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.     

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.     

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.     

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°.     

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.     

Linewidth studies on the NI(4S−4P) resonance multiplet

We have measured the linewidths of the NI multiplets $\text{[2p}{}^2\text{3p}{}^4\text{D}{}^0\text{?2p}{}^2\text{3s}{}^4\text{P, λ8691}\text{A}\text{?}\text{; 2p}{}^2\text{3p}{}^4\text{P}{}^0\text{?2p}{}^2\text{3s}{}^4\text{P, λ8212}\text{A}\text{?}\text{; 2p}{}^2\text{3s}{}^4\text{P?2p}{}^3{}^4\text{S}{}^0\text{, λ1200}\text{A}\text{?}\text{]}$ produced in the dissociative excitation of N₂ by energy electrons. The infrared transitions excite the N(⁴P) resonance state by cascade and they account for > 50% of the total N(⁴P) cross section at 100 eV. Both the i.r. and v.u.v. lines are found to be highly Doppler broadened ( ～ 25 times the thermal Doppler line width). These results indicate that dissociative excitation of N₂ produces N (⁴P) atoms with sufficient kinetic energy so that the $\text{λ1200}\text{A}\text{?}$ resonance radiation $\text{[2p}{}^2\text{3s}{}^4\text{P ?2p}{}^3{}^4\text{S}{}^0\text{]}$ emitted by these excited atoms would be optically thin in the Earth's upper atmosphere. We also found that the line strength ratios for the resolved components of the $\text{λ1200}\text{A}\text{?}$ triplet excited by dissociative excitation differ from those predicted by the multiplicities of the states involved and used in current entrapment models; the intensity ratios also vary with the energy of the incident electron. These developments introduce new complications into the analysis of the terrestrial ultraviolet dayglow.     

A self-consistent evaluation of the rate constants for the production of the OI 6300 Å airglow

The quenching rate k_N2 of $\text{O(}{}^1\text{D)}$ by N₂ and the specific recombination rate α_1D of O₂⁺ leading to $\text{O(}{}^1\text{D)}$ are re-examined in light of available laboratory and satellite data. Use of recent experimental values for the $\text{O(}{}^1\text{D)}$ transition probabilities in a re-analysis of AE-C satellite 6300 Å airglow data results in a value for k_N2 of 2.3 × 10^?11 cm³s^?1 at thermospheric temperatures, in excellent agreement with the laboratory measurements. This implies a value of J_O2 = 1.5 × 10^?6s^?1 for the O₂ photodissociation rate in the Schumann-Runge continuum. The specific recombination coefficient α_1D = 2.1 × 10^?7cm³s^?1 is also in agreement with the laboratory value. Implications for the suggested $\text{N}\text{(}{}^2\text{D) + O}{}_2\text{→ O(}{}^1\text{D) + NO}$ reaction are discussed.     

Eolian intracrater deposits on Mars: Physical properties and global distribution

An investigation of Martian intracrater materials has been made using their thermophysical properties as derived from Viking IRTM observations. Over one-fourth of all craters larger than 25 km in diameter between ?50°S and 50°N have localized deposits of coarse material on the floor which are associated with dark “splotches” seen visually. Assuming homogeneous, unconsolidated materials, the measured thermal inertias of these deposits ($\text{I = 0.003 × 10}{}^{\mathrm{?3}}\text{to}\text{0.026 × 10}{}^{\mathrm{?3}}\text{cal cm}{}^{\mathrm{?2}}\text{sec}{}^{\mathrm{<mtext>?</mtext><mtext>1</mtext><mtext>2</mtext>}}\text{°}\text{K}{}^{\mathrm{?1}}$) imply effective grain sizes ranging from 0.1 mm to 1 cm, with a modal value of 0.9 mm. These deposits are coarser and darker than the surrounding terrains and the majority of the Martian surface, but are not compositionally distinct from materials with similar albedos. They occur more frequently in the south, in regions of relatively coarse material (0.2 to 2 mm), and in relatively dark areas. These features most likely formed by entrapment of marginally mobile material which can be transported into, but not out of, crater depressions by the wind. Very few have recognizable dune forms: those that do have effective grain sizes less than 0.5 mm. The majority of the “splotch” deposits are coarser than the dune-forming materials found in the north polar region and inside extreme southern latitude craters and probably form low, broad zibar dunes or lag deposits. Intracrater deposits are noticeably lacking from the interior of the large, northern hemisphere low-inertia region of Arabia (?10°S to 30°N, 300° to 360°W), interpreted to be a sink for suspended dust, but do occur around the perimeter of this region. This distribution suggests that the intracrater features have been buried in the interior of Arabia and that the dust deposit is less extensive at the margins and may currently be expanding. The occurrence of regional dust deposits in the north may be related to the maximum wind activity currently occurring in the southern hemisphere and suggests that the location of regional sinks may migrate with time as the solar insolation maximum migrates.     

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).     

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.     

Airglow from the inner comas of comets

The intensities of radiation from the inner comas of comets which are composed primarily of water and carbon monoxide have been calculated. Only “airglow” emissions initiated by the absorption of extreme ultraviolet radiation have been considered. The photoionizations of H₂O, CO, CO₂, and N₂ are the most important emission sources, although photoelectron excitation is also considered. Among the emission features for which intensities were calculated are H₂O⁺ ($\text{A}\text{?}{}^2\text{A}{}_{\mathrm{1?}}\text{X}\text{?}{}^2\text{B}{}_1$), CO⁺ (first negative), CO (fourth positive), CO (Cameron), CO₂⁺ ($\text{B}\text{?}{}^2\text{?}{}_{\mathrm{u}}\text{?}\text{X}\text{?}{}^2\text{II}{}_{\mathrm{g}}$), N₂ (Vegard-Kaplan), N₂⁺ (first negative), and OI (1304 Å). In the inner coma (collision region) these airglow mechanisms are shown to be possible competitors with the usually assumed resonance scattering and flourescence excitation mechanisms which are appropriate for the outer coma and tail.     

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.     

Titan: Far-infrared and microwave remote sensing of methane clouds and organic haze

It is shown that Titan's surface and plausible atmospheric thermal opacity sources—gaseous N₂, CH₄, and H₂, CH₄ cloud, and organic haze—are sufficient to match available Earth-based and Voyager observations of Titan's thermal emission spectrum. Dominant sources of thermal emission are the surface for wavelenghts λ ? 1 cm, atmospheric N₂ for 1 cm ? λ ? 200 μm,, condensed and gaseous CH₄ for 200 μm ? λ ? 20 μm, and molecular bands and organic haze for λ ? 20 μm. Matching computed spectra to the observed Voyager IRIS spectra at 7.3 and 52.7° emission angles yields the following abundances and locations of opacity sources: CH₄ clouds: 0.1 g cm^? at a planetocentric radius of 2610–2625 km, 0.3 g cm^?2 at 2590–2610 km, total 0.4 ± 0.1 g cm^–2 above 2590 km; organic haze: $\text{4 ± 2 × 10}{}^{\mathrm{?6}}\text{,}\text{g cm}\text{,}{}^{\mathrm{?2}}$ above 2750 km; tropospheric H₂: 0.3 ± 0.1 mol%. This is the first quantitative estimate of the column density of condensed methane (or CH₄/C₂H₆) on Titan. Maximum transparency in the middle to far IR occurs at 19 μm where the atmospheric vertical absorption optical depth is $\text{?0.6}$ A particle radius $\text{r}\text{? 2 μ}\text{m}$ in the upper portion of the CH₄ cloud is indicated by the apparent absence of scattering effects.     

The 63 μm radiation field in the earth's thermosphere and its influence on the atomic hydrogen temperature

The thermal escape of hydrogen from the Earth's atmosphere is strongly affected by its temperature at the exobase. It has been suggested recently that the hydrogen temperature might be significantly lower than the thermospheric temperature as a result of a collisional exchange of energy with atomic oxygen. The tendency is to cool the hydrogen since the energy of the excited ${}^3\text{P}{}_1$ level of oxygen can be lost from the atmosphere via magnetic dipole emission of the 63 μm line (${}^3\text{P}{}_2\text{?}{}^3\text{P}{}_1$). We present here a detailed calculation of the net cooling effect as a function of altitude throughout the thermosphere. The calculations have been performed for both day and night conditions and for periods of maximum and minimum solar activity conditions. It is found that its effect on ΔT/T varies from a very small value to a maximum of ～3%. We also provide the theoretical framework for describing deviations of the 63 μm emission from local thermodynamic equilibrium and show that these effects can cause the emission to be reduced by as much as 40% near 500 km.     

Electric fields and potential patterns in the high-latitude ionosphere for different situations in interplanetary space

The potential ? of the electric field at high latitudes has been obtained by solving numerically the second order differential equation in spherical coordinates:

$\text{?}\text{1}\text{2}\text{(rσ}{}^{\mathrm{H}}\text{?}{}_{\mathrm{\theta }}\text{)}{}_{\mathrm{\theta }}\text{+}\text{1}\text{r}\text{(σ}{}^{\mathrm{H}}\text{?λ)}{}_{\mathrm{\lambda }}\text{+}\text{1}\text{r}\text{(σ}{}^{\mathrm{P}}\text{?}{}_{\mathrm{\lambda }}\text{)}{}_{\mathrm{\theta }}\text{?(σ}{}^{\mathrm{P}}\text{?}{}_{\mathrm{\theta }}\text{)}{}_{\mathrm{\lambda }}\text{=}\text{1}\text{r}\text{(rψ}{}_{\mathrm{\theta }}\text{)}{}_{\mathrm{\theta }}\text{+}\text{1}\text{r}{}^2\text{ψ}{}_{\mathrm{\lambda \lambda }}$

, where θ is colatitude, λ is longitude, σ^H and σ^P are the height-integrated Hall and Perdersen ionospheric conductivities, r = sinθ, and ψ is the current function. The boundary condition is ? = 0 on the geomagnetic parallel θ = 34°. Values of ψ are determined from geomagnetic field variations at the Earth's surface from geomagnetic field variations at the Earth's surface for various conditions in interplanetary space. σ^P and σ^H are taken to vary with season, local time, tilt of the geomagnetic dipole axis (UT), and intensity of corpuscular precipitation (the model proposed by Wallis and Budzinski, 1981). The model distributions of ?^M and E^M = -▽?^m so obtained are compared with observational results. The feasibility has been demonstrated of interpreting the statistical results and individual measurement data in terms of a unified dynamic model of ionospheric electric fields. The model makes allowance for the changes of electromagnetic “weather” in interplanetary space.     

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.     

Cross-section for the dissociative photoionization of hydrogen by 584 Å radiation: The formation of protons in the Jovian ionosphere

The cross-section for dissociative photoionization of hydrogen by 584 Å radiation has been measured, yielding a value of 5 × 10^?20 cm². The process can be explained as a transition from the $\text{X}{}^1\text{Σ}{}_{\mathrm{g}}{}^{\mathrm{+}}$ ground state to a continuum level of the $\text{X}{}^2\text{Σ}{}_{\mathrm{g}}{}^{\mathrm{+}}$ ionized state of H₂ The branching ratio for proton (H⁺) vs molecular ion (H₂⁺) production at this energy is 8 × 10^?3. This process is quite likely an important source of protons in the Jovian ionosphere near altitudes where peak ionization rates are found.     

On the photodissociation of water vapour in the mesosphere

The photodissociation of water vapour in the mesosphere depends on the absorption of solar radiation in the region (175–200 nm) of the O₂ Schumann-Runge band system and also at H-Lyman alpha. The photodissociation products are OH + H, OH + H, O + 2H and H₂ + O at Lyman alpha; the percentages for these four channels are 70, 8, 12 and 10%, respectively, but OH + H is the only channel between 175 and 200 nm. Such proportions lead to a production of H atoms corresponding to practically the total photodissociation of H₂O, while the production of H₂ molecules is only 10% of the H₂O photodissociation by Lyman alpha.The photodissociation frequency (s^?1) at Lyman alpha can be expressed by a simple formula

$\text{J}{}_{\mathrm{<mtext>Ly</mtext><mtext>\alpha </mtext>}}\text{H}{}_2\text{O=4.5 ×10}{}^{\mathrm{?6}}\text{1+0.2}\text{F}{}_{10.7}\text{?65}\text{100}\text{exp}\text{[?4.4 ×10}{}^{\mathrm{?19}}\text{N}{}^{0.917}\text{]}$

where F_{10.7 cm} is the solar radioflux at 10.7 cm and N the total number of O₂ molecules (cm^?2), and when the following conventional value is accepted for the Lyman alpha solar irradiance at the top of the Earth's atmosphere ($\text{Δλ = 3.5}\text{A}\text{?}$) q_Lyα,∞ = 3 × 10¹¹ photons cm^?2 s^1?.The photodissociation frequency for the Schumann-Runge band region is also given for mesospheric conditions by a simple formula

$\text{J}{}_{\mathrm{<mtext>SRB</mtext>}}\text{(}\text{H}{}_2\text{O}\text{) = J}{}_{\mathrm{<mtext>SRB</mtext><mtext>,\infty </mtext>}}\text{(}\text{H}{}_2\text{O}\text{) exp [?10}{}^{\mathrm{?7}}\text{N}{}^{0.35}\text{]}$

where J_SRB,∞(H₂O) = 1.2 × 10^?6 and 1.4 × 10^?6 s^?1 for quiet and active sun conditions, respectively.The precision of both formulae is good, with an uncertainty less than 10%, but their accuracy depends on the accuracy of observational and experimental parameters such as the absolute solar irradiances, the variable transmittance of O₂ and the H₂O effective absorption cross sections. The various uncertainties are discussed. As an example, the absolute values deduced from the above formulae could be decreased by about 25-20% if the possible minimum values of the solar irradiances were used.     

Changes in the concentration of mesospheric ozone during the total solar eclipse

Measurements of dayglow radiance of $\text{O}{}_2\text{(}{}^1\text{Δ}{}_{\mathrm{g}}\text{)}$ and OH(7,2) bands are reported. Ground based photometers were used to monitor zenith radiance of 1270 and 694 nm emissions during the total solar eclipse of 16 February 1980. Altitude distribution of 1270 nm intensity was derived from ground based observations. A set of altitude distributions of $\text{O}{}_2\text{(}{}^1\text{Δ}{}_{\mathrm{g}}\text{)}$ were thus obtained throughout the eclipse. These altitude distributions were converted into ozone distributions using the rate equations for formation and loss of ozone and $\text{O}{}_2\text{(}{}^1\text{Δ}{}_{\mathrm{g}}\text{)}$ molecules. Results indicate an increase in the ozone concentration at mid-eclipse. OH(7,2) emission did not show enhancement during totality. This may mean that there was no increase in OH concentration during the eclipse.