Thermal inertia mapping of Mars from 60°S to 60°N

Twenty-micrometer brightness temperatures are used to derive the thermal inertia for 81% of the Martian surface between latitudes ±60°. These data were acquired by the two Viking Infrared Thermal Mappers in 1977 and 1978 following the two global dust storms of 1977. The spatial resolution used is 2° in latitude by 2° in longitude and the total range in derived inertia is $\text{1}\text{to}\text{15 × 10}{}^{\mathrm{?3}}\text{cal cm}{}^{\mathrm{?2}}\text{sec}{}^{\mathrm{<mtext>?</mtext><mtext>1</mtext><mtext>2</mtext>}}\text{°}\text{K}{}^{\mathrm{?1}}$. The distribution of thermal inertia is strongly bimodal with all values of thermal inertia less than $\text{4 × 10}{}^{\mathrm{?3}}\text{cal cm}{}^{\mathrm{?2}}\text{sec}{}^{\mathrm{<mtext>?</mtext><mtext>1</mtext><mtext>2</mtext>}}\text{°}\text{K}{}^{\mathrm{?1}}$ being associated with three disjoint bright regions mostly in the northern hemisphere. Sufficient dust is raised in global storms to provide fine material adequate to produce these low-inertia areas but the specific deposition mechanism has not been defined. At the low resolution used, no complete exposures of clean rock were found. There is some tendency for darker material to be associated with higher thermal inertia, although the trend is far from one to one. The distribution of high- and low-inertia areas is sufficiently nonrandom to produce a variation in whole-disk brightness temperature with central meridian longitude. This variation and the change in surface kinetic temperature associated with dust storms are factors in establishing the whole-disk brightness temperature at radio and infrared wavelengths and will be important for those who use Mars as a calibration source.     

Physical properties of the particles composing the Martian dust storm of 1971–1972

Infrared spectra obtained from the Mariner 9 spacecraft during the 1971–1972 dust storm are used to derive information on the composition and particle size distribution of the dust and to study the time evolution of the storm. The dust is not composed of pure granite, basalt, basaltic glass, obsidian, quartz, andesite, or montmorillonite. The infrared spectra suggest that the dust is a mixture of materials, dominated by igneous silicates with >62;60% SiO₂, or weathering products such as clay minerals, but the dust could possibly have a significant component of lower SiO₂ materials such as basalt. Substantial quantities of carbonates, nitrates, or carbon suboxide are excluded from the mixture. All infrared, visible, and ultraviolet data on the Martian surface composition seem consistent with a mixture of basalt and clay minerals or high SiO₂ igneous rocks, with a surface patina of oxides of iron. For all candidate compositions, the data are best matched with a size distribution that approximates a differential power law function of slope ?4. This size distribution is quite similar to terrestial size distributions in regions remote from sources of dust. The relative abundance of particles between 1- and 10-μm radius did not change during the Mariner 9 mission; thus suspended particles did not experience Stokes-Cunningham fallout but instead were supported by turbulence with an eddy diffusion coefficient, $\text{K}{}_{\mathrm{<mtext>e</mtext>}}\text{? 7 × 10}{}^6\text{cm}{}^2\text{sec}{}^{\mathrm{?1}}$. The aerosol optical depth, standardized to 0.3-μm wavelength, varied from about 1.5 early in the mission to about 0.2 at Orbit 200.     

Small-scale turbulence in the atmosphere of Venus

Previous studies based on radio scintillation measurements of the atmosphere of Venus have identified two regions of small-scale temperature fluctuations located in the vicinity of 45 and 60 km. A global study of the fluctuations near 60 km, which are consistent with wind-shear-generated turbulence, was conducted using the Pioneer Venus measurements. The structure constants of refractive index fluctuations c_n² and temperature fluctuations c_T² increase poleward, peak near 70° latitude, and decrease over the pole; c_n² varies from 2 × 10^?15 to 1.5 × 10^?14m^{$\text{2}\text{3}$} and c_T² from 4 × 10^?3 to 7 × 10^?2°K²m^{$\text{?2}\text{3}$}. These results indicate greater turbulent activity at the higher latitudes. In the region near 45 km the refractive index fluctuations and the corresponding temperature fluctuations are substantially lower. Based on the analysis of one representative occultation measurement, $\text{c}{}_{\mathrm{<mtext>n</mtext>}}{}^2\text{= 2 × 10}{}^{\mathrm{?16}}\text{m}{}^{\mathrm{<mtext>?2</mtext><mtext>3</mtext>}}\text{and}\text{c}{}_{\mathrm{<mtext>T</mtext>}}{}^2\text{= 7.3 × 10}{}^{\mathrm{?4}}\text{°}\text{K}{}^2\text{m}{}^{\mathrm{<mtext>?2</mtext><mtext>3</mtext>}}$ in the 45-km region. The fluctuations in this region also appear to be consistent with wind-shear-generated turbulence. The turbulence level is considerably weaker than that at 60 km; the energy dissipation rate ε is 4.9 × 10^?5m²sec^?3 and the small-scale eddy diffusion coefficient K is 2 × 10³ cm² sec^?1.     

10 μm polarimetry of ceres

Linear polarimetry of Ceres at 10 μm is presented. These data represent the first published polarization measurements of an asteroid in the thermal infrared. It is found that Ceres is polarized at the 0.2-0.6% level. This data set is compared with theoretical models of the linear polarization of emitted radiation from a spherical plane. These models are used to derive the pole position and thermal inertia of Ceres. Ceres is best fit with a thermal inertia of $\text{0.0010±0.0003}\text{cal}\text{cm}{}^{\mathrm{?2}}\text{°}\text{K}{}^{\mathrm{?1}}\text{sec}{}^{\mathrm{<mtext>1</mtext><mtext>2</mtext>}}$ and a pole orientation of β_p = 36° ± 5°, λ_p = 270° ± 3°. It is concluded that 10μm polarimetry is a potentially powerful technique for remotely sensing the pole orientation and thermal inertia of asteroids.     

Io: Observational constraints on internal energy and thermophysics of the surface

Infrared observations of the Io eclipse of 12 April 1980 in five broad bands from 3 to 30 μm define the thermal emission spectrum both during and after eclipse. A substantial fraction of the emitted radiation during eclipse arises from hot spots; the equivalent global average heat flow is 1.5 ± 0.3 W m^?2, corresponding to an internal source of (6 ± 1) × 10¹³ W. The hot spot spectra can be matched by components with color temperatures of 200–600°K covering 1–2% of the surface. Comparison with observations over the past 8 years suggests that, while the flux at the hottest temperatures may be highly variable, there is no evidence for major changes in the total heat flow, which is emitted primarily in the spectral region 10–20 μm. The heating curves of the surface were observed at 10 and 20 μm; when corrected for the hot spot contribution they indicate a typical global thermal inertia for Io of $\text{(0.2 ± 0.1) × 10}{}^{\mathrm{?3}}\text{cal cm}{}^{\mathrm{?2}}\text{sec}\text{?}{}^{\mathrm{<mtext>1</mtext><mtext>2</mtext>}}\text{K}{}^{\mathrm{?1}}$, similar to that of the other Galilean satellites.     

Geopotential coefficients of order 29, from analysis of the orbit of satellite 1967-104B

The orbit of the satellite 1967-104B has been analysed as it passed through 29:2 resonance with the Earth's gravitational field between January 1977 and September 1978. From the changes in inclination and eccentricity the following lumped 29th-order geopotential harmonic coefficients were obtained: $\text{10}{}^9\text{C}\text{?}{}_{29}{}^{0.2}\text{= 4.1 ± 0.8}$, $\text{10}{}^9\text{S}\text{?}{}_{29}{}^{0.2}\text{= 10.3 ± 2.4}$, $\text{10}{}^9\text{C}\text{?}{}_{29}{}^{1.1}\text{= ? 160 ± 19}$, $\text{10}{}^9\text{S}\text{?}{}_{29}{}^{1.1}\text{= 79 ± 10}$, $\text{10}{}^9\text{C}\text{?}{}_{29}{}^{\mathrm{?1.3}}\text{= 38 ± 14}$, $\text{10}{}^9\text{S}\text{?}{}_{29}{}^{\mathrm{?1.3}}\text{= 19 ± 5}$. These values have been compared with existing comprehensive geopotential models: the best agreement is with the model of Rapp (1981).     

The determination and analysis of the orbit of NIMBUS 1 rocket, 1964-52B: Variations in orbital inclination

The influence of aerodynamic drag and the geopotential on the motion of the satellite 1964-52B is considered. A model of the atmosphere is adopted that allows for oblateness, and in which the density behaviour approximates to the observed diurnal variation. A differential equation governing the variation of the orbital inclination combining the effects of air drag with those of the Earth's gravitational field is given.The 310 observed values of inclination are modified by the removal of perturbations due to luni-solar attraction, solid Earth and ocean tides, solar radiation pressure, low-order long-periodic tesseral harmonic perturbations and changes due to precession. The method of removal of these effects is given in some detail.The variations in inclination due to drag are analysed to give four values of the average atmospheric rotation rate at heights of 296–476 km at latitude 0–54°. These values are as expected from previous analyses.The analysis of the change in inclination due to solar radiation pressure shows that this rapidly tumbling cylindrical satellite may be considered as equivalent to a spherical satellite of a given area-to-mass ratio.Analysis of the inclination near 15:1 resonance with the geopotential yields values of lumped geopotential harmonics of order 15 and 30, namely, $\text{10}{}^9\text{C}\text{?}{}^{0.1}{}_{15}\text{= ?31.2 ± 2.3 10}{}^9\text{S}\text{?}{}^{0.1}{}_{15}\text{= ?4.4 ± 3.2 10}{}^9\text{C}\text{?}{}^{0.2}{}_{30}\text{= 39.0 ± 10.7 10}{}^9\text{S}\text{?}{}^{0.2}{}_{30}\text{= 51.8 ± 10.0}$     

A dynamical effect in the ionospheric f1 layer

Incoherent scatter observations of the ionospheric F₁ layer above Saint-Santin (44.6°N) are analyzed after correction of a systematic error at 165 and 180 km altitude. The daytime valley observed around 200 km during summer for low solar activity conditions is explained in terms of a downward ionization drift which reaches ?30 m s^?1 around 180 km. Experimental determinations of the ion drift confirm the theoretical characteristics required for the summer daytime valley as well as for the winter behaviour without a valley. The computations require an effective dissociative recombination rate of $\text{2.3 × 10}{}^{\mathrm{?7}}\text{(}\text{300/}\text{T}{}_{\mathrm{e}}\text{)}{}^{0.7}\text{(}\text{cm}{}^3\text{s}{}^{\mathrm{?1}}\text{)}$ and ionizing fluxes compatible with solar activity conditions at the time when the valley is observed.     

Martian glaciation and the flow of solid CO2

The flow law determined experimentally for solid CO₂ establishes that a hypothesis of glacial flow of CO₂ at the Martian poles is not physically unrealistic. Compression experiments carried out under 1 atm pressure and constant strain rate demonstrate that the strength of CO₂ near its sublimation point is considerably less than the strength of water ice near its melting point. The data fit a power law “creep” equation of the form

$\text{?}\text{?}\text{= (4 × 10}{}^6\text{) σ}{}^{3.9}\text{exp}\text{(}\text{?12 200}\text{RT}\text{)}$

, where ? is compressive strain rate (sec^?1), σ is compressive stress (bars), R is the gas constant in calories per mole, and T is absolute temperature. The exponent of σ of 3.9 contrasts with a value near 3.1 for water ice, and indicates that the strain rate is somewhat more sensitive to stress for CO₂ than for water. Likewise, the low activation energy for creep, 12 200 cal mole^?1, illustrates that CO₂ is not highly sensitive to temperature and is thus likely to flow over a broad range of temperatures below its melting point. Strength values for CO₂ are of the order of one-tenth to one-third the strength of ice under equivalent conditions.A plausible glacial model for the Martian polar caps can be constructed and is helpful in explaining the unique character of the polar regions. CO₂-rich layers deposited near the pole would have flowed outward laterally to relieve high internal shear stresses. The topography of the polar caps, the uniform layering of the layered deposits, and the general extent of the polar “sediments” could all be explained using this model. Flow of CO₂ rather than water ice greatly reduces the problems with Martian glaciation. Nevertheless, problems do remain, in particular the large amounts of CO₂ necessary, the need to increase vapor pressure and temperature with depth in the polar deposits, and the lack of good observational evidence of flor features. Within the limits of the present knowledge of surface conditions of Mars, CO₂ glaciation appears to be a realistic alternate working hypothesis for the origin of the polar features.     

Mariner 10 infrared radiometer results: Temperatures and thermal properties of the surface of Mercury

Mariner 10 infrared brightness temperatures of the surface of Mercury at 11 and 45 μm are presented. The data were obtained during the first flyby along a nera-equatorial swath extending from 17 hours local time through local midnight to 9 hours local time. For an assumed emissivity of 0.9, derived surface thermal inertias are between $\text{0.0031}\text{and}\text{0.0031}\text{cal cm}{}^{\mathrm{?2}}\text{sec}{}^{\mathrm{<mtext>?1</mtext><mtext>2</mtext>}}\text{K}{}^{\mathrm{?1}}$ and the implied minimum predawn surface kinetic temperature for the warm pole at longitude 270° is near 93 K. Several pronounced thermal inhomogeneities were seen, one of which appears to coincide with a region of high radar reflectivity. The derived thermal properties and the electrical skin depth and loss tangent fall within the range of values found on the Moon.     

The determination and analysis of the orbit of NIMBUS 1 ROCKET, 1964-52B: Part 2: Variations in orbital eccentricity

The influence of aerodynamic drag and the geopotential on the motion of the satellite 1964-52B is considered. A model of the atmosphere is adopted that allows for oblateness, and in which the density behaviour approximates to the observed diurnal variation. A differential equation governing the variation of the eccentricity, e, combining the effects of air drag with those of the Earth's gravitational field is given. This is solved numerically using as initial conditions 310 computed orbits of 1964-52B.The observed values of eccentricity are modified by the removal of perturbations due to luni-solar attraction, solid Earth and ocean tides, solar radiation pressure and low-order long-periodic tesseral harmonic perturbations. The method of removal of these effects is given in some detail. The behaviour of the orbital eccentricity predicted by the numerical solution is compared with the modified observed eccentricity to obtain values of atmospheric parameters at heights between 310 and 430 km. The daytime maximum of air density is found to be at 14.5 hours local time. Analysis of the eccentricity near 15th order resonance with the geopotential yielded values of four lumped geopotential harmonics of order 15, namely: $\text{10}{}^9\text{C}{}^{\mathrm{1,0}}{}_{15}\text{= ?78.8 ± 7.0}$, $\text{10}{}^9\text{S}{}^{\mathrm{1,0}}{}_{15}\text{= ?69.4 ± 5.3}$, $\text{10}{}^9\text{C}{}^{\mathrm{?1,2}}{}_{15}\text{= ?41.6 ± 3.5}$$\text{10}{}^9\text{S}{}^{\mathrm{?1,2}}{}_{15}\text{= ?26.1 ± 8.9}$, at inclination 98.68°.     

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

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

Analysis of the orbital inclination of Tansei 3 rocket 1977-12B

The orbit of Tansei 3rocket(1977-12B) has been determined at 47 epochs between 1 October 1977 and 19 March 1979 using over 1700 observations and the RAE orbit refinement program PROP6. The rate of change of the inclination was examined to evaluate values of the atmospheric rotation rate, Λ rev day^?1. Analysis yielded the value Λ = 1.1 ± 0.05 at height 315 ± 30 km, average conditions; or alternatively Λ = 1.1 ± 0.1 at height 347 ± 12 km, slight winter bias and Λ = 1.07 ± 0.1 at height 270 ± 18 km, average conditions, supplying further evidence of a decrease in rotation rates from the 1960s to the 1970s.Analysis of the inclination at 15th-order resonance yielded the lumped harmonic values

$\text{10}{}^9\text{C}{}^{\mathrm{0,1}}{}_{15}\text{= 13.4 ± 6.2, 10}{}^9\text{S}{}^{\mathrm{0,1}}{}_{15}\text{= 0.7 ± 13.3}$

for inclination 65.485°.     

15th order resonance terms using the decaying orbit of TETR-3

The orbit of TETR-3 (1971-83B), inclination: 33°, passed through resonance with 15th order geopotential terms in February 1972. The resonance caused the orbit inclination to increase by 0.015°. Analysis of 48 sets of mean Kepler elements for this satellite in 1971–1972 (across the resonance) has established the following strong constraint for high degree, 15th order gravitational terms (normalized):

$\text{10}{}^9\text{(}\text{C, S}\text{)}{}_{15}\text{= (28.3 ± 3.0, 7.4 ± 3.0) =}\text{0.001(}\text{C, S}\text{)}{}_{\mathrm{15,15}}\text{?0.015(}\text{C, S}\text{)}{}_{\mathrm{17,15}}\text{+0.073(}\text{C, S}\text{)}{}_{\mathrm{19,15}}\text{?0.219(}\text{C, S}\text{)}{}_{\mathrm{21,15}}\text{+0.477(}\text{C, S}\text{)}{}_{\mathrm{23,15}}\text{?0.781(}\text{C, S}\text{)}{}_{\mathrm{25,15}}\text{+1.000(}\text{C, S}\text{)}{}_{\mathrm{27,15}}\text{?0.0963(}\text{C, S}\text{)}{}_{\mathrm{29,15}}\text{+0.622(}\text{C, S}\text{)}{}_{\mathrm{31,15}}\text{?0.119(}\text{C, S}\text{)}{}_{\mathrm{33,15}}\text{?0.290(}\text{C, S}\text{)}{}_{\mathrm{35,15}}\text{+0.403(}\text{C, S}\text{)}{}_{\mathrm{37,15}}\text{?0.223(}\text{C, S}\text{)}{}_{\mathrm{39,15}}\text{?0.058(}\text{C, S}\text{)}{}_{\mathrm{41,15}}\text{+…}$

This result combined with previous results on high inclination 15th order and other resonant orbits suggests that the coefficients of the gravity field beyond the 15th degree are smaller than Kaula's rule ($\text{10}{}^{\mathrm{?5}}\text{l}{}^2$).     

The locating of hydromagnetic whistler sources and determination of their generating proton spectra

Results are given of the calculations of the group delay time propagating τ(ω, φ₀) of hydromagnetic whistlers, using outer ionospheric models closely resembling actual conditions. The τ(ω, φ₀) dependencies were compared with the experimental data of τ_exp(ω, φ₀) obtained from sonagrams. The sonagrams were recorded in the frequency range $\text{? ? (0.5?2.5) Hz}$ at observation points located at geomagnetic latitudes φ₀ = (53?66)° and in the vicinity of the geomagnetic poles. This investigation has led us to new and important conclusions.The wave packets (W.P.) forming hydromagnetic whistlers (H.W.) are mainly generated in the plasma regions at L = 3.5?4.0. This is not consistent with ideas already expressed in the literature that their generation region is $\text{L ? 3?10}$. The overwhelming majority of the τ_exp values differ considerably from the times at which wave packets would, in theory, propagate along the magnetic field lines corresponding to those of the geomagnetic latitudes φ₀ of the observation points. The second important fact is that the W.P. frequency ω is less than $\text{Ω}{}_{\mathrm{H}}$ everywhere along its propagation trajectory, including the apogee of the magnetic force line ($\text{Ω}{}_{\mathrm{H}}$ is the proton gyrofrequency). Proton flux spectra $\text{E ? (30?120)}$ keV, responsible for H.W. generation, were determined. Comparison of the Explorer-45 and OGO-3 measurements published in the literature, with our data, showed that the proton flux density energy responsible for the H.W. excitation $\text{N}{}_{\mathrm{p}}\text{(}\text{MV}{}_6{}^2\text{2}\text{) ? (5 × 10}{}^{\mathrm{?3}}\text{?10}{}^{\mathrm{?1}}\text{)}\text{H}{}_{\mathrm{a}}{}^2\text{8π}$ where H_a is the magnetic field force in the generation region of these W.P. The electron concentration is $\text{N}{}_{\mathrm{a}}\text{? (10}{}^2\text{?10}{}^3\text{) cm}{}^{\mathrm{?3}}$. The values given in the literature are $\text{N}{}_{\mathrm{a}}\text{? (10}{}^{\mathrm{?10}}\text{?10}{}^3\text{) cm}{}^{\mathrm{?3}}$. The e data considered also leads to the conclusion that the generating mechanism of the W.P. studied probably always co-exists with the mechanism of their amplification.     

On the change of period of Dy Pegasi

In the course of testing a 3-channel photon-counting, high-speed photometer, we observed DY Peg on three nights, 1980 September 18, 19, and November 15, and obtained 7 light maxima. Combining with results observed over the past 30 years, we re-calculated the formula for the maximum epoch to be

$\text{T}{}_{\max }\text{=}\text{H.J.D.}\text{2432751.96182 + 0.072926362}\text{E}\text{? 2.28 × 10}{}^{\mathrm{?13}}\text{E}{}^2\text{± 0.00008 ± 0.000000001 ± 0.09}$

and the period decay rate to be (3.1± 0.1) × 10^?8yr^?1.     

Photometric properties of zodiacal light particles

From published ground-base, spacecraft, and rocket photometry and polarimetry of the zodiacal light, a number of optical and physical parameters have been derived. It was assumed that the number density, mean particle size, and albedo vary with heliocentric distance, and shown that average individual interplanetary particles have a small but definite opposition effect, a mean single-scattering albedo in the V band at 1-AU heliocentric distance of 0.09 ± 0.01, and a zero-phase geometric albedo of 0.04. Modeled by a power law, both albedos decrease with increasing heliocentric distance as r^?0.54. The corresponding exponents for changes in mean particle size and number density are related in a simple way. The median orbital inclination of zodiacal light particles with respect to the ecliptic is 12°, close to the observed median value for faint asteroids and short-period comets. Furthermore, the color of dust particles and its variation with solar phase angle closely resemble those of C asteroids. These findings are, at least, consistent with the zodiacal cloud originating primarily from collisions among asteroids. Finally, a value of ?10¹⁸?_ErmE g was derived for the mass of the zodiacal cloud, where ?_E is the mean particle radius (in micrometers) at 1-AU-heliocentric distance. For extinction in the ecliptic, $\text{Δm = 10}{}^{\mathrm{?5}}\text{?}{}^{\mathrm{<mtext>?1</mtext><mtext>2</mtext>}}\text{mag}$ was obtained, where ? is the solar elongation in degrees.     

Quantal calculation of the lifetime of N+(5S) and the λ2145Å auroral mystery feature

The calculated radiative lifetime of the metastable ion is 6.4 × 10^?3s. Used in conjunction with the results of measurements by Erdman, Espy and Zipf this sets 1.3 × 10^?18 cm² as the upper limit to the cross section for the formation of $\text{N}{}^{\mathrm{+}}\text{(}{}^5\text{S}\text{)}$ in e - N₂ collisions at 100eV which leaves the possibility that the process is responsible for the $\text{λ2145}\text{A}\text{?}$ feature in auroras only just open. The cross section for the formation of $\text{N}{}^{\mathrm{+}}\text{(}{}^5\text{S}\text{)}$ in e — N collisions is large. However for this process to lead to the observed intensity of $\text{λ2145}\text{A}\text{?}$ relative to $\text{λ3914}\text{A}\text{?}$ the N:N₂ abundance ratio would have to be as high as 1.6 × 10^?2.     

Airborne spectroscopy and spacecraft radiometry of Venus in the far infrared

In March 1979, the spectrum of Venus was recorded in the far infrared from the G.P. Kuiper Airborne Observatory when the planet subtended a phase angle of 62°. The brightness temperature was observed to be 275°K near 110 cm^?1, dropping to 230°K near 270 cm^?1. Radiance calculations, using temperature and cloud structure formation from the Pioneer Venus mission and including gaseous absorption by the collision-induced dipole of CO₂, yield results consistently brighter than the observations. Supplementing the spectral data, Pioneer Venus OIR data at similar phase angles provide the constraint that any additional infrared opacity must be contained in the upper cloud, H₂SO₄ to the Pioneer-measured upper cloud structure serves to reconcile the model spectrum and the observations, but cloud microphysics strongly indicates that such a high particle density haze $\text{(N ? 1.6 × 10}{}^7\text{cm}{}^{\mathrm{?3}}\text{)}$ is implausible. The atmospheric environment is reviewed with regard to the far infrared opacity and possible particle distribution modifications are discussed. We conclude that the most likely possibility for supplementing the far-infrared opacity is a population of large particles $\text{(}\text{r}\text{? 1 μm)}$ in the upper cloud with number densities less than 1 particle cm^?3 which has remained undetected by in situ measurements.