UVIS observations of the FUV OI and CO 4P Venus... B. Hubert , J.C. Gérard , J. Gustin

Icarus 207 (2010) 549–557
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UVIS observations of the FUV OI and CO 4P Venus dayglow during the Cassini flyby
B. Hubert a,*, J.C. Gérard a, J. Gustin a, V.I. Shematovich b, D.V. Bisikalo b, A.I. Stewart c, G.R. Gladstone d
Laboratoire de Physique Atmosphérique et Planétaire, Université de Liège, 17, Allée du 6 Août, Bât. B5c, B-4000 Liège, Belgium
Institute of Astronomy of the Russian Academy of Sciences, 48, Pyatnitskaya Street, 119017 Moscow, Russia
Laboratory for Atmospheric and Space Physics, University of Colorado, 1234, Innovation Drive, Boulder, CO 80303, USA
Southwest Research Institute, 6220, Culebra Road, San Antonio, TX 78228-0510, USA
a r t i c l e
i n f o
Article history:
Received 1 April 2009
Revised 18 December 2009
Accepted 19 December 2009
Available online 4 January 2010
Venus, Atmosphere
Ultraviolet observations
Radiative transfer
a b s t r a c t
We analyze FUV spatially-resolved dayglow spectra obtained at 0.37 nm resolution by the UVIS instrument during the Cassini flyby of Venus. We use a least-squares fit method to determine the brightness
of the OI emissions at 130.4 and OI 135.6 nm, and of the bands of the CO fourth positive system which
are dominated by fluorescence scattering. We compare the brightness observed along the UVIS foot track
of the two OI multiplets with that deduced from a model of the excitation of these emissions by photoelectron impact on O atoms and resonance scattering of the solar 130.4 nm emission. The large optical
thickness 130.4 nm emission is accounted for using a radiative transfer model. The airglow intensities
are calculated along the foot track and found to agree with the observed 130.4 nm brightness within
10%. The modeled OI 135.6 nm brightness is also well reproduced by the model. The oxygen density
profile of the VTS3 model is found to be consistent with the observations. We find that self-absorption
of the (0, v00 ) bands of the fourth positive emission of CO is important and we derive a CO vertical column
of about 6.4 1015 cm2 in close agreement with the value provided by the VTS3 empirical atmospheric
Ó 2009 Elsevier Inc. All rights reserved.
1. Introduction
The atomic oxygen triplet at 130.2, 130.4 and 130.6 nm arises
from the dipole-allowed O(3S0) ? O(3P) transition. The OI
130.4 nm dayglow triplet is the result of resonance scattering of
photons emitted in the broad solar emission line, photoelectron impact on ground state 3P atomic oxygen and, to a minor extent, electron impact and photodissociation of CO and CO2. Carbon dioxide
absorbs the OI 130.4 nm emission below 125–130 km. The emission is optically thick so that the airglow brightness is not proportional to the oxygen content above the altitude of CO2 absorption,
but multiple scattering occurs between the region of initial photon
production and the outside observer. Atomic oxygen becomes the
main constituent 15–20 km above the homopause located near
135 km (Von Zahn et al., 1983), where molecular diffusive processes control vertical transport. Sources for O atoms in the Venus
thermosphere are photodissociation of CO2 and CO and, less importantly, dissociative recombination in the ionosphere.
The first observations of the Venus oxygen dayglow were made
using a rocket-borne spectrometer by Moos et al. (1969) and Moos
* Corresponding author. Address: University of Liège, Institut d’Astrophysique et
de Géophysique, 17, Allée du 6 Août, Bât. B5c, B-4000 Liège, Belgium. Fax: +32 (0)4
366 97 11.
E-mail address: B.Hubert@ulg.ac.be (B. Hubert).
0019-1035/$ - see front matter Ó 2009 Elsevier Inc. All rights reserved.
and Rottman (1971) from rocket launches on 5 December 1967
(near solar maximum conditions) and 25 January 1971 (solar maximum conditions), respectively. The reported disk intensity was
5.5 ± 0.5 kR on 25 January 1971 (Rottman and Moos, 1973). The
OI triplet was also detected during the Mariner 10 flyby on 5 February 1974 (near solar minimum) with a brightness of 17 kR
(Broadfoot et al., 1974). The spectrometers on board Venera 11
and 12 (which entered the Venus atmosphere on 25 and 21
December 1978, respectively, as the Sun activity was rising towards solar maximum conditions) measured a disk intensity of
6.4 kR (Bertaux et al., 1981). Stewart and Barth (1979) obtained
mid-resolution (1.3 nm) dayglow spectra with the Orbiting
UltraViolet Spectrometer (Stewart, 1980) on board Pioneer Venus
(PV–OUVS), that reached Venus in December 1978. Paxton and
Meier (1986) analyzed the OI 130.4 nm emission observed with
PV–OUVS and compared the observed limb scans with the results
of a Monte Carlo radiative transfer model. They found atomic
oxygen densities consistent with in situ measurements. A large
set of 130.4 nm dayglow images obtained between 1980 and
1990, both near periapsis and during the inbound orbital segment,
was analyzed by Alexander et al. (1993). The emission brightness
widely varied with the geometry of the observations with typical
local values between 2 and 8 kR. The observations were compared
with model calculations including resonance scattering of the solar
line and photoelectron impact as sources of O(3S0) atoms. Radiative
B. Hubert et al. / Icarus 207 (2010) 549–557
transfer was solved using the same Monte Carlo method as Paxton
and Meier (1986). Spherical geometry effects were included using
an approximation to the Chapman function. They found that the
intensity of the 130.4 nm triplet varies linearly with the F10.7 solar
activity proxy. A high-contrast asymmetry in local time was observed at latitudes poleward of 30°, consisting of a generally larger
oxygen dayglow brightness in the evening sector than in the morning sector. It was interpreted as a factor of 2 increase in O density
at the nighttime terminator over corresponding day times. Enhanced eddy mixing in the morning hours was proposed as a possible cause of this density asymmetry. Further dayglow
observations were made at 0.4 nm resolution by Feldman et al.
(2000) who used the Hopkins Ultraviolet Telescope (HUT) instrument on board the Space Shuttle to observe the Venus disk in
the wavelength range from 82 to 184 nm on 13 March 1995 (near
solar minimum). They identified and provided disk brightness of
spectral signatures from CI, CII, OI, Ar, NI, CO and CO2 and measured a 130.4 nm disk brightness of 2.8 ± 0.28 kR.
The O(5S0) ? O(3P) transition is a spin-forbidden doublet at
135.6–135.8 nm and is therefore solely excited by photoelectron
impact. Its brightness was estimated 2.7 ± 0.5 kR by Rottman and
Moos (1973). The 135.6 nm emission however falls within the
spectral interval of the bright fourth positive system of molecular
CO (CO 4P), as well as the OI 130.4 nm emission, and the detailed
analysis of these emissions was investigated in several studies.
First, Stewart and Barth (1979) showed the presence of limb
brightening of the 135.6 nm doublet in the PV–OUVS dayglow
spectra. Second, Durrance et al. (1980) measured the FUV dayglow
spectrum between 125 and 143 nm with PV–OUVS and suggested
that the (14, 3) and (14, 4) CO 4P bands are blended respectively
with the OI 130.4 and 135.6 nm emissions at the OUVS spectral
resolution. Third, these bands were resolved in high-resolution
(0.04 nm) observations acquired with the International Ultraviolet Explorer (IUE) satellite (Durrance, 1981). The bright OI
130.4 nm was also very conspicuous in the IUE spectrum. Finally,
the 135.6 nm limb brightening was interpreted by Durrance et al.
(1981) as a consequence of the contamination by the (14, 4) CO
4P band. Recently, a 135.6 nm disk brightness of 605 ± 50 R was
observed in the HUT spectrum by Feldman et al. (2000).
Rottman and Moos (1973) found that photoelectron impact on
O atoms was the major source of OI emissions at 130.4 and
135.6 nm. Fox and Dalgarno (1981) calculated the O(3S0) and
O(5S0) excitation rates by electron impact and showed that impact on O atoms is about an order of magnitude larger than on
CO2. Meier et al. (1983) included both fluorescence and electron
impact sources and used a Monte Carlo model for radiative transfer of the 130.4 nm emission with a partial frequency redistribution. They were able to reproduce PV–OUVS limb scans, but they
needed to decrease the O densities by a factor of 2. This discrepancy was clarified by Paxton and Meier (1986) who obtained
good agreement without any O density adjustment following
revision of the electron impact excitation cross section. Alexander
et al. (1993) also considered electron impact on O and solar resonance scattering as the important sources of O(3S0) atoms. In the
example they presented, the photoelectron source dominates below 180 km while resonance scattering is the main source at
higher altitude. In a recent study, Gérard et al. (2008) revisited
the PV–OUVS observations of the OI 130.4 nm emission. They
compared OUVS spin scans with outputs from a model including
a Monte Carlo computation of the photoelectron population, excitation rate of the OI 3S0 upper state of the 130.4 nm multiplet by
the solar radiation and electron impact, and radiative transfer of
these emissions through the Venus upper atmosphere. They
found that the solar resonance source is the globally dominant
contribution to the OI 130.4 nm emission from the atmosphere
of Venus. The calculated I(130.4)/I(135.6) 8 intensity ratio was
also in good agreement with the value observed in the HUT
Finally, 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–X1R transition. Potential sources include photodissociative excitation and electron
impact dissociation of CO2, electron impact excitation of CO, dissociative recombination of COþ
2 and fluorescent scattering by ground
state CO.
In the present study, we present the spatially-resolved spectra
obtained with the Ultraviolet Imaging Spectrograph (UVIS) during
the Cassini flyby of Venus. We compare the observed and modeled
brightness of the 130.4 and 135.6 nm OI emissions. The absolute
brightness and the ratio of these two emissions, independent of
the absolute instrumental calibration, is also examined. The brightness of the CO 4P band system is determined along the track of the
Cassini spacecraft and the contribution of the (14, 4) CO 4P band to
the measurement of the 135.6 nm emission is also discussed.
2. Observations
The Cassini spacecraft was launched on 15 October 1997. On its
long journey to Saturn, the spacecraft took a gravitational assist to
gain energy from Venus on 24 June 1999. The UVIS instrument on
board Cassini (Esposito et al., 1998) obtained a series of FUV spectra during this flyby, at a time period of rising solar activity, when
the F10.7 solar index was 214 at Earth distance. During the flyby,
Cassini reached an altitude of closest approach of 602 km. The
spacecraft had to be oriented so that its 4-m antenna shielded
the payload from the Sun. This required UVIS to look in a direction
nearly perpendicular to the Sun-spacecraft line, so that the phase
angle remained close to 90°. A total of 55 records of 32 s has been
obtained along the track. Twenty-two of them showed dayglow
emissions, as the UVIS field of view intersected the illuminated
disk of Venus. The latitude of tangent point of the line of sight varied from 24°N to 15°S. The solar zenith angle decreased along
the track from 90° at the morning terminator to 0° when Cassini
left Venus and UVIS observed the planetary limb. Fig. 1 shows
the track geometry and illustrates the variation of the solar zenith
angle and emission angle, i.e. the angle between the line of sight
and local zenith at the altitude of UV emission, during the Cassini
flyby. Periapsis occurred at 20:30:07 UT. The spectral range of UVIS
extends from 111.5 to 191.3 nm in the FUV range, thus including
the very bright Lyman-a emission of atomic hydrogen as well as
the CO A1P ? X1R+ 4P band system and the OI emissions at
130.4 and 135.6 nm. The high-resolution slit of UVIS was used, producing spectra with an observational line spread function of
0.37 nm FWHM, that is about 3.5 times less than PV–OUVS, and
similar to the HUT instrument. The UVIS FUV slit is composed of
1024 pixels in the dispersion direction and 64 pixels in the spatial
direction. The full spectral resolution has been used during the Venus observations, while the spatial direction has been rebinned by
16 pixels, leaving a resolution of four pixels along the spatial direction. Each record presented here is the sum of the four spatial
These conditions offered the possibility to collect measurements of the absolute brightness of various emission features with
a high signal to noise ratio. The UVIS field of view along the slit is
64 mrad, corresponding to 450 km projected on the planet surface from an altitude of 7000 km. The slit was oriented nearly
perpendicular to the ecliptic plane. The spacecraft moved
500 km during the 32 s integration period of each record. Spatial
pixels have been summed within each record in order to improve
the signal/noise ratio. Consequently, the UVIS spectra represent
local measurements of the FUV emission, in contrast with a full
B. Hubert et al. / Icarus 207 (2010) 549–557
Fig. 1. (a) Trace of UVIS’s field of view across Venus during Cassini’s swingby on June 24, 1999. The center of the disk is at 0° latitude, hour angle 120°. The solid curves on the
disk are the traces of the middle and the two ends of UVIS’s slit, which subtended 64 1 mrad. The line of sight for the middle of the slit is shown every 60 s from closest
approach (C/A) – 600 s to C/A + 60 s, and also at the times of first and last contact with the disk. The lengths of the line of sight at first contact, C/A, and last contact were 7700,
1200, and 3000 km. During each 32-s UVIS integration time the spacecraft travelled more than 400 km along its trajectory. (b) Time variation of the solar zenith angle and
emission angle of the foot track of the UVIS field of view during the sequence of dayglow observations from Cassini. Each diamond symbol corresponds to a spectral record; a
crossed diamond is marked every five record number.
disk measurement from Earth orbit, such as the HUT spectrum.
Therefore, UVIS offered the possibility to obtain a spectral resolution comparable to that of HUT, while collecting spatially resolved
observations of the absolute brightness of the Venus FUV
3. Analysis
3.1. Method
The UVIS instrument and its operation were described by
Esposito et al. (2004). The in-flight calibration was used to convert
the count into physical units. In addition, an empirically derived
background noise level of 4.5 104 count per pixel due to the
Cassini radioisotope thermoelectric generators has been removed,
and a flat-field correction derived from observations of star Spica
(Steffl et al., 2004) has been applied. Another source of background
signal that affects the recorded spectra is scattering of HI Lyman-a
photons from the grating that reach the detector. This wavelengthdependent background, proportional to the Lyman-a intensity, has
been modeled, scaled to the observed H Lyman-a and subtracted
from the data.
We first describe the spectral identification and the derivation
of the brightness of the emission features. Fig. 2 presents the average brightness of the 22 calibrated spectra in the wavelength range
125–180 nm, obtained as the UVIS FUV slit intersected the illuminated disk. The statistical photon noise is very small and associated
1 r error bars are comparable or less than the line thickness of
the plot, except close to 180 nm where it becomes more significant.
B. Hubert et al. / Icarus 207 (2010) 549–557
Fig. 2. Averaged FUV dayglow spectrum obtained by the UVIS instrument during the Cassini flyby of Venus. Background counts and contribution from the bright Lyman-a
line have been removed. Carbon and oxygen lines have been identified. The brightest bands of the CO 4P system are labeled as (v0 , v00 ).
Several bands of the CO 4P system have been identified. The solar
Lyman-a line plays a key role in the excitation of the CO 4P band
system because an incidental resonance with the (14, 0) band produces optical pumping of the CO 4P system (Durrance, 1981 and
references therein). The spectrum also shows lines from CI and
CII transitions. At the UVIS resolution, the CO 4P (14, 4) band and
the OI 3P–3S0 transition at 135.6 nm appear blended. The OI
130.4 nm optically thick line is the second brightest spectral feature after Lyman-a (which is not shown).
This spectrum is the complex result of several physical processes. First, the solar Lyman-a line produces optical pumping of
the CO 4P (14, v00 ) band system through resonance with the
(14, 0) band (Durrance, 1981 and references therein). Cascading
and radiative transfer of the CO 4P bands is important, especially
for those transitions connecting to v00 = 0, since transitions to the
ground vibrational level can be optically thick. In addition, the
CO(A1P) state is also excited by photoelectron impact on CO, photodissociation of CO2, photoelectron impact on CO2 and dissociative recombination of COþ
CO þ hm ! COðA1 PÞ k < 150 nm
CO þ e ! COðA1 PÞ þ e Ee > 8:02 eV
CO2 þ hm ! COðA1 PÞ þ O k < 92:3 nm
CO2 þ e ! COðA1 PÞ þ O þ e Ee > 13:5 eV
COþ2 þ e ! COðA1 PÞ þ O
The dominant source of CO(A1P) airglow is fluorescent scattering. It contributes about 90% of the total intensity and it is the
source of the v0 = 14 progression, among others (Fox and Dalgarno,
Second, as mentioned before, the Venus thermosphere is optically thick to the OI 130.4 nm emission. The primary excitation
of its upper state results from photoelectron impact on oxygen
atoms and from resonance scattering of the solar 130.4 nm emission line:
Oð3 PÞ þ hm130:4 ! Oð3s 3 S0 Þ
Oð3 PÞ þ e ! Oð3s 3 S0 Þ Ee > 9:49 eV
Third, the OI 135.6 nm emission is excited by photoelectron impact on oxygen atoms:
Oð3 PÞ þ e ! Oð3s 5 S0 Þ þ e Ee > 9:14 eV
Radiative transfer of this emission through the thermosphere of
Venus may be neglected because this transition is spin forbidden,
so that this emission may be considered as optically thin.
3.2. Least-squares fitting procedure
The wavelength range of the CO 4P emission also contains several atomic lines which are partly blended. To separate the individual contributions, the spectrum of each record is analyzed using a
least-squares fit method where the wavelengths of the CO 4P
bandheads and the most important atomic emissions present in
the spectra are preset. Multiplets from neutral atomic carbon at
126.1, 127.7, 132.9, 143.2, 146.3, 156.1 and 165.7 nm are accounted for, as well as the CII multiplet at 133.5 nm, the OI multiplet at 130.217, 130.486 and 130.603 nm and the OI doublet at
135.560 and 135.851 nm. A single wavelength is assumed for the
CI and CII emissions (which are not resolved by the instrument),
but the individual wavelength of each line of the oxygen multiplets
are accounted for since we want to carry a detailed analysis of
these emissions. The wavelength of the A–X (v0 , v00 ) bands are taken
from Kurucz (1976). The intensity of the bands of the (v0 , v00 ) progressions are fitted assuming that the relative intensity of the
members of each v0 progression is proportional to the AAv0 Xv00 Einstein coefficients derived from Kurucz (1976). However, the intensities of the (v0 , 0) bands, which can be self-absorbed, are
determined separately. The synthetic spectra which are leastsquares fitted take into accounts the broadening of the emissions
by the instrumental line spread function. The synthetic spectra
are degraded to the UVIS spectral resolution by combining the
modeled lines and bands with the point spread function of the
instrument. Finally, the procedure determines the individual intensity of all CO 4P bands present in the 129–180 nm interval, avoiding the potentially troublesome wings of the Lyman-a line, and
assuming a single wavelength for each band (the bandhead wavelength). To fit the OI 130.4 nm emission, we use the relative intensity of the three lines in the multiplet provided by our radiative
transfer model. Since the 3P–5S0 transition is optically thin, the relative intensity of the components of the OI 135.6 nm multiplet is
simply proportional to their emission probability. The CI and CII
lines are also fitted individually. Note that, in the fitting procedure,
a very small weight is attributed to wavelength bins between 175
and 180 nm, to cope with the contamination by the CO Cameron
bands. Tables 1 and 2 list the observed brightness of the OI
B. Hubert et al. / Icarus 207 (2010) 549–557
Table 1
Variation of the brightness of the observed (column 4) and modeled (column 5)
130.4 nm intensity for UVIS records intersecting the sunlit disk during the Cassini
Venus flyby. The values of the solar zenith angle (SZA) and emission angle (EMA) are
indicated in columns 2 and 3.
UVIS record
OI 130.4 nm dayglow
SZA (°)
EMA (°)
UVIS fit (kR)
Model (kR)
Difference (%)
0.06 ± 0.01
0.07 ± 0.01
0.26 ± 0.01
0.82 ± 0.02
1.19 ± 0.03
1.57 ± 0.03
1.90 ± 0.03
2.23 ± 0.04
2.60 ± 0.04
3.01 ± 0.04
3.53 ± 0.04
4.04 ± 0.05
4.60 ± 0.05
4.94 ± 0.05
4.99 ± 0.05
5.44 ± 0.05
5.86 ± 0.06
6.57 ± 0.06
7.02 ± 0.06
7.32 ± 0.06
7.29 ± 0.06
8.17 ± 0.07
8.93 ± 0.07
8.72 ± 0.07
Table 2
Variation of the brightness of the observed (column 4) and modeled (column 5)
135.6 nm intensity for each UVIS records intersecting the sunlit disk during the
Cassini Venus flyby. The values of the solar zenith and emission angles are indicated
in columns 2 and 3.
Record #
OI 135.6 nm dayglow
SZA (°)
EMA (°)
UVIS fit (kR)
Model (kR)
Difference (%)
0.0 ± 0.004
0.0 ± 0.004
0.02 ± 0.01
0.15 ± 0.01
0.20 ± 0.01
0.24 ± 0.02
0.28 ± 0.02
0.34 ± 0.02
0.37 ± 0.02
0.40 ± 0.02
0.45 ± 0.02
0.55 ± 0.02
0.60 ± 0.02
0.59 ± 0.02
0.68 ± 0.02
0.74 ± 0.02
0.80 ± 0.02
0.97 ± 0.03
1.04 ± 0.03
1.13 ± 0.029
1.20 ± 0.03
1.34 ± 0.03
1.73 ± 0.04
2.16 ± 0.04
130.4 nm and OI 135.6 nm emissions for each record, determined
with the least-squares fitting procedure.
3.3. The CO fourth positive bands
Fig. 3 shows the CO 4P brightness along the Cassini track
plotted versus the record number. These values are obtained by
Fig. 3. Total intensity of the CO 4P band system (solid line) and contribution of the
(0, 1) band multiplied by 10 (dotted line), observed with UVIS instrument as a
function of decreasing solar zenith angle along the UVIS footprint. The dash-dotted
line is the observed contribution of the (v0 , 0) transitions, the dashed line represents
the calculated intensity that the (v0 , 0) transitions would have without selfabsorption.
summing all CO 4P contributions to the observed spectra determined with the least-squares fitting method and excluding other
emission features. The value of the solar zenith angle (SZA) is also
indicated along the upper horizontal axis. Note that not only the
solar zenith angle varied along the UVIS footprint, but also the
viewing geometry as illustrated in Fig. 1. The brightness reaches
80 kR at the limb and decreases with increasing solar zenith angle. A detailed examination shows that the CO 4P (0, 1) band has a
peak brightness of 5.8 kR at the limb. We find that the brightness
of the (0, 1) band along the Cassini track is generally larger than the
0.75 kR reported by Feldman et al. (2000) for the sunlit disk near
solar minimum (F10.7 82), whereas the total CO 4P brightness is
larger than the 25 ± 5 kR obtained by Rottman and Moos (1973),
who did not identify all the CO 4P bands. Table 3 lists the total
brightness of the CO 4P system determined with the least-squares
fitting procedure. A more detailed comparison between these and
the disk brightness observed with HUT will be made in later section. Fig. 3 also shows the brightness of the A–X (v0 , 0) emissions
summed over all v0 values (dash-dotted lines), to be compared to
the intensity obtained if the (v0 , 0) relative intensities were simply
proportional to the AAv0 –Xv00 coefficients within the progression. The
large difference shows that self-absorption plays a very important
role in the radiative transfer of the CO 4P bands in the Venus
The CO column may be estimated from the spectrum using the
remote sensing method suggested by Durrance (1981). The intensity of a CO band excited almost entirely by fluorescence scattering
may be used to convert the observed emission rate into a CO column density once the relevant g-factor is known. For this purpose,
we use the intensity of the A–X (14, 3) band and the corresponding
g-factor obtained by scaling the g -factor used by Feldman et al.
(2000) according to the Lyman-a solar flux. Feldman et al. used recent measurements of the CO 4P oscillator strength by Eidelsberg
et al. (1999), the Lyman-a line profile measured by Lemaire et al.
(1998) and the Lyman-a solar flux measured with UARS-SOLSTICE
(Woods et al., 1996). Unfortunately, direct measurement of the Lyman-a solar flux is not available for the time of the Cassini flyby,
and we must use a simple interpolation between the available
measured fluxes before and after the day of interest. From the
B. Hubert et al. / Icarus 207 (2010) 549–557
Table 3
Variation of the brightness of the observed intensity of the total CO fourth positive
emission for UVIS records intersecting the sunlit disk during the Cassini Venus flyby.
The values of the solar zenith and emission angles are indicated in columns 2 and 3.
Record #
CO fourth positive dayglow
SZA (°)
EMA (°)
UVIS fit (kR)
0.41 ± 0.03
0.43 ± 0.03
0.50 ± 0.04
1.61 ± 0.07
3.03 ± 0.08
4.36 ± 0.11
5.78 ± 0.12
6.93 ± 0.13
7.37 ± 0.14
8.88 ± 0.15
9.79 ± 0.15
10.83 ± 0.16
12.12 ± 0.17
12.47 ± 0.18
13.88 ± 0.17
14.96 ± 0.19
16.92 ± 0.20
19.29 ± 0.21
20.29 ± 0.22
21.19 ± 0.23
23.77 ± 0.24
28.26 ± 0.26
33.12 ± 0.28
41.34 ± 0.32
fitted brightness of the (14, 3) band of 56 R observed during record
25 at 11°S and 16:15 LT, we find a slant CO column of
9.5 1015 cm2. We estimate the rms uncertainty on the CO column to be 14%, due to propagation of the Poisson noise. Our
observation corresponds to a vertical column of 6.4 1015 cm2.
This value is close to 8.7 1015 cm2, the vertical column of CO
above the altitude of unit optical depth at the wavelength of the
(14, 3) band provided by the VTS3N model for the conditions of
the UVIS observation. This column is smaller than the disk-averaged value of 3.4 1016 cm2 found by Feldman et al. (2000) for
solar minimum conditions. We note that the wavelength interval
covered by the (14, 3) (k 131.6 nm) band at the UVIS resolution
is contaminated by the (10, 1) (k 131.81 nm) and the (12, 2)
(k = 131.59 nm) CO 4P bands, not accounted for in the estimate
of Feldman et al. (2000), but taken into account in our fitting procedure. Moreover, this band is close to the very bright OI 130.4 nm
emission, with a possible additional contamination. If we neglect
any contamination and simply integrate the spectrum over the
(14, 3) wavelength interval, we find a slant column of
4.8 1016 cm2 and a vertical column of 3.2 1016 cm2, in better agreement with the column given by Feldman et al. (2000). Indeed, comparing the fitted intensity of the (14, 3) band and the
intensity obtained by straightforward integration of the (14, 3)
wavelength interval, we find that the (14, 3) band itself contributes
only 17% to the observed feature, all the rest being attributed to
contamination according to our fitting method (for the spectral record number 25, i.e. for the observation having the smallest possible emission angle). We also note that the g-factor method is
approximate since the expression of the g-factor (Durrance,
1981) includes a dependence on the solar flux intensity and spectral distribution, and that, in addition, it does not explicitly account
for multiple scattering and self-absorption, which could be important for transition involving the ground state. The g -factors can
thus in principle not be considered as natural constants, unlike
the usual atomic parameters such as the Einstein coefficients. Ideally, the temperature profile of the atmospheric gas should be accounted for when computing g-factors. However, the
temperature may vary with time, and from place to place in the
upper atmosphere of Venus. The appropriate detailed temperature
profile of each observation is thus hard to determine as it depends
on the solar EUV flux and dynamical processes. Its determination is
beyond the scope of the present work.
Recently, several authors have revisited the transition parameters and wavelengths of the CO 4P system. Morton and Noreau
(1994) gave a more accurate set of wavelength and oscillator
strength for the (v0 , 0) transitions. Smith et al. (1994) measured
the oscillator strength of the (v0 , 0) transitions for 11 v0 14,
and Stark et al. (1998) extended these measurements to
8 v0 22. Eidelsberg et al. (1999) also measured the oscillator
strength of the CO 4P system for 11 v0 23 and v00 = 0. These
studies show significant differences with the previous values from
Kurucz (1976). However, the full set of v0 , v00 values are needed for
our least-squares fit procedure, so that we cannot consistently
adopt these recent results here. Indeed, the absolute values of
the transition probabilities are unimportant to our fitting method.
Only the relative intensity of the bands with v00 > 0 within a given
progression can influence our fitted intensities. However, the discrepancies between the Kurucz (1976) transition parameters and
the more recent ones suggest that our results may need revisions
when more accurate values will be known for the full CO 4P
3.4. The OI emissions
The solar UV flux is obtained from the SOLAR2000 (Version
2.27) empirical model (Tobiska, 2004) giving the solar flux at Earth
distance for a given date. We account for the angle between Venus,
the Sun and the Earth to correct for the solar rotation and the
resulting delay between the solar photon flux reaching the Earth
and Venus. The solar flux is then corrected for the distance between the Sun and Venus. The composition of the Venus upper
atmosphere is obtained from the VTS3 empirical model (Hedin
et al., 1983) which provides neutral densities (CO2, O, CO, N2, He,
and N) and temperature. This model was based on measurements
of the main neutral constituents in the thermosphere performed
above 145 km with the mass spectrometer on board Pioneer Venus. The numerical formulation relies on modified Bates temperature profiles and the related diffusive equilibrium density profiles.
The numerical model used to calculate the photoelectron production and energy degradation in the Venus atmosphere was described by Shematovich et al. (2008) and Gérard et al. (2008).
Energetic electrons are produced by photoionization of the major
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 ionization collisions with the ambient atmospheric gas. Cross sections
used to calculate the energy loss associated with inelastic collisions
with CO2, CO and O were listed by Shematovich et al. (2008). The
energetic electrons lose their excess kinetic energy in collisions
with the atmospheric particles. The Direct Simulation Monte Carlo
(DSMC) method is used to solve atmospheric kinetic systems in the
stochastic approximation. The evolution of the system of modeling
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. The region of the atmosphere under study is divided into
49 vertical cells. The excitation rates of the O(3S) and O(5S) states
by electron impact are then directly calculated using the calculated
energy distribution function, the O density distribution and the relevant excitation cross sections.
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 solar
B. Hubert et al. / Icarus 207 (2010) 549–557
130.4 nm flux is obtained from the model of Woods and Rottman
(2002) that sets up a proxy relating the solar UV flux and the
F10.7 index. The shape of the solar 130.4 nm lines is taken from
Gladstone (1992). 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 this study we use angle-averaged partial frequency redistribution. In addition, the 130.4 nm photons can be absorbed by
CO2, especially at altitudes below the peak of the O(3S0) emission.
A vertical unit optical depth for 130.4 nm absorption by CO2
reaches unity near 130 km. The effects of spherical geometry are
included in the radiative transfer code to calculate the photon slant
optical paths. Consideration of sphericity becomes important for
viewing and solar zenith angles larger than 70°. Cascades from
upper lying levels feed the O(3p 3S0) state, mostly through the 3p
3S0–3s0 3D0 transition at 844.6 nm and may contribute to enhance
the 130.4 nm emission. The efficiency of these cascades and their
complexity due to radiative transfer were discussed in detail by
Meier (1991). As indicated by Alexander et al. (1993), the total cascade contribution from these additional sources amounts to 15%
over the optically thin excitation rate. This contribution has been
taken into account by using the ‘‘optically thick” Zipf and Erdman
(1985) cross sections.
Fig. 4a shows the OI 130.4 nm brightness measured along the
footprint of the UVIS line of sight. The root mean square uncertainties are generally below 3%. We now compare the observed OI
130.4 nm intensity with model calculations for the geometry
appropriate for each UVIS spectral record. The dashed line in
Fig. 4a shows the calculated 130.4 nm intensity in the UVIS viewing direction when only the electron impact primary source is considered. It clearly appears that an additional primary source of
photons needs to be included. Resonance scattering of the solar
130.4 nm multiplet provides an important contribution. The total
calculated brightness including both primary sources of 130.4 nm
photons is close to the observations (dotted lines). The model overestimates the OI 130.4 nm brightness by less than 0.7 kR for nearly
all the observations. We note that, due to geometric complications,
we cannot reliably simulate the 130.4 nm brightness near the limb
and the terminator. On average, the relative error is 11%, ignoring
the poorly modeled limb and terminator values. This difference is
quite small, considering that we use a value for the solar
130.4 nm flux that is based on a proxy, and that the F10.7 proxy value is based on observations made from Earth. We find, in agreement with Gérard et al. (2008), that the solar 130.4 nm flux is
the dominant primary source of photons feeding radiative transfer
of the 130.4 nm oxygen multiplet in the Venus atmosphere.
The The brightness of the OI 135.6 nm emission deduced from
our least-squares fitting method, which separates the CO 4P
(14, 4) band contribution from the 135.6 multiplet, is shown in
Fig. 4b. The 1 r rms uncertainties are below 6% for nearly all
the observations. The excitation of the O(5S0) state is solely due to
photoelectron impact on oxygen atoms. The vertical thickness at
the line center of the 135.6 nm strongest component is s 0.04 at
the level of unit optical depth for CO2 absorption. Therefore, multiple scattering of the O 3P–5S0 transition is neglected in this study.
The O(5S0) production rate is simulated with the model for the conditions of each UVIS record. The calculated volume emission rate is
then integrated along the UVIS line of sight to obtain the simulated
OI 135.6 nm brightness for each successive position of the UVIS slit
projection. The agreement between simulation and observation is
excellent (within a few percent), except for a viewing geometry
close to the limb, for the same reason as mentioned for the
130.4 nm emission. We find that the CO 4P (14, 4) band always contributes less than 25% to the brightness of the spectral feature at
135.5 nm in the UVIS observations. This result is in contrast with
that of Durrance (1981) who found the (14, 4) band contributes
Fig. 4. (a) Observed OI 130.4 nm intensity (solid line). The modeled OI 130.4 nm
emission is shown for comparison (dotted line). The contribution of the 130.4 nm
photons produced by the impact of electrons on oxygen to the calculated OI
130.4 nm intensity is indicated by the dashed line. (b) Observed intensity of the OI
135.6 nm emission following (solid line) and before (dotted line) removal of the CO
4P contribution. The modeled OI 135.6 nm emission is shown for comparison (long
dashes). (c) Observed (solid lines) and modeled (dotted line) OI 130.4/135.6 ratio.
The error bars in panel c indicate the 1 r uncertainties due to propagation of the
Poisson noise.
B. Hubert et al. / Icarus 207 (2010) 549–557
2/3 to the 135.6 nm feature, based on IUE observations. Different
atmospheric composition for both observations and/or different solar UV spectral distribution modifying the relative importance of
the primary sources of photons for both emissions may account
for the discrepancy. An inaccurate fitting of the brightness of both
emissions cannot be ruled out either. The 1 r rms uncertainty
on the fitted (14, 4) intensity remains below 16% for most of the
observations, although it is larger than 25% for very dim spectra obtained close to the terminator. The good consistency that we find
between the observed and modeled OI 135.6 nm emission suggests
that both the model and the analysis of the observations are
The ratio of the brightness of the 130.4 and 135.6 nm emissions
can be used as further check of the validity of the consistency between observations and simulations. The ratio deduced from the
UVIS observations is plotted in Fig. 4c. It is generally larger than
that deduced from the HUT observations (Feldman et al., 2000)
which is 4.6. It is however less than the value of 8 found by Durrance (1981). We find that the agreement between the modeled
and observed ratio varies along the track of the UVIS line of sight.
It is much better with observations obtained away from the terminator, i.e. when the intensity is high, and away from the limb, i.e.
when the geometry is sufficiently simple. One possibility is that
the atmosphere used in the simulations corresponds to the actual
situation prevailing in the Venus atmosphere only for the locations
of the second half of the flyby. Uncertainties on the solar flux may
also contribute to the discrepancy. Setting the relative contribution
of OI 135.6 nm line to 1/3 of the observed feature obtained by Durrance (1981) would reduce our estimated OI 135.6 nm brightness,
which would increase the OI-130.4/OI-135.6 ratio by, roughly, a
factor 2–3. This would not really improve the agreement between
modeling and observations. The logical conclusion from the good
agreement between the observed and modeled OI 130.4 nm and
135.6 nm emissions as presented in Fig. 4a and b is that the oxygen
density profile deduced from the VTS3 model and used in our model is consistent with the UVIS observations, despite the discrepancy
for the ratio of these two emissions for observations close to the
3.5. Comparison with HUT disk intensity observations
It is of interest to further compare the present measurements
with preceding observations of comparable spectral range and resolution by Feldman et al. (2000). They described observations obtained with HUT from the Space Shuttle Astro 2 mission on 13
March 1995. Venus was then at a western elongation of 40° and
a phase of 60°. The phase angle values were close to 99° during
the UVIS observations, but the solar activity was low for HUT
and very high during the UVIS flyby. We estimate the appropriate
values for the F10.7 index to be 82 and 214 respectively. Feldman
et al. reported the disk-averaged brightness of 13 emissions identified within their spectral range of 82–184 nm. The HUT spectrum
integrated the dayglow over the sunlit fraction of the disk, whereas
UVIS observed only a relatively narrow strip of Venus running from
the dusk terminator to the vicinity of the subsolar limb. Therefore,
for comparison of the UVIS spectra with the HUT data, we estimate
the equivalent full-disk brightness of a given emission observed
with UVIS by fitting to the UVIS data a simple function of the illumination and emission cosines, l = cos h and l0 = cos h0, where h is
the illumination angle and h0 the emission angle. We adopt a function which has the form:
B ¼ B00 ððl0 þ aÞ=ð1 þ aÞ > 0Þp =ðl < bÞp > 0
where 0 p , p 1 determine the variation of the signal across the
disk. The quantity a simulates the extension of an optically thick
Table 4
Disk brightness of several spectral features of the Venus FUV airglow observed by the
HUT spectrometer and deduced from UVIS observations.
k (nm)
UVIS/HUT ratio
OI, CO A–X (14, 4)
CO A–X (0, 1)
emission beyond the terminator, while b effectively simulates the
smearing of the signal near the limb by the motion of the UVIS field
of view. B00 is a normalization factor equivalent to the subsolar nadir brightness. The equivalent disk-averaged brightness can then be
computed by integrating over the disk at the appropriate phase angle. Table 4 compares these estimates with the HUT results, together with their ratios. The ratios range between 1.23 and 1.64,
which is not inconsistent with the difference in solar activity during
HUT and UVIS observations. It is beyond the scope of the present
work to carry these comparisons any further, except as described
above in this text.
4. Summary and conclusions
We have analyzed observations of the OI and CO fourth positive
dayglow obtained with UVIS during the Venus flyby by Cassini.
These observations provide high quality spatially-resolved spectra
extending from the evening terminator to the vicinity of the subsolar point. A least-squares fit method is used to determine the
brightness of the individual spectral features present in the observed spectra. The CO 4P brightness peaks at 80 kR at the limb,
in good general agreement with previous observations. The corresponding CO vertical column is 9.5 1015 cm2, a value close to
that provided by the VTS3 model.
We compare the observed intensity of the OI emissions at
135.6 nm and 130.4 nm with those calculated using an airglow
model. In this model the photoelectron impact contribution is calculated based on a Monte Carlo method solving the steady state
photoelectron energy distribution function. We estimate the contribution of the CO 4P (14, 4) band to the OI 135.6 nm doublet to
be is less than 25%. Once this contribution is removed, the OI
135.6 nm brightness deduced from the observations is well reproduced by numerical simulations of this emission for the specific
conditions of the UVIS observations. The OI 130.4 nm multiplet is
predominantly excited by photoelectron impact on O atoms and
resonance scattering. Multiple scattering in the optically thick Venus atmosphere with partial frequency redistribution has been
considered. The observed and modeled OI 130.4 nm brightness
agree with each other within 11% on average. The ratio of the observed OI 130.4 to the OI 135.6 nm emissions lies within the values
previously reported in the literature.
The general agreement between the observed and modeled
brightness of the OI 130.4 and OI 135.6 nm emissions lead to the
conclusion that the O density profile from the VTS3 model by Hedin et al. (1983) gives values of the oxygen density profile in the
upper atmosphere in agreement with the UVIS observations of
the OI airglow. A similar conclusion was reached by Paxton and
Meier (1986) from the analysis of PV–OUVS observations of the
130.4 nm dayglow. They found that the O densities used to fit
the OUVS observations were consistent with the measurements
performed with Pioneer Venus bus neutral mass spectrometer.
The VTS3 model largely relies on the database collected near
periapsis with the Orbiting Neutral Mass Spectrometer (ONMS)
on board the Pioneer orbiter. The altitude of the emission peak of
the OI emissions is close to 140 km, only few kilometers below
B. Hubert et al. / Icarus 207 (2010) 549–557
the lowest altitude probed with the ONMS instrument. Since these
airglow observations were made at low latitude where ONMS data
were collected, it is expected that the VTS3 model provides a realistic extrapolation of the observed densities.
This study is based on observations by the Cassini project. B.
Hubert and J.-C. Gérard are supported by the Belgian National Fund
for Scientific Research (FNRS). Funding for this research was managed by the PRODEX program of the European Space Agency in collaboration with the Belgian Science Policy Office and by FRFC Grant
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