encounters has resulted in a determination of two second degree gravity harmonics, J₂ = (6.0 ± 2.0) × 10⁻⁵ and C₂₂ = (1.0 ± 0.5) × 10⁻⁵, referred to an equatorial radius of 2439 km, plus an indication of a gravity anomaly in the region of closest approach of Mariner 10 to Mercury in March 1975 amounting to a mass deficiency of about GM = −0.1 km³sec⁻². An analysis is included that defends the integrity of previously published values for the mass of Mercury (H. T. Howard et al. 1974, Science 185, 179–180; P. B. Esposito, J. D. Anderson, and A. T. Y. Ng 1978, COSPAR: Space Res. 17, 639–644). This is in response to a published suggestion by R. A. Lyttleton (1980, Q. J. R. Astron. Soc. 21, 400–413; 1981, Q. J. R. Astron. Soc. 22, 322–323) that the accepted values may be in error by more than 30%. We conclude that there is no basis for being suspicious of the earlier determinations and obtain a mass GM = 22,032.09 ± 0.91 km³sec⁻² or a Sun to Mercury mass ratio of 6,023,600 ± 250. The corresponding mean density of Mercury is 5.43 ± 0.01 g cm⁻³. The one-sigma error limits on the gravity results include an assessment of systematic error, including the possibility that harmonics other than J₂and C₂₂ are significantly different from zero. A discussion of the utility of the DSN radio range data obtained with Mariner 10 is included. These data are most applicable to the improvement of the ephemeris of Mercury, in particular the determination of the precession of the perihelion.     

Spectrophotometry of Saturn and its rings from 60 to 180 microns

We observed Saturn at far-infrared and submillimeter wavelengths during the Earth's March 1980 passage through the plane of Saturn's rings. Comparison with earlier spectroscopic observations by D. B. Ward [Icarus32, 437–442 (1977)], obtained at a time when the tilt angle of the rings was 21.8°, permits separation of the disk and ring contributions to the flux observed in this wavelength range. We present two main results: (1) The observed emission of the disk between 60 and 180 μm corresponds to a brightness temperature of 104 ± 2°K; (2) the brightness temperature of the rings drops approximately 20°K between 60 and 80 μm. Our data, in conjunction with the data obtained by other observers between 1 μm and 1 mm, permit us to derive an improved estimate for the total Saturnian surface brightness of (4.84 ± 0.32) × 10^?4W cm^?2 corresponding to an effective temperature of 96.1 ± 1.6°K. The ratio of radiated to incident power, P_R/P_I, is (1.46 ± 0.08)/(1 - A), where A is the Bond albedo. For A = 0.337 ± 0.029, P_R/P_I = 2.20 ± 0.15 and Saturn's intrinsic luminosity is L_S = (2.9 ± 0.5) × 10^?10L_⊙.     

Kinematics of the Gould Belt based on open clusters

We have redetermined the kinematic parameters of the Gould Belt using currently available data on the motion of nearby young (log t < 7.91) open clusters, OB associations, and moving stellar groups. Our modeling shows that the residual velocities reach their maximum values of ?4 km s^?1 for rotation (in the direction of Galactic rotation) and +4 km s^?1 for expansion at a distance from the kinematic center of ≈300 pc. We have taken the following parameters of the Gould Belt center: R ₀ = 150 pc and l ₀ = 128°. The whole structure is shown to move relative to the local standard of rest at a velocity of 10.7 ± 0.7 km s^?1 in the direction l = 274° ± 4° and b = ?1° ± 3°. Using the derived rotation velocity, we have estimated the virial mass of the Gould Belt to be 1.5 × 10⁶ M _⊙.     

Martian volatiles: Their degassing history and geochemical fate

Observations of Mars and cosmochemical considerations imply that the total inventory of degassed volatiles on Mars is 10² to 10³ times that present in Mars' atmosphere and polar caps. The degassed volatiles have been physically and chemically incorporated into a layer of unconsolidated surface rubble (a “megaregolith”) up to 2km thick. Tentative lines of evidence suggest a high concentration (～5g/cm²) of ⁴⁰ Ar in the atmosphere of Mars. If correct, this would be consistent with a degassing model for Mars in which the Martian “surface” volatile inventory is presumed identical to that of Earth but scaled to Mars' smaller mass and surface area. The implied inventory would be: (⁴⁰Ar) = 4g/cm², (H₂O) = 1 × 10⁵g/cm², (CO₂) = 7 × 10³g/cm², (N₂) = 3 × 10²g/cm², (Cl) = 2 × 10³g/cm², and (S) = 2 × 10²g/cm². Such a model is useful for testing, but differences in composition and planetary energy history may be anticipated between Mars and Earth on theoretical grounds. Also, the model demands huge regolith sinks for the volatiles listed.If the regolith were in physical equilibrium with the atmosphere, as much as 2 × 10⁴g/cm² of H₂O could be stored in it as hard-frozen permafrost, or 5 × 10⁴g/cm² if equilibrium with the atmosphere were inhibited. Spectral measurements of Martian regolith material and laboratory measurement of weathering kinetics on simulated regolith material suggest large amounts of hydrated iron oxides and clay minerals exist in the regolith; the amount of chemically bound H₂O could be from 1 × 10⁴ to 4 × 10⁴g/cm². In an Earth-analogous model, a 2 km mixed regolith must contain the following concentrations of other volatile-containing compounds by weight: carbonates = 1.5%, nitrates = 0·3%, chlorides = 0.6%, and sulfates = 0.1%. Such concentrations would be undetectable by current Earth-based spectral reflectance measurements, and (except the nitrates) formation of the “required” amounts of these compounds could result from interaction of adsorbed H₂O and ice with primary silicates expected on Mars. Most of the CO₂ could be physically adsorbed on the regolith.Thus, maximum amounts of H₂O and other volatiles which could be stored in the Mars regolith are marginally compatible with those required by an Earth-analogous model, although a lower atmospheric ⁴⁰Ar concentration and regolith volatile inventory would be easier to reconcile with observational constraints. Differences in the ratios of H₂O and other volatiles to ⁴⁰Ar between surface volatiles on the real Mars and on an Earth-analogous Mars could result from and reflect differences in bulk composition and time history of degassing between Mars and Earth. Models relating Viking-observable parameters, e.g., (⁴⁰Ar) and (³⁶Ar), to the time history and overall intensity of Mars degassing are given.     

Diffusion-deposition patterns in Martian streaks

The spreading angle of a number of light and dark Martian streaks is determined from selected Mariner 9 images. The resulting frequency distributions of spreading half-angles have maxima at ～5° for light, and ～7° for dark streaks; however the dark streaks have a secondary maximum spreading angle at ～14°. The smaller values, which include most streaks, are interpreted as crater-wake spreading phenomena. The larger value, found in only a few dark streaks or “tails,” may result from atmospheric diffusion and subsequent deposition of material from isolated sources such as vents or blowouts. An atmospheric diffusion-deposition analysis is presented, assuming this streak origin, from which it is possible to deduce the eddy diffusivity, K, in Mars' boudary layer. Calculated K values are found to agree with various theoretical estimates. They lie in the range 10⁷ and 10⁹ cm² sec^?1 and exhibit the proper scale dependence. Thus it appears that, in addition to streak-derived wind direction patterns and speed information, it is possible in a few cases to derive information on Mars' boundary-layer turbulence from streak-spreading measurements.     

On the structure of Mars' atmosphere at 120–220 km

Altitude dependences of [CO₂] and [CO₂⁺] are deduced from Mariner 6 and 7 CO₂⁺ airglow measurements. CO₂ densities are also obtained from n_e radio occultation measurements. Both [CO₂] profiles are similar and correspond to the model atmosphere of Barth et al. (1972) at 120 km, but at higher altitudes they diverge and at 200–220 km the obtained [CO₂] values are three times less the model. Both the airglow and radio occultation observations show that a correction factor of 2.5 should be included into the values for solar ionization flux given by Hinteregger (1970). The ratio of [CO₂⁺]/n_e is 0.15–0.2 and, hence, [O]/[CO₂] is ～3% at 135 km. An atmospheric and ionospheric model is developed for 120–220 km. The calculated temperature profile is characterized by a value of T ≈ 370°K at h ? 220 km, a steep gradient (～2°/km) at 200-160 km, a bend in the profile at 160 km, a small gradient (～0.7°/km) below and a value of T ≈ 250°K at 120 km. The upper point agrees well with the results of the Lyman-α measurements; the steep gradient may be explained by molecular viscosity dissipation of gravity and acoustical waves (the corresponding energy flux is 4 × 10^?2 erg cm^?2sec^?1 at 180 km). The bend at 160 km may be caused by a sharp decrease of the eddy diffusion coefficient and defines K ≈ 2 × 10⁸cm²sec^?1; and the low gradient gives an estimate of the efficiency of the atmosphere heating by the solar radiation as ? ≈ 0.1.     

Altitude profile of H in the atmosphere of Venus from Lyman α observations of Venera 11 and Venera 12 and origin of the hot exospheric component

Two extreme ultraviolet (EUV) spectrophotometers flown in December 1978 on Venera 11 and Venera 12 measured the hydrogen Lyman α emission resonantly scattered in the atmosphere of Venus. Measurements were obtained across the dayside of the disk, and in the exosphere up to 50,000 km. They were analyzed with spherically symmetric models for which the radiative transfer equation was solved. The H content of the Venus atmosphere varies from optically thin to moderately thick regions. A shape fit at the bright limb allows one to determine the exospheric temperature T_c and the number density n_c independently of the calibration of the instrument or the exact value of the solar flux. The dayside exospheric temperature was measured for the first time in the polar regions, with T_c = 300 ± 25°K for Venera 11 (79°S) and T_c = 275 ± 25°K (59°S) for Venera 12. At the same place, the density is n_c = 4_?2⁺³ × 10⁴ atom.cm^?3, and the integrated number density N_t from 250 to 110 km (the level of CO₂ absorption) is 2.1 × 10¹² atom.cm^?2, a factor of 3 to 6 lower than that predicted in aeronomical models. This probably indicates that the models should be revised in the content of H-bearing molecules and should include the effect of dynamics. Across the disk the value of N_t decreases smoothly with a total variation of two from the morning side to the afternoon side. Alternately it could be a latitude effect, with less hydrogen in the polar regions. The nonthermal component if clearly seen up to 40,000 km of altitude. It is twice as abundant as at the time of Mariner 10 (solar minimum). Its radial distribution above 4000 km can be simulated by an exospheric distribution with T = 10³⁰K and n = 10³ atom.cm^?3 at the exobase level. However, there are less hot atoms between 2000 and 4000 km than predicted by an ionospheric source. A by-product of the analysis is a determination of a very high solar Lyman α flux of 7.6 × 10¹¹ photons (cm² sec Å)^?1 at line center (1 AU) in December 1978.     

Vertical distribution of water vapor and mars model lower and middle atmosphere

Observations and model calculations of water vapor diffusion suggest that about half the amount of water vapor is distributed with constant mixing ratio in the Martian atmosphere, the other half is the excess water vapor in the lower troposphere. During 24 hr the total content of water vapor may vary by a factor of two. The eddy diffusion coefficient providing agreement between calculations and observations is K = (3–10) × 10⁶ cm² sec^?1 in the troposphere. An analytical expression is derived for condensate density in the stratosphere in terms of the temperature profile, the particle radius r, and K. The calculations agree with the Mars 5 measurements for r = 1.5 μm, condensate density 5 × 10^?12 g/cm³ in the layer maximum at 30 to 35 km, condensate column density 7 × 10^?6 cm^?2, K = (1?3) × 10⁶ cm² sec^?1, and the temperature profile T = 185 ? 0.05z ? 0.01z² at 20 to 40 km. Condensation conditions yield a temperature of 160°K at 60 km in the evening; the scale height for scattered radiation yields T = 110°k at 80 to 90 km. The Mars model atmosphere has been developed up to 125 km.     

Two-micron spectrophotometry of Uranus and its rings

Multiaperture K photometry and 2.0- to 2.5-μm spectrophotometry of Uranus and its ring system are presented. The photometric results are used, together with a previously published measurement, to set limits on the geometric albedos of Uranus and the rings at ～2.2 μm: (0.74 ± 0.02) × 10(su?4) ≤ p_K (Uranus) ≤ (1.5 ± 0.3) × 10^?4, and (2.7 ± 0.6) × 10^?2 ≤ p_K (rings) ≤ (3.4 ± 0.1) × 10^?2. Reflectance spectra of Uranus and Uranus plus rings show features in the planet's spectrum which are attributed to gaseous CH₄ absorption, and a 2.20-μm feature in the combined spectrum which may be due to the rings. This feature is tentatively identified with either the 2.26-μm absorption feature of NH₃ frost, or the 2.2-μm OH band exhibited by certain silicate minerals. The results of JHK photometry of Uranus' satellite, Ariel (U1), indicate that the infrared colors of this object are very similar to those of the satellites U2, U3, and U4.     

Radio emission from the quiet sun and local sources at wavelengths 5, 10.7, 12, and 95 cm from observations of the January 4, 2011 eclipse in Crimea

The radio radii of the Sun at wavelengths of 5, 10.7, 12, and 95 cm have been determined from eclipse observations as R₅ ?? (1.0 ± 0.015)R _??, R _10,12 = (1.05 ± 0.003)R _??, and R ₉₅ = (1.2 ± 0.02)R _??. The bright-ness temperatures of quiet solar disk areas at these wavelengths have turned out to be Td ₅ = (22 ± 2) × 10³, Td ₁₀ = (44 ± 3) × 10³, Td ₁₂ = (47 ± 3) × 10³, and Td ₉₅ = (1000 ± 30) × 10³ K. There were local sources of radio emission with angular sizes from 1.9 to 2.4 arcmin and brightness temperatures from 80 × 10³ to 1.75 × 10⁶ K above sunspot groups at short wavelengths of 5, 10.7, and 12 cm. The radio flux from the local sources at 95 cm turned out to be below the detection threshold of 1.0 × 10^?22 W m^?2 Hz^?1. Comparison of the values obtained with the results of observations of another eclipse on August 1, 2008, occurred at the epoch of minimum of the 11-year solar cycle has shown that the radio radius of the Sun at 10.7 and 12 cm increased from 1.016 R _?? to 1.05 ± 0.003R _??, the height of the emitting layer at these wavelengths moved from 11 × 10³ km to (30 ± 7) × 10³ K, and the brightness temperature of the quiet Sun rose from (35.8 ± 0.4) × 10³ K to (44 ± 3) × 10³ K at 10.7 cm and from (37.3 ± 0.4) × 10³ K to (47 ± 3) × 10³ K at 12 cm. Consequently, the parameters of the solar atmosphere changed noticeably in 2 years in connection with the beginning of the new solar cycle 24. The almost complete absence of local sources at the longest wavelength of 95 cm suggests that the magnetic fields of the sunspot groups on January 4, 2011, were weak and did not penetrate to the height from where their emission could originate. If this property is inherent in most sunspot groups of cycle 24, then it can be responsible for its low flare activity.     

Estimation of the galactic spiral pattern speed from Cepheids

To study the peculiarities of the Galactic spiral density wave, we have analyzed the space velocities of Galactic Cepheids with propermotions from the Hipparcos catalog and line-of-sight velocities from various sources. First, based on the entire sample of 185 stars and taking R ₀ = 8 kpc, we have found the components of the peculiar solar velocity (u _⊙, v _⊙) = (7.6, 11.6) ± (0.8, 1.1) km s^?1, the angular velocity of Galactic rotation Ω₀ = 27.5 ± 0.5 km s^?1 kpc^?1 and its derivatives Ω′₀ = ?4.12 ± 0.10 km s^?1 kpc^?2 and Ω″₀ = 0.85 ± 0.07 km s^?1 kpc^?3, the amplitudes of the velocity perturbations in the spiral density wave f _R = ?6.8 ± 0.7 and f _θ = 3.3 ± 0.5 km s^?1, the pitch angle of a two-armed spiral pattern (m = 2) i = ?4.6° ± 0.1° (which corresponds to a wavelength λ = 2.0 ± 0.1 kpc), and the phase of the Sun in the spiral density wave χ_⊙ = ?193° ± 5°. The phase χ_⊙ has been found to change noticeably with the mean age of the sample. Having analyzed these phase shifts, we have determined the mean value of the angular velocity difference Ω _p ? Ω, which depends significantly on the calibrations used to estimate the individual ages of Cepheids. When estimating the ages of Cepheids based on Efremov’s calibration, we have found |Ω _p ? Ω₀| = 10 ± 1_stat ± 3_syst km s^?1 kpc^?1. The ratio of the radial component of the gravitational force produced by the spiral arms to the total gravitational force of the Galaxy has been estimated to be f _r0 = 0.04 ± 0.01.     

The global distribution of O3 on mars

Models are developed to describe the photochemistry of ozone on Mars. Catalytic reactions involving H, OH and HO₂ play a major role at low latitudes where they ensure a vertical column density for O₃ of less than 2 × 10^?4 cm atm. The source for odd hydrogen (H + OH + HO₂) is relatively smaller at high latitudes in winter due to the small concentrations of H₂O present there at that time. Odd hydrogen is also efficiently removed from the high-latitude winter atmosphere by condensation of H₂O₂. The role of catalytic chemistry is reduced accordingly and the vertical column density of O₃ may be as large as 5.7 × 10^?3 cm atm in accord with earlier observations carried out by Barth and co-workers with instruments on Mariner 9.     

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.     

A seasonal climate model of the atmospheres of the giant planets at the voyager encounter time: I. Saturn's stratosphere

A radiative seasonal model which incorporates a multilayer radiative transfer treatment at wave-lengths longward of 7 μm is presented and applied to Saturn's stratosphere. Opacities due to H₂-He, CH₄, C₂H₂, and C₂H₆ are included. Season-dependent insolation is shown to produce a strong hemispheric asymmetry decreasing with depth at the Voyager encounter times, and seasonal amplitudes of 30°K at the poles are predicted in the high stratosphere. The ring-modulated dependence of the insolation and the orbital eccentricity are shown to have a significant effect. Calculations agree closely with the Voyager 1 and 2 radio occultation ingress profiles recorded at 76°S and 36.5°S for CH₄/H₂ = 3.5 + 1.4/? 1.0 × 10^?3;the estimated errors include modeling systematic errors and uncertainties in the occultations profiles. The possible role of aerosols in the stratospheric heating is analyzed. The Voyager 2 egress profile recorded at 31°S cannot be reproduced by calculations. Some constraints on the C₂H₂ and C₂H₆ abundances are derived. The upper portion of the occultation profiles (p < 3mbar) can be matched for C₂H₂/H₂ = 1.0 + 1.3/?0.6 × 10^?7, C₂H₆/H₂ = 1.5 + 1.8/?0.9 × 10^?6 at 76°S and C₂H₂/H₂ = 4 + 6/?4 × 10^?8, C₂H₆/H₂ = 6 + 9/?6 × 10^?7 at 36.5°N. At the northern occultation latitude, the discrepancy with the concentrations derived from analysis of IRIS spectra by R. Courtin, D. Gautier, A. Marten, B. Bézard, and R. Hanel (1984, Astrophys. J.287) can be explained by a sharp variation of the mixing ratios of these gases with altitude in the upper stratosphere. Other interpretations are discussed.     

Distribution and relations of 4- to 10-km-diameter craters to global geologic units of Mars

By correlating the 1:25,000,000 geologic map of Mars of Scott and Carr (1977) with 4- to 10-km-diameter crater density data from Mariner 9 images, the average crater density for 23 of the equatorial geologic-geomorphic units on Mars was computed. The correlation of these two data sets was accomplished by digitizing both the crater density data and geologic map at the same scale and by comparing them in a computer. This technique assigns the crater density value found in the corresponding location on the geologic data set to a discrete computer file assigned each of the 23 geologic units. By averaging the crater density values accumulated in each file, an “average” crater density for each geologic unit was obtained. Condit believes these average crater density values are accurate indicators of the relative age of the geologic units considered. The statistical validity of these average values is strongest for the geologic units of the largest areal extent. The relative ages as obtained from the average crater density values for the seven largest geologic units, from youngest to oldest, are: Tharsis volcanic material, 21 ± 4 craters/10⁶km²; smooth plains material, 57 ± 14 craters/10⁶km²; rolling plains material, 66 ± 16 craters/10⁶km²; plains materials, 80 ± 17 craters/10⁶km²; ridged plains material, 128 ± 25 craters/10⁶km²; hilly and cratered material, 137 ± 38 craters/10⁶km²; and cratered plateau material, 138 ± 27 craters/10⁶km².     

An estimate of the PH3, CH3D,and GeH4 Abundances on Jupiter from the Voyager IRIS data at 4.5 μm

The abundances of PH₃, CH₃D, and GeH₄ are derived from the 2100- to 2250-cm^?1 region of the Voyager 1 IRIS spectra. No evidence is seen for large-scale variations of the phosphine abundance over Jovian latitudes between ?30 and +30°. In the atmospheric regions corresponding to 170–200°K, the derived PH₃/H₂ value is (4.5 ± 1.5) × 10^?7 or 0.75 ± 0.25 times the solar value. This result, compared with other PH₃ determinations at 10 μm, suggests than the PH₃/H₂ ratio on Jupiter decreases with atmospheric pressure. In the 200–250°K region, we derive, within a factor of 2, CH₃D/H₂ and GeH₄/H₂ ratios of 2.0 × 10^?7 and 1.0 × 10^?9, respectively. Assuming a C/H value of 1.0 × 10^?3, as derived from Voyager, our CH₃D/H₂ ratio implies a D/H ratio of 1.8 × 10^?5, in reasonable agreement with the interstellar medium value.     

The penetration limit of thin films

One sensor of the Helios micrometeoroid experiment is covered by a thin film consisting of 3000 Å parylene and 750 Å aluminium. Micrometeoroids must penetrate this film before they are detected. In order to study the effects of the film on the detection of micrometeoroids simulation experiments were performed with iron, aluminium, glass and polyphenylene projectiles in the mass range of 5 × 10^?13g < m < 2 × 10^?10g and in the speed range of 1.5 km/sec <ν < 13 km/sec. The bulk densities of the projectiles ranged from 1.25 g/cm³ (polyphenylene) to 7.9 g/cm³ (iron). By measuring the speed of the projectiles before and after the film penetration the speed loss Δν caused by the film was determined. The angle of incidence was varied in three steps (0°, 30° and 60°). This deceleration strongly depends on the projectiles' densities: Vertically impacting iron projectiles of m = 10^?11g and ν₁ = 3 km/sec were subject to a relative speed loss of Δν/ν₁ = 4%, aluminium projectiles of the same mass and speed showed Δν/ν₁ = 8%, glass projectiles Δν/ν₁ = 9% and polyphenylene projectiles Δν/ν₁ = 14%. The total charge of the plasma produced upon impact on a gold target of a projectile which had penetrated the film before that was compared with the plasma produced by a projectile without a penetration. For iron projectiles these two signals did not differ significantly even at an angle of incidence of 60°. Whereas polyphenylene projectiles showed an attenuation of the charge signal by a factor of 10 after the penetration already at an angle of incidence of 0°. Polyphenylene projectiles impacting the film at an angle of incidence of 60° could no longer be detected behind the film. This experiment defined the penetration limit of the Helios film. Comparison with other penetration data yielded a penetration formula which is applicable to projectiles with diameters in the submicron to centimeter range. This penetration formula gives the penetration limit of a film as a function of the projectile's mass, speed and density.     

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.