The Venus ultraviolet oxygen dayglow and aurora: Model comparison with observations

ARTICLE IN PRESS

Planetary and Space Science 56 (2008) 542–552 www.elsevier.com/locate/pss

The Venus ultraviolet oxygen dayglow and aurora: Model comparison with observations J.-C. Ge´rarda,, B. Huberta, V.I. Shematovichb, D.V. Bisikalob, G.R. Gladstonec a

Laboratoire de Physique Atmosphe´rique et Plane´taire, Universite´ de Lie`ge, Belgium Institute of Astronomy, Russian Academy of Sciences, Moscow, Russian Federation c Southwest Research Institute, San Antonio, TX, USA

b

Received 14 May 2007; received in revised form 15 November 2007; accepted 19 November 2007 Available online 5 February 2008

Abstract We compare the intensity of the OI 130.4 and 135.6 nm emissions calculated using the soft electron precipitation measured on board the Pioneer Venus (PV) Orbiter with the auroral brightness observed with the ultraviolet spectrometer (OUVS) on board the PV. For this purpose, we use a new electron transport model based on a Monte Carlo implementation of the Boltzmann equation and a multi-stream radiative transfer model to calculate the effects of multiple scattering on the intensity field of the 130.4-nm triplet. We show that the consideration of the enhancement of the emergent 130.4-nm to the 135.6-nm intensity by multiple scattering in the optically thick Venus atmosphere increases the auroral 130.4/135.6 ratio by a factor of about 3. We find agreement with the mean 130.4/135.6 ratio observed with PV-OUVS using the typical suprathermal electron energy spectrum reported from PV in situ measurements showing a characteristic energy of about 14 eV. To account for the average OI auroral emissions, the required precipitated energy flux is 2  103 mW m2, that is about 30% of the measured suprathermal night-side soft electron spectrum used as a reference. The calculated brightness of the CO Cameron bands is about twice as large as the weak observed emission, but within the error bars of the observations and the uncertainties of the dissociative excitation cross-section of CO2. The electron transport model, coupled with calculations of excitation processes is also applied to an analysis of the FUV oxygen day airglow observations made with PV-OUVS and the Hopkins Ultraviolet Telescope (HUT) spectrograph. Comparisons indicate that the model accounts for both the disc-averaged intensities observed with the HUT spectrograph, the limb scans and the 130.4-nm images obtained with PV-OUVS. The relative contribution of resonance scattering of the solar line and photoelectron impact to the excitation of the 130.4-nm triplet depends on the altitude, but is globally dominated by resonance scattering. The intensity of the 130.4-nm dayglow emission does not vary proportionally with the O density in the lower thermosphere, but provides nevertheless a useful tool to remotely probe the atomic oxygen density and its variations. r 2007 Elsevier Ltd. All rights reserved. Keywords: Venus thermosphere; Aurora; Electron precipitation; Ultraviolet airglow; Atomic oxygen

1. Introduction Ultraviolet dayglow emissions of Venus are produced by the direct or indirect interaction of solar ultraviolet radiation and photoelectrons with neutral constituents. The processes involved include photoionisation and excitation, excitation by photoelectron impact, electron impact ionisation, dissociative excitation and resonance scattering. Variable auroral emission has also been Corresponding author. Tel.: +32 43669774; fax: +32 43669711.

E-mail address: [email protected] (J.-C. Ge´rard). 0032-0633/$ - see front matter r 2007 Elsevier Ltd. All rights reserved. doi:10.1016/j.pss.2007.11.008

observed on the night side, suggesting that the atmosphere there is subject to a flux of energetic but soft electrons. 1.1. The dayglow Among commonly observed dayglow emissions, the OI emissions at 130.4 and 135.6 nm are particularly important, as their production rates are dominated by processes directly involving ground state O atoms. Observations and models of the ultraviolet dayglow of Venus have been reviewed by Paxton and Anderson (1992) and Fox and Bougher (1991). Observations were obtained from spectro-

ARTICLE IN PRESS J.-C. Ge´rard et al. / Planetary and Space Science 56 (2008) 542–552

meters and spectrographs on board Mariner 10 (Broadfoot et al., 1974) and Venera 11 and 12 (Bertaux et al., 1981). The Pioneer Venus Orbiter (PVO) also carried a UV spectrometer (PV-OUVS) with modest spectral resolution that was also able to provide spin-scan monochromatic images of the airglow distribution (Stewart, 1980). In 1990, Galileo flew by Venus and observations from the UV spectrometers were reported by Hord et al. (1991). In March 1995, FUV spectra of the Venus disc were obtained between 82 and 184 nm at 0.4 nm resolution with the Hopkins Ultraviolet Telescope (HUT) on board the Space Shuttle by Feldman et al. (2000). Emissions from HI, OI, OII, CI and CO were observed and compared with earlier observations. PV-OUVS limb scans at 130.4 nm were modelled by Paxton and Meier (1986) who found that their model could adequately fit the observations using O densities from the VIRA model (Hedin et al., 1983), based on in situ observations. They concluded that remote sensing of this emission is potentially able to provide the global distribution of atomic oxygen. Monochromatic spin-scan images of the 130.4-nm dayglow emission collected by the PV near apoapsis were analysed by Alexander et al. (1993). They found a local time asymmetry, which they interpreted as an enhancement of the O density by a factor of 2 near the evening terminator, poleward of 301. Atomic oxygen becomes the main constituent 15–20 km above the homopause located near 135 km, where molecular diffusive processes control vertical transport. Carbon dioxide, which dominates the lower atmosphere, absorbs the OI 130.4-nm emission below 125–130 km. Sources for O atoms in the Venus thermosphere are photodissociation of CO2 and CO and, to a minor extent, dissociative recombination in the ionosphere. The OI 130.4-nm dayglow emission is excited by resonance scattering of a fraction of the photons emitted in the broad solar emission line and by photoelectron impact on ground state 3P atomic oxygen. Therefore, the 130.4-nm intensity partly reflects the distribution of O atoms above the altitude of CO2 absorption. However, since the O3P–3S1 transition is optically thick, multiple scattering complicates the interpretation of the measured brightness in terms of oxygen density. The excitation of the 135.6-nm emission is limited to electron impact since the transition from the 3P ground state of O to the 5S1 state is spin-forbidden. The fourth positive (4P) system of CO is the most intense band system in the 120–180 nm region of the dayglow spectrum of Venus. It arises from the dipole allowed A1P–X1S transition. Potential sources include photodissociative excitation and electron impact dissociation of CO2, electron impact excitation of CO, dissociative recombination of CO+ and fluorescent scattering by 2 ground-state CO (Fox and Dalgarno, 1981). The Cameron band system of CO dominates the 180–260 nm region of the ultraviolet spectrum and corresponds to the a3P–X1S dipole forbidden transition. The CO (a3P) upper state is thus metastable, with a relatively short radiative lifetime of

543

8 ms. The potential sources are the same as those for the 4P band system, except for fluorescent scattering. The ultraviolet doublet at 288.3 and 289.6 nm is produced by the (0, 0, 0)–(0, 0, 0) transition between the B2S state and the X2P ground state of CO+ 2 . This band system is produced by photoionisation and electron impact ionisation of CO2 and by fluorescent scattering of solar ultraviolet radiation. Carbon lines have also been observed (Paxton and Anderson, 1992) but will not be discussed in this study. 1.2. The aurora Highly variable UV emissions have been observed over an 8-year period with the Pioneer Venus Ultraviolet Spectrometer (PV-OUVS) on the night-side of Venus (Philipps et al., 1986; Fox and Stewart, 1991). They appeared mostly as bright spots of 130.4-nm emission, although OI 135.6 nm and CO Cameron bands were also observed. The 130.4-nm intensity varied from levels near the instrumental background (4 R) to 20 R, with occasional peaks up to 100 R. The mean maximum brightness level was close to 50 R, while the night-side global mean intensity was 10 R at solar maximum, decreasing to 4 R at solar minimum. The dependence on the emission angle did not follow a cos-1 law and indicated that the emission was optically thick. A dawn–dusk brightness asymmetry was observed, with the most intense emissions in the evening sector. No correlation was found between the brightness and the solar wind parameters, indicating that the solar wind does not directly modulate the aurora. However, intensity enhancements preceding the arrival of interplanetary shocks at Venus have been observed (Philipps et al., 1986). The 135.6-nm emission is weaker with a disc-averaged vertical brightness 1.6 R at solar maximum and 0.8 R at solar minimum. Its intensity increases towards the limb as expected for an optically thin emission. Finally, the CO Cameron bands were also detected with an estimated brightness of 25750 R at the nadir. Suprathermal electrons have been detected above the atmosphere in the Venus umbra by the PV retarding Potential Analyzer (RPA) (Knudsen and Miller, 1985) between 1000 and 2500 km, and between 1500 and 2000 km by the detectors on board Venera 8 and 9 (Gringauz et al., 1979). These suprathermal electrons fall into two groups. The first one (less energetic) is similar to the photoelectron spectra observed in the dayside ionosphere This population has been interpreted as belonging to plasma transported from the day to the night-side ionosphere. It is generally well-fitted by a Maxwellian energy distribution with a characteristic energy E0 of about 7 eV. The second (more energetic) group presents similarities with the population observed in the upper mantle and solar wind near the terminator. It is also well-fitted by a Maxwellian distribution, but the characteristic energy is close to 14 eV. The two components are sometimes mixed, but regions exist where

ARTICLE IN PRESS 544

J.-C. Ge´rard et al. / Planetary and Space Science 56 (2008) 542–552

one or the other component is dominant. The relative importance of the two populations appears to be dependant on temporal and local conditions. A more detailed analysis by Spenner et al. (1996) showed that a suprathermal flux of about 12-eV electrons seems to be present at all altitudes down to at least 200 km. This harder population is believed to be associated with solar wind plasma making its way into the umbra of Venus, precipitating into the atmosphere and possibly causing the observed UV emission enhancements. Its origin is possibly connected to night-side ionospheric holes frequently observed which are threaded by an enhanced, near-vertical magnetic field penetrating into the ionosphere and presumably extending into the solar wind wake. Spenner et al. (1996) suggested that solar wind electrons may enter the ionosphere along the channels and diffuse outward into the surrounding ionosphere. Fox and Stewart (1991) compared the intensity of the auroral emissions observed by PV-OUVS to that expected from the precipitation of the suprathermal electrons with the energy spectrum measured with the RPA on the night side. They used the downward travelling portion of the omni-directional flux of 14-eV electrons, carrying an energy flux of 6  103 mW m2 s1, reported by Knudsen and Miller (1985) to estimate the column excitation rates of O (3S), O (5S) and CO (a3P). Their discrete local loss model produced intensities of 6.4 R for 130.4 nm, 1.6 R for 135.6 nm and 18 R for CO Cameron if the ‘‘reference’’ spectrum of suprathermal electrons was multiplied by 0.08. They calculated a 130.4/135.6 nm ratio of 4, to be compared with the observed ratio ranging between 5.572.7. However, multiple scattering of the 130.4-nm photons by the optically thick Venus atmosphere was not accounted for in their model, even though the upward 130.4-nm intensity under these conditions may be enhanced by a factor of more than 2. The observed weakness of the Cameron emission is also an indicator that the precipitation is soft and does not reach regions dominated by CO2. In this study, we apply a recently developed model including electron transport in the Venus thermosphere using a Monte Carlo approach. This transport model is described by Shematovich et al. (2007) and validated by comparing calculated CO Cameron and CO+ 2 doublet intensities in the dayglow of Mars with observations made with the SPICAM instrument on board Mars Express. It is, together with calculations of other processes and treatment of multiple scattering, applied here to the analysis of existing observations of the Venus aurora and dayglow with emphasis on the oxygen emissions. 2. The model Excitation processes leading to the production of dayglow and auroral emissions involve the interaction of solar ultraviolet radiation and energetic electrons with major neutral atmospheric species. Therefore, modelling their

production rate requires a detailed treatment of the absorption of solar UV photons as they travel through the CO2-dominated atmosphere, dissociation and ionisation of the constituents. Since photoelectrons produced in the ionisation processes may carry enough energy to excite optical transitions, collisions leading to energy degradation of these energetic electrons must be properly accounted for in order to accurately calculate volume excitation rates. In the case of auroral precipitation, only direct excitation by inelastic collisions with the main constituents is involved. In this section, we first describe the model atmosphere used in this study and the calculation of the steady state energetic electron energy distribution. We then review the mechanisms producing the OI 130.4 and 135.6 nm and the CO Cameron band emissions. Finally, we describe the role of multiple scattering in the distribution of the 130.4-nm radiation field. 2.1. Neutral atmosphere, cross-sections and solar flux Information about the composition and structure of the thermosphere of Venus, obtained largely from in situ measurements, has been summarised in several empirical models. In this study, we use the Pioneer Venus Thermospheric Model developed by Hedin et al. (1983), with its extension below 140 km by the Venus International Reference Atmosphere (VIRA, 1985) model. The model provides neutral densities (CO2, O, CO, N2, He and N) and temperature in the Venus thermosphere. It is based on PVO-data from the Orbiter Neutral Mass Spectrometer (ONMS) and on some density data from the entry probe. The model formulation relies on modified Bates temperature profiles and the related diffusive equilibrium density profiles. To calculate the photoionisation rates, we use the SOLAR2000 EUV solar spectrum model (Tobiska, 2004), which provides fluxes in 39 wavelength bins and emission lines between 1.86 and 105.0 nm. This code is available from the web site /http://www.spacewx.com/Products. html#s2k_rg_gradeS. The photoionisation and absorption cross-section data and the branching ratios for CO2, CO, O, N2 were taken from Huebner et al. (1992). 2.2. Electron production, transport and thermalisation The numerical model used to calculate the photoelectron production and energy degradation in the Venus atmosphere has been described by Shematovich et al. (2007) and will be only briefly summarised here. The same model is applied to soft electron auroral precipitation by replacing the solar-induced ionisation by the production of secondary auroral electrons following inelastic collisions. In the daytime thermosphere of Venus, energetic electrons are produced by photoionisation of the main atmospheric constituents by EUV and X-ray solar radiation. These newly formed electrons are transported in the thermosphere where they lose their kinetic energy in elastic, inelastic and ionisation collisions with the ambient

ARTICLE IN PRESS J.-C. Ge´rard et al. / Planetary and Space Science 56 (2008) 542–552

atmospheric gas. 8 0 > < eðE Þ þ X  0 eðEÞ þ X ! eðE Þ þ X > : eðE 0 Þ þ X þ þ eðE Þ s

(1)

where E and E0 (oE) are the kinetic energies of the primary electron before and after a collision, X stands for CO2, CO, O or N2, X* and X+ are atmospheric species in excited and ionised states and Es is the energy of the secondary electron formed in the ionising collision. If the collision produces ionisation, a secondary electron is created and is randomly assigned an isotropically distributed pitch angle and an energy in agreement with the procedure given by Green and Sawada (1972) and Jackman et al. (1977). The list of cross-sections used to calculate the energy loss associated with inelastic collisions with CO2, CO and O is given by Shematovich et al. (2007). The following expression is used to calculate the photoelectron production rate, Pe(E, z): Z 1 XX Pe ðE; zÞ ¼ nk ðzÞ dl I 1 ðlÞ expðtðl; zÞÞsik pk ðl; E l Þ, k

l

0

(2) where the slant optical depth t is given by R1 X nk ðz0 Þdz0 a tðl; zÞ ¼ sk ðlÞ z cos y k and nk is the number density of the kth neutral constituent, sik ðlÞ and sak ðlÞ are the wavelength-dependant total ionisation and absorption cross-section, respectively, pk(l, El) the branching ratio for the excited ion state with and an ionisation energy El, E ¼ ElEl, El the energy corresponding to wavelength l, lk the ionisation threshold wavelength for neutral species k, IN(l) the incident solar radiation flux at wavelength l and y is the solar zenith angle (SZA). The newly created electrons lose their excess kinetic energy in collisions with the atmospheric particles. Their kinetics and transport is described by the kinetic Boltzmann equation. The Direct Simulation Monte Carlo (DSMC) method is used to solve atmospheric kinetic systems in the stochastic approximation (Shematovich et al., 1994, 2007; Bisikalo et al., 1995; Ge´rard et al.; 2000). The evolution of the system of modelling particles due to collisional processes and particle transport is calculated from the initial to the steady state. The lower boundary is set at an altitude 100 km and the upper boundary is fixed at 250 km where the atmospheric gas is practically collisionless. The region of the atmosphere under study is divided into 49 vertical cells. The production rates of states excited by electron impact are easily calculated using the calculated energy distribution function (EDF), the neutral density distribution and the relevant excitation cross-sections. We now illustrate results of calculations of EDF obtained with the Monte Carlo code by solving the Boltzmann equation for photoelectrons. The EDF calculated at 139 km including both the degraded primary

545

photoelectrons and the secondary component is shown in Fig. 1. The thermal core dominates below 1–2 eV, a dip is observed between 2 and 4 eV as a consequence of energy loss by excitation of the CO2 vibrational levels. Finally, the higher-energy tail is produced by photoionisation processes and the various peaks correspond to emission features in the solar spectrum. Suprathermal electron fluxes with energies up to 50 eV were measured in the Venus dayside ionosphere, and their dependence on altitude, SZA and solar F10.7 flux was analysed by Spenner et al. (1997). It was found that the average flux is nearly constant for SZAs between 01 and 701, and decreased with increasing altitude. The flux increased with the F10.7 index. To make an accurate comparison, we use the measured flux for F10.7 ¼ 180, SZA ¼ 201, at an altitude of 250 km. The measured upgoing flux is compared with the calculated upward flux in Fig. 2. It is seen that the calculated and measured fluxes are in good agreement with the observations both in shape and magnitude. The calculated electron flux shows a shape similar to the EDF presented in Fig. 1, with the same characteristic features such as the thermal core, a bump due to the secondary electrons and the structured tail distribution of photoelectrons. We note that these calculations were performed using a ‘‘transparent’’ upper boundary condition assuming that no reflection of the upward electron flux occurs in the ionosphere. However, there are some indications (see, e.g. Spenner et al., 1997) that the Venus ionopause can reflect the upward flux. We made runs with a ‘‘reflecting’’ boundary condition and found essentially an increase of the downward flux with little effect on the upward flux values.

Fig. 1. Energy distribution function of photoelectrons at 139 km, local noon calculated for maximum solar activity conditions (F10.7 ¼ 180).

ARTICLE IN PRESS 546

J.-C. Ge´rard et al. / Planetary and Space Science 56 (2008) 542–552

distribution of these emissions for various characteristic electron energies and determine the emergent intensity, including the effect of radiative transfer on the 130.4-nm triplet emission. To calculate the energy distribution of precipitating auroral electrons in the Venus thermosphere, we use the model where the photoelectron source is replaced by an initial auroral energy distribution at the top of the atmospheric region under study. The initial auroral electron beam degrades by collisions and produces secondary electrons. Other loss and excitation processes are treated identically for photoelectrons and secondary auroral electrons. 3.1. Auroral precipitation

Fig. 2. Measured (triangles) and calculated (solid line) upward fluxes of photoelectrons on Venus day side at 250 km for solar maximum conditions (F10.7 ¼ 180).

2.3. Multiple scattering To calculate the effects of multiple scattering on the 130.4-nm triplet radiation field and the emerging intensity, we use a resonance line radiative transfer code (Gladstone, 1985). The process of frequency redistribution allows photons to escape an optically thick atmosphere by scattering in frequency from the core of the line into the optically thin line wings. In these conditions, the emergent intensity can be up to five times as large as the intensity obtained by integrating the production rate along a line of sight. In this study, we use angle-averaged partial frequency redistribution. In the case of the Venus dayglow and aurora, the 130.4-nm photons may be absorbed by CO2 at altitudes below the peak of the O (3S1) emission. For both the aurora and the airglow, the optical depth for 130.4-nm absorption by CO2 reaches unity near 130 km. In the terrestrial atmosphere, the 135.6-nm emission is marginally optically thick. In the Venus thermosphere, the vertical thickness at the line centre of the strongest component is t0.04 at 128 km, the altitude of unit optical depth for CO2 absorption. Therefore, we neglect multiple scattering of the O 3P–5S1 transition in this study.

The precipitation of energetic electrons on the Venus night side is sporadic and variable as described by Spenner et al. (1996). To simulate the interaction of the electron beam with the atmosphere, we use an isotropic Maxwellian distribution with a characteristic energy (equals half of the mean energy) of 14 eV and a unit isotropic energy flux (1 mW m2) at the top of the region under study. As an example, Fig. 3 shows the upward and downward auroral fluxes calculated at 139 km, close to the altitude of the maximum energy deposition. The scatter at energies above approximately 20 eV is caused by stochastic numerical noise in the solution of the Boltzmann equation. The characteristics of the energy spectra of the auroral electrons are very similar to those of the photoelectrons shown in Fig. 2.

3. The Venus ultraviolet aurora As described before, the 130.4, 135.6 nm and, to some extent, CO Cameron band emissions have been sporadically observed with PV-OUVS on the Venus night side. We use the electron transport model to calculate the vertical

Fig. 3. Downward and upward fluxes of auroral electrons at 139 km for an incident Maxwellian precipitation with a characteristic energy E0 ¼ 14 eV carrying a total energy flux of 1 mW m2 s1 at solar maximum conditions (F10.7 ¼ 200).

ARTICLE IN PRESS J.-C. Ge´rard et al. / Planetary and Space Science 56 (2008) 542–552

3.2. Auroral emissions To investigate the FUV emissions associated with the precipitation of soft auroral electrons, the column emission rates associated with three different electron characteristic energies have been calculated with the Monte Carlo model described previously. A unit isotropic energy flux is precipitated through the upper boundary and the primary electron energy degradation is followed in parallel with the formation of secondary electrons following ionising collisions with the ambient neutrals. Table 1 lists the four cases discussed as follows. All runs were made for local midnight at the equator. Fig. 4 shows the volume emission rate of O (3S1), O (5S1) and CO (a3P) states for an auroral precipitation characterised by a Maxwellian distribution with a characteristic energy of 14 eV and an energy flux of 1 mW m2. These states are the upper levels of the OI multiplets at 130.4 and 135.6 nm and of the CO Cameron bands, respectively. The 3S1 and 5S1 states are produced by electron impact on O, CO2 and CO. According to Fox and Dalgarno (1981), the contributions to these atomic oxygen emissions from dissociative excitation of CO2 and CO are approximately one order of magnitude less than electron impact on atomic oxygen. Therefore, in this model, we

Table 1 Characteristics of auroral Maxwellian precipitation used in the model calculations Case Case Case Case

E0 ¼ 14 eV, F10.7 ¼ 180, UT ¼ 00 LT E0 ¼ 7 eV, F10.7 ¼ 180, UT ¼ 00 LT E0 ¼ 21 eV, F10.7 ¼ 180, UT ¼ 00 LT E0 ¼ 14 eV, F10.7 ¼ 80, UT ¼ 00 LT

1 2 3 4

260 240

Altitude (km)

220 200

130.4

135.6 CAMCO2

180 CAMCO 160 140 120

10-2

100 102 -3 -1 Emission rate (cm s )

104

Fig. 4. Calculated volume emission rate of OI and CO Cameron emissions for an incident Maxwellian auroral precipitation with the same characteristics as in Fig. 3.

547

consider the contribution of electron impact on atomic oxygen only. The excitation cross-sections for electron impact on O are taken from the recent critical review by Johnson et al. (2005). The cross-section for the 135.6-nm emission is based on the value determined by Stone and Zipf (1974), reduced by a factor of 2.8, following the recommendation of Zipf and Erdman (1985). It reaches a value of 1.2  1017 cm2 at 16 eV. The 130.4-nm emission crosssection given by Zipf and Erdman (1985) has been adopted for electron impact excitation on atomic oxygen. The peak value is equal to 1.7  1017 cm2 at 20 eV. We note that uncertainties exist concerning the importance of cascades from upper level of the O atoms such as the 844.6 nm and the 102.7 nm transitions which feed the 3S1 state. The potential importance of these cascade processes was discussed in detail for the Earth’s atmosphere by Meier (1991) and by Meier et al. (1983) for the Venus case. The Cameron bands are produced by electron impact on CO2 and CO: CO2 þ e ! CO ða3 PÞ þ O þ e CO þ e ! CO ða3 PÞ þ e To calculate the dissociative excitation of Cameron bands, we use the semi-empirical representation by Sawada et al. (1972) scaled to the measurements by Ajello (1971) as has been frequently used in earlier work. However, the CO (a3P) upper state of the Cameron bands is metastable with a relatively long radiative lifetime and is both directly excited by electron impact and indirectly by internal cascade. Furthermore, being a dissociative fragment of CO2, CO (a3P) molecules are formed with an excess kinetic energy and such excited molecules can escape from the excitation region. Accordingly, the emission of the Cameron system is very weak and easily blended with other emissions, unless the excitation energy is small (Furlong and Newell, 1996). Following the discussions by Itikawa (2002) and Furlong and Newell (1996), we rescaled the Cameron system emission cross-section to the peak value of 2.4  1016 cm2 at 80 eV, based on the measurements by Erdman and Zipf (1983). We note that this value is more than three times larger than the value of 7  107 cm2 deduced by Conway (1981) from his analysis of the Martian dayglow. Because of the above-mentioned experimental difficulties, the peak value is likely to have a large uncertainty of at least a factor of 2. The 130.4- and 135.6-nm oxygen emissions peak at a slightly higher altitude than the Cameron bands, reflecting the changing chemical composition with altitude in the atmosphere of Venus. Table 2 lists the oxygen emission rate (in R) observed with the PV-OUVS instrument (Fox and Stewart, 1991) and the column excitation rates (in kR) calculated for the four cases listed in Table 1. The calculated 130.4/135.6 column production ratio varies between 1.8 and 2.2; a range of values to be compared with the observed ratio of about 4. However, the calculated

ARTICLE IN PRESS J.-C. Ge´rard et al. / Planetary and Space Science 56 (2008) 542–552

548

Table 2 Observed disc-averaged nadir auroral brightness at solar maximum (in R) (Fox and Stewart, 1991) and auroral emission rate (in kR) calculated for a 1 mW m2 energy flux in cases listed in Table 1

OI 130.4 column OI 130.4 nm (RT) OI 135.6 nm CO Cameron bands 130.4/135.6 (column) 130.4/135.6 (RT)

10 1.8 25750 5.572.7

Case 1 (kR)

Case 2 (kR)

Case 3 (kR)

Case 4 (kR)

1.7 5.3 0.8 5.5 2.1 6.5

2.1 6.5 1.2 6.0 1.8 5.5

1.4 4.3 0.6 5.4 2.2 6.8

1.8 5.5 0.8 5.5 2.1 6.6

emerging intensity ratio is increased by multiple scattering of the 130.4-nm photons in the optically thick Venus atmosphere and ranges between 5.5 and 6.8, as also shown in Table 2. Part of the downward flux of photons will ultimately be absorbed by CO2 at the bottom of or below the emission layer. Consequently, observations of the 130.4/135.6 ratio can only be used to determine the characteristic energy of the auroral electrons if the effect of multiple scattering in the O (3P–3S) transition is accounted for. The role of multiple scattering is illustrated in Fig. 5 comparing the volume excitation rate of O (3S) atoms with the source function in the radiative transport equation. As already outlined before, the 130.4/135.6 calculated ratio varies between 1.8 and 2.2 for the column emission rate, and between 5.5 and 6.8 for the nadir intensity, as a result of multiple scattering through the Venus atmosphere. It is seen that, although both exhibit a peak at about the same altitude, the final source function is considerably larger than the primary excitation rate as a consequence of multiple scattering of the 130.4-nm photons. Table 2 also lists the emerging 130.4-nm intensity for a nadir observation. The altitude of the maximum excitation rate is 13971.5 km and absorption of photon by CO2 in the emission column is negligible above 130 km. The observed 130.4/135.6 ratio of 5.572.7 is larger than the values (2) predicted for a 130.4 nm optically thin emission but quite close to those (5.5–6.8) obtained when multiple scattering is included. We note, however, that the observed variability is large. From the comparison of the 130.4-nm intensities listed in columns 1 and 2 (case 1), we conclude that if the electron spectrum has a characteristic energy of 14 eV, the average precipitated auroral energy flux is about 2  103 mW m2, that is about 30% of the flux carried by the suprathermal energy spectrum reported by Knudsen and Miller (1985). This value is somewhat larger than the 8–28% deduced by Fox and Stewart (1991) from comparison with their model. The presence of the CO Cameron bands was also marginally detected with PV-OUVS, with an average nadir brightness on the order of 25 R and a large range of variability and uncertainty. For Case 1, the calculated Cameron emission rate is nearly equal to the 130.4-nm triplet brightness after accounting for multiple scattering, to be compared with the observed ratio of about 2.5. The

130.4 nm

220 RT Source Function 200 Altitude (km)

Observations (R)

180

160

140

Primary Source

120 100

101

102 103 104 -3 Volume emission rate (cm s-1)

105

Fig. 5. Dotted line—primary source of 130.4-nm photons produced by auroral electron impact on OI (E0 ¼ 14 eV). Solid line—final source accounting for multiple scattering used in the formal solution of the radiative transfer equation. The large amplification of the primary source function is a consequence of the large optical thickness of the OI 3P–3S triplet.

observed Cameron/135.6-nm ratio is 14, while the calculated ratio is 8.3 for Case 1 and 10.3 for Case 3. The predicted CO Cameron band intensity for Case 1 and a flux of 2.3  103 mW m2 is 13 R, that is less by a factor of 2 than the reported emission rate. However, we note that an accurate evaluation of a molecular emission spread over a wide wavelength range is more inaccurate than for an atomic line. Although puzzling, this discrepancy may also reflect uncertainties on the excitation cross-section of the CO (a3P) state by electron impact. In summary, our study confirms earlier models indicating that soft electron precipitation, probably associated with solar wind electron precipitation in the planet’s umbra, produces the weak sporadic auroral emissions detected by the PV. Consideration of multiple scattering of the 130.4-nm emission confirms that the optical thickness of the O3P–3S1 transition enhances the 130.4/ 135.6 ratio up to 5.5–6.8 for characteristic electron energies of 7–21 eV.

ARTICLE IN PRESS J.-C. Ge´rard et al. / Planetary and Space Science 56 (2008) 542–552

4. Airglow emissions

350 130.4 nm

300

Altitude (km)

The volume excitation rates of the dayglow emissions by photoelectron impact may also be calculated using the calculated EDFs described in Section 2, number density profiles and relevant excitation cross-sections. We can then examine the model predictions for the oxygen FUV dayglow, with particular emphasis on the 130.4-nm emission. As for the auroral case, multiple scattering of the 130.4-nm photons redistributes the primary source of photons and alters the emergent intensity field. We first illustrate profiles for the source of O (3S1) excitation and compare with angular scans from PV-OUVS. We then compare the calculated emergent intensity of some FUV dayglow features with observations from spectrometers on board Venera, HUT and the PV. In addition to photoelectron impact, the 3S1 state is also excited by resonance scattering of solar 130.4-nm photons:

549

250

200

150

O ð3 PÞ þ hn130:4 ! O ð33 S Þ. 100 1

10

100

Primary source

1000

(cm-3 s-1)

Fig. 6. Volume excitation rate of O (3S1) atoms in the dayglow for solar maximum conditions (F10.7 ¼ 240, F 10:7A ¼ 177) (see text). Solid line—the photoelectron component, dotted line—the resonant scattering contribution.

400

130.4 nm

350

Altitude (km)

In order to compare with observations, we consider both sources (resonance scattering of solar radiation and photoelectron impact). To calculate the resonance scattering contribution, we use the solar flux determined from SOLSTICE observations (Rottman, 2000) from the proxy of Woods and Rottman (2002) based on the daily solar flux at 10.7 cm (F10.7) and its 81-day running mean value (F 10:7A ). The distribution of the solar flux between the three lines of the 130.4-nm triplet follows the recommendations of Gladstone (1992). The solar flux at 130.4 nm at the orbital distance of Venus varies from 2.0 to 2.5  1010 photons cm2 s1 for the cases illustrated below. The excitation efficiency (or g-factor) of the triplet by resonance scattering is 2.5  105 s1 for the solar activity conditions illustrated here. Fig. 6 shows the volume excitation rate of O (3S1) atoms for solar maximum conditions (F10.7 ¼ 240, F 10:7A ¼ 177) calculated for equatorial latitudes and noon local time (SZA ¼ 01). The plot indicates that the two processes have a different vertical distribution with the photoelectron impact component showing a peak at 141 km, while the contribution from resonance scattering maximises about 12 km lower. The peak values of the two contributions are close, but the vertically integrated emission rate arising from resonant scattering is about two times as large as the photoelectron component. The vertical distribution of the resonant scattering contribution stems from the complex interplay between the increase of the O density with decreasing altitude (causing increasing attenuation of the UV solar radiation) and the growing importance of photon scattering further away from the line centre as the central wavelengths progressively saturate. Therefore, the vertical distribution of absorption of photons scattering in the 103.4-nm triplet strongly depends on wavelength. Consequently, when combined with the 130.4-nm solar emission line profile, it results in a vertical distribution of resonance scattering with the peculiar shape shown in Fig. 6. Fig. 7

300 RT Source Function 250

200

Primary Source

150

100

100

101

102 103 104 -3 s-1) Volume emission rate (cm

105

106

Fig. 7. Calculated emission rate of the OI 130.4-nm line for the same solar activity conditions as in Fig. 6. Dotted line—total primary excitation rate, solid line—final source function including multiple scattering. The large amplification of the primary source function is a consequence of large optical thickness of the OI 3P–3S triplet.

illustrates the effect of multiple scattering of 130.4-nm photons. The source function, including multiple scattering, is enhanced nearly three orders of magnitude over the initial total (photoelectron+resonance scattering) source.

ARTICLE IN PRESS J.-C. Ge´rard et al. / Planetary and Space Science 56 (2008) 542–552

0

0

50

50

50

100

θ (DEG)

0

θ (DEG)

θ (DEG)

550

100

150

150 0

5

10

Brightness (kR)

15

100

150 0

5

10

15

Brightness (kR)

0 2 4 6 8 10 12 14 Brightness (kR)

Fig. 8. Angular distribution of the calculated brightness of OI 130.4-nm dayglow emission simulated for Pioneer Venus orbits 164 (a), 188 (b) and 206 (c). The observation angle is measured from zenith. The solar conditions were F10.7 ¼ 188 (a), 240 (b) and 159 (c) and the F 10:7A index were 176, 177 and 172, respectively. Solid curve—distribution for the standard O density profile, dotted line—sensitivity when the O density is reduced to 50% of the standard value. Diamonds—observations from the PV-OUVS instrument (Paxton and Meier, 1986). The model profiles account for the 2.51 field of view of the instrument.

The calculated total intensity for nadir viewing from outside the Venus atmosphere is 8.7 kR. This value compares favourably with observations by the PV-OUVS instrument as will be discussed in Section 5. A comparison with a limb scan from PV-OUVS (Paxton and Meier, 1986) is shown in Fig. 8. The scans were obtained during the PV orbits 164, 188 and 206, when the spacecraft was at altitudes of 161.6, 169.3 and 161.8 km, respectively, for SZAs of 411, 16.31 and 341. The F10.7 index was 188, 240 and 159 and the 3-month average F 10:7A indices were 176, 177 and 171, respectively. The calculated angular scan, from the zenith (y ¼ 01) to the nadir (y ¼ 1801) is represented by the solid line. Both the angle of the maximum and the brightness are very well reproduced by the model. Other angular scans have been modelled showing a similarly good agreement with the observations. The calculated 130.4-nm intensities for nadir observations are 6.8, 8.7 and 6.9 kR for the conditions of orbits 164, 188 and 206, respectively. Resonance scattering of the solar 130.4-nm triplet contributes 70% to the total primary source of excitation, the remaining 30% are due to photoelectron impact. This is at variance with previous theoretical studies (Fox and Dalgarno, 1981; Paxton and Meier, 1986), which found that the solar contribution is small in comparison with the photoelectron source. We argue that, in the Earth’s thermosphere, resonance scattering of the solar emission line is the dominant source in the 130.4-nm dayglow, and that the difference of the ionisation cross-sections of the N2 and O2 molecules (dominant species of the Earth thermosphere) and of CO2 (the dominant constituent of the Venus lower thermosphere) can probably not modify the situation in such a manner that the solar contribution would become negligible in the atmosphere of Venus. To examine the sensitivity of the OI 130.4-nm intensity to the atomic oxygen density, Fig. 8 also shows the angular distribution calculated for an O density profile divided by a factor of 2. The peak intensities are decreased by approximately 30%. The changes in the

contribution of the two sources are, as expected, somewhat different. The resonance scattering source drops by 27%, while the photoelectron impact is reduced by 40%. The nadir contributions decrease by 18% and 36%, respectively, resulting in a 22% reduction of the total nadir intensity. The reduction of both contributions is thus nonlinear. The non-linearity of the solar contribution is clearly a consequence of the optical thickness of the OI 3 P–3S1 transition and the role of multiple scattering. The dependence of the photoelectron contribution on the O density indirectly results from the change of density of the target O atoms which redistributes the altitude of the absorption of the EUV solar flux and the energy loss of the ejected photoelectrons. Comparing the observed intensity with the two model curves, one concludes that in all three cases, an O density profile reduced by about 25–30% with respect to the Hedin/VIRA model provides the best fit to the OUVS limb scans. The nadir intensity observed with PV-OUVS near the sub-solar point was typically 5–6 kR (Alexander et al., 1993), while the disc-averaged brightness deduced from the HUT spectrum (Feldman et al., 2000) was close to 2.8 kR. The Venera 11/12 spectrometers (Bertaux et al., 1981) measured 6.2 kR in December 1978, at a time approaching solar maximum. Both the PV-OUVS and the Venera observations were made near solar maximum, while the HUT observations were made for quiet conditions with F10.7 ¼ 76 and F 10:7A ¼ 82 (13 March 1995). It should be remembered that the comparison between observations and model outputs is complicated in the present case by the uncertainty affecting the cross-section of the photoelectron impact source of excited oxygen. Only effective emission cross-sections are available attempting to account for cascades from upper lying states of the oxygen atom to the 3S1 state. This effective cross-section is atmosphere-dependant owing to the complex coupling between radiative transfer, absorption and branching of the various emissions involved in this problem. In some

ARTICLE IN PRESS J.-C. Ge´rard et al. / Planetary and Space Science 56 (2008) 542–552

earlier studies, this cross-section was left as a free parameter to be adjusted on the data (Meier et al., 1983). The input solar flux is somewhat uncertain as well. The data used to set up the proxy model have an uncertainty ranging from 10% to 20% depending on wavelength, due to instrument calibration issues. Another 10% uncertainty arises from the use of the F10.7 and F 10:7A indices that alone cannot account for the whole complexity of the physics governing the sun. Indeed, a 20–30% lower solar contribution would nicely close up the observations and our theoretical results presented in Fig. 8. As discussed previously, since the 3P–5S1 transition is forbidden, no resonance scattering process has been considered for this transition. Photoelectron impact on CO2, CO and O is the only important source of OI 135.6nm emission. In fact, electron impact on O largely dominates the production of 5S1 atoms in the Venus thermosphere. The calculated column excitation rate for PV orbit 188 is 7.4  108 photons cm2 s1, corresponding to an emission rate of 0.74 kR for a nadir viewing direction, and accounting for the weak absorption by CO2. At the low resolution of the PV-OUVS instrument, the 135.6-nm line is blended with the (14.5) band of the CO 4P system (Durrance et al., 1980). Following subtraction of the CO contribution, the observed 130.4/135.6 ratio is 8 for a sum of spectra in which the illumination, emission and phase angles were less than 601. The limb intensity of the Cameron bands was estimated 9.5 kR from PV orbit 185 (Meier et al., 1983), a value very close to our calculated limb emission rate of 9.4 kR. The 135.6-nm brightness observed in the HUT spectrum is about 0.6 kR, leading to a disc-averaged 130.4/135.6 ratio of 4.670.4. Comparing this observed 130.4/135.6 intensity ratio with our calculated value of 5.8 for a view angle of 301 and 4.5 for discaveraged brightness, we find that our model reproduces well these observations. Note that our disc-average value is a first order approximation that does not account for the variation of the SZA across the disc, but assumes a constant SZA of 301.

5. Conclusions A new numerical Monte Carlo algorithm to simulate energetic electron transport in the Venus thermosphere has been applied to the simulation of auroral electron interaction with the Venus night-side thermosphere. The brightness of the OI 130.4, OI 135.6 and CO Cameron bands have been calculated using an energy spectrum based on RPA measurements of suprathermal electrons on board the PV. Multiple scattering of the 130.4-nm triplet is considered and we find that the emergent intensity is increased by a factor of about 3 in comparison with the optically thin case. When this effect is taken into account, the predicted 130.4/135.6 ratio for precipitation of soft electron of 14 eV is about 6. This value is in good agreement with the measurements made with the PV-

551

OUVS instrument which detected a weak auroral emission with a ratio of 5.572.7. The average precipitated flux deduced from the typical brightness of the auroral emission is about 2  103 mW m2 at solar maximum, corresponding to 30% of the suprathermal electron spectrum reported by Knudsen and Miller (1985). It is also in satisfactory agreement with earlier estimates of the column production rates of the ultraviolet emission by Fox and Stewart (1991) which did not include a treatment of the O (3S1) multiple scattering. The calculated Cameron band intensity is about 6 R, a value lower than the 12 R predicted by the model for 2  103 mW m2 of precipitation. One must, however, consider the uncertainties on the excitation cross-section of the CO a3P state and on the observed brightness. For example, using the Conway (1981) crosssection would lead to an underestimate of the observed mean brightness. The comparison of calculated OI 130.4- and 135.6-nm dayglow intensities with PV-OUVS and HUT observations indicates that the model satisfactorily reproduces the global intensity of both emissions observed with HUT, as well as the PV-OUVS 130.4-nm limb brightness distribution and the 130.4/135.6 nm intensity ratio. The primary production of 130.4-nm photons is dominated by resonant scattering of solar photons at all altitudes, except near the emission peak, where photoelectron impact plays an important role. When integrated over the full column, the resonance scattering source is larger than the photoelectron source. We also note that the calculations were made with the VIRA standard atmosphere. This model is statistical and is not expected to accurately reproduce specific conditions. Nevertheless, the brightness of the calculated OI emission is in general good agreement with the observations. We conclude that the current knowledge of the oxygen aurora/dayglow emissions make it possible to use future FUV observations to survey the global distribution of atomic oxygen in the Venus thermosphere and to remotely quantify the characteristics of night-side auroral precipitation. We showed that the intensity of the OI 130.4-nm emission does not vary proportionally to the O density in the lower thermosphere, but provides a tool to probe the abundance of atomic oxygen in the dayside lower thermosphere and its variations with latitude, local time and solar activity conditions.

Acknowledgements J.C. Ge´rard and B. Hubert are supported by the Belgian Fund for Scientific Research (FNRS). This work was also funded by grant RFBR 05-02-17165, the Belgian National Fund for Collective Fundamental Research (FRFC grant 2.4517.02) and by the PRODEX Programme managed by the European Space Agency in collaboration with the Belgian Federal Science Policy Office.

ARTICLE IN PRESS 552

J.-C. Ge´rard et al. / Planetary and Space Science 56 (2008) 542–552

References Ajello, J.M., 1971. Emission cross sections of CO2 by electron impact in the interval 1260–4500 A˚. II. J. Chem. Phys. 55, 3169–3177. Alexander, M.J., Stewart, A.I.F., Solomon, S.C., Bougher, S.W., 1993. Local time asymmetries in the Venus thermosphere. J. Geophys. Res. 98, 10849–10871. Bertaux, J.L., Blamont, J.E., Lepine, V.M., Kurt, V.G., Romanova, N.N., Smirnov, A.S., 1981. Venera 11 and Venera 12 observations of E.U.V. emissions from the upper atmosphere of Venus. Planet. Space Sci. 29, 149–166. Bisikalo, D.V., Shematovich, V.I., Ge´rard, J.-C., 1995. Kinetic model of the formation of the hot oxygen geocorona. II. Influence of O+ ion precipitation. J. Geophys. Res. 100, 3715–3720. Broadfoot, A.L., Kumar, S., Belton, M.J.S., McElroy, M.B., 1974. Ultraviolet observations of Venus from Mariner 10: preliminary result. Science 183, 1315–1318. Conway, R.R., 1981. Spectroscopy of the Cameron bands in the Mars airglow. J. Geophys. Res. 87, 4767–4775. Durrance, S.T., Barth, C.A., Stewart, A.I.F., 1980. Pioneer Venus observations of the Venus dayglow spectrum 1250–1430 A. Geophys. Res. Lett. 7, 222–224. Erdman, P.W., Zipf, E.C., 1983. Electron-impact excitation of the Cameron System (a 3P-X1S) of CO. Planet. Space Sci. 31, 317–321. Feldman, P.D., Burgh, E.B., Durrance, S.T., Davidsen, A.F., 2000. Farultraviolet spectroscopy of Venus and Mars at 4 Angstrom resolution with the Hopkins Ultraviolet Telescope on ASTRO-2. Ap. J. 538, 395–400. Fox, J.L., Bougher, S.W., 1991. Structure, luminosity and dynamics of the Venus thermosphere. Space Sci. Rev. 55, 357–489. Fox, J.L., Dalgarno, A., 1981. Ionization, luminosity and heating of the upper atmosphere of Venus. J. Geophys. Res. 86, 629–639. Fox, J.L., Stewart, A.I.F., 1991. The Venus ultraviolet aurora: a soft electron source. J. Geophys. Res. 96, 9821–9828. Furlong, J.M., Newell, W.R., 1996. Total cross section measurement for the metastable a 3P state of CO. J. Phys. B 29, 331–338. Ge´rard, J.-C., Hubert, B., Bisikalo, D.V., Shematovich, V.I., 2000. A model of the Lyman-a line profile in the proton aurora. J. Geophys. Res. 105, 15795–15806. Gladstone, G.R., 1985. Radiative transfer of resonance lines with internal sources. J. Quant. Spectr. Radiat. Transfer 33, 453–458. Gladstone, G.R., 1992. Solar OI 1304 A triplet line profiles. J. Geophys. Res. 97, 19125–19519. Green, A.E.S., Sawada, T., 1972. Ionization cross sections and secondary electron distributions. J. Atmos. Terr. Phys. 34, 1719–1728. Gringauz, K.I., Verigin, M.I., Breus, T.K., Gombosi, T., 1979. Interaction of electrons in the optical umbra of Venus with the planetary atmosphere: the origin of the nighttime ionosphere. J. Geophys. Res. 84, 2123–2128. Hedin, A.E., Niemann, H.B., Kasprzak, W.T., Seiff, A., 1983. Global empirical model of the Venus thermosphere. J. Geophys. Res. 88, 73–83. Hord, C.W., Barth, C.A., Esposito, L.W., Mcclintock, W.E., Pryor, W.R., Simmons, K.E., Stewart, A.I.F., Thomas, G.E., Ajello, J.M., Lane, A.L., West, R.E., Sandel, B.R., Broadfoot, A.L., Hunten, D.M., Shemansky, D.E., 1991. Galileo ultraviolet spectrometer experiment: initial Venus and interplanetary cruise results. Science 253, 1548–1550. Huebner, W.F., Keady, J.J., Lyon, S.P., 1992. Solar photo rates for planetary atmospheres and atmospheric pollutants. Astrophys. Space Sci. 195, 1–289.

Itikawa, Y., 2002. Cross sections for electron collisions with carbon dioxide. J. Phys. Chem. Ref. Data 31, 749–767. Jackman, C.H., Garvey, R.H., Green, A.E.S., 1977. Electron impact on atmospheric gases. I-Updated cross sections. J. Geophys. Res. 82, 5081–5090. Johnson, P.V., McConkey, J.W., Tayal, S.S., Kanik, I., 2005. Collisions of electrons with atomic oxygen: current status. Can. J. Phys. 83, 589–616. Knudsen, W.C., Miller, K.L., 1985. Pioneer Venus superthermal electron flux measurements in the Venus umbra. J. Geophys. Res. 90, 2695–2702. Meier, R.R., 1991. Ultraviolet spectroscopy and remote sensing of the upper atmosphere. Space Sci. Rev. 58, 1–185. Meier, R.R., Anderson Jr., D.E., Stewart, A.I.F., 1983. Atomic oxygen emissions observed from Pioneer Venus. Geophys. Res. Lett. 10, 214–217. Paxton, L., Anderson, D.E.N., 1992. Venus and Mars: atmospheres, ionospheres, and solar wind interactions. In: Proceedings of the Chapman Conference, Balatonfured, Hungary, Washington, DC, American Geophysical Union, pp. 113–189. Paxton, L.J., Meier, R.R., 1986. Reanalysis of Pioneer orbiter ultraviolet spectrometer data: OI 1304 intensities and atomic oxygen densities. Geophys. Res. Lett. 13, 229–232. Philipps, J.L., Stewart, A.I.F., Luhmann, J.G., 1986. The Venus ultraviolet aurora: observations at 130.4 nm. Geophys. Res. Lett. 13, 1047–1050. Rottman, G., 2000. Variations of solar ultraviolet irradiance observed by the UARS SOLSTICE—1991 to 1999. Space Sci. Rev. 94, 83–91. Sawada, T., Strickland, D.J., Green, A.E.S., 1972. Electron energy deposition in CO2. J. Geophys. Res. 77, 4812–4818. Shematovich, V.I., Bisikalo, D.V., Ge´rard, J.-C., 1994. A kinetic model of the formation of the hot oxygen geocorona. I. Quiet geomagnetic conditions. J. Geophys. Res. 99, 23,217–23,228. Shematovich, V.I., Bisikalo, D.V., Ge´rard, J.-C., Cox, C., Bougher, S.W., Leblanc, F., 2007. Monte Carlo model of the electron transport for the calculations of Mars dayglow emissions. J. Geophys. Res., in press. Spenner, K., Knudsen, W.C., Lotze, W., 1996. Suprathermal electron fluxes in the Venus nightside ionosphere at moderate and high solar activity. J. Geophys. Res. 101, 4557–4564. Spenner, K., Dobe, Z., Nagy, A.F., Knudsen, W.C., Lotze, W., 1997. Photoelectron fluxes in the Venus dayside ionosphere. J. Geophys. Res. 102, 2577–2583. Stewart, A.I., 1980. Design and operation of the Pioneer Venus orbiter ultraviolet spectrometer. IEEE Trans. Geosci. Remote Sensing GE18, 65–70. Stone, J., Zipf, E.C., 1974. Electron impact excitation of the 3S1 and 5S1 states of atomic oxygen. J. Chem. Phys. 60, 4237–4243. Tobiska, W.K., 2004. SOLAR2000 irradiances for climate change, aeronomy and space system engineering. Adv. Space. Res. 34, 1736–1746. Venus International Reference Atmosphere (VIRA), 1985. Kliore, A.J., Moroz, V.I., Keating, J.M. (Eds.), Adv. Space Res. 5(11), Pergamon Press. Woods, T., Rottman, G., 2002. Ultraviolet variability over time periods of aeronomic interest, atmospheres in the solar system: comparative aeronomy. In: Mendillo, M., Nagy, A., Waite, J.H. (Eds.), Geophysical Monograph 130. American Geophysical Union, Washington, DC, p. 221. Zipf, E.C., Erdman, P.W., 1985. Electron impact excitation of atomic oxygen: revised cross sections. J. Geophys. Res. 90, 11087–11090.