Planetary and Space Science 59 (2011) 1524–1528
Contents lists available at ScienceDirect
Planetary and Space Science
journal homepage: www.elsevier.com/locate/pss
Measurements of the helium 584 Å airglow during the Cassini flyby of Venus
J.-C. Gérard a,n, J. Gustin a, B. Hubert a, G.R. Gladstone b, L.W. Esposito c
a
b
c
Laboratoire de Physique Atmosphérique et Planétaire, Université de Lie ge, 17, allée du 6 août, B5c, B-4000 Lie ge, Belgium
Southwest Research Institute, P.O. Drawer 28510, 6220 Culebra Road, San Antonio, TX 78238-5166, USA
Laboratory for Atmospheric and Space Physics, University of Colorado, Campus Box 392, Boulder, CO 80309-0392, USA
a r t i c l e i n f o
a b s t r a c t
Article history:
Received 24 February 2011
Received in revised form
17 June 2011
Accepted 21 June 2011
Available online 7 July 2011
The helium resonance line at 584 Å has been observed with the UltraViolet Imaging Spectrograph
(UVIS) Extreme Ultraviolet channel during the flyby of Venus by Cassini at a period of high solar
activity. The brightness was measured along the disk from the morning terminator up to the bright
limb near local noon. The mean disk intensity was 320 R, reaching 700 R at the bright limb. These
values are slightly higher than those determined from previous observations. The sensitivity of the
584 Å intensity to the helium abundance is analyzed using recent cross-sections and solar irradiance
measurements at 584 Å. The intensity distribution along the UVIS footprint on the disk is best
reproduced using the EUVAC solar flux model and the helium density distribution from the VTS3
empirical model. It corresponds to a helium density of 8 106 cm 3 at the level of where the CO2 is
2 1010 cm 3.
& 2011 Elsevier Ltd. All rights reserved.
Keywords:
Venus
Helium density
Airglow
Cassini
Flyby
Radiative transfer
1. Introduction
Helium atoms in the Venus atmosphere originate from outgassing following radioactive decay of uranium and thorium in
the planetary crust or from the capture of solar wind a-particles.
The helium budget on Venus has been discussed by Pollack and
Black (1982), Prather and McElroy (1983), Chassefie re et al.
(1986), and Krasnopolsky and Gladstone (2005). The thermal
escape rate is negligibly small and he main loss process is
ionization of helium atoms above the ionopause followed by
interaction with the solar wind sweeping out He þ ions. The
efficiency of this loss process depends on the helium density in
the upper atmosphere. Krasnopolsky and Gladstone (2005)
derived a mixing ratio of 976 ppmV in the middle and lower
atmosphere and concluded that the efficiency of the solar wind
a-particles capture is 0.1. Barabash et al. (2007) detected with the
Aspera-4 instrument on board Venus Express the presence of He þ
ions escaping through the planetary plasma wake. They measured
an unexpectedly high He þ relative abundance, possibly due to an
effective acceleration of the ions caused by the polarization field.
The importance of the knowledge of the helium density is linked
to the planetary evolution since the present helium abundance
n
Corresponding author. Tel.: þ32 4 3669775; fax: þ32 4 3669711.
E-mail addresses: jc.gerard@ulg.ac.be (J.-C. Gérard),
J.Gustin@ulg.ac.be (J. Gustin), randy.gladstone@swri.org (G.R. Gladstone),
larry.esposito@lasp.colorado.edu (L.W. Esposito).
0032-0633/$ - see front matter & 2011 Elsevier Ltd. All rights reserved.
doi:10.1016/j.pss.2011.06.018
depends on the relative rates of the outgassing and He þ capture
sources during the evolution of the planet.
The first detection of helium in the Venus atmosphere was made
with the ultraviolet spectrometer on board Mariner 10. From the
bright limb profile of the HeI resonance line at 584 Å, Kumar and
Broadfoot (1975) derived a helium density of (271) 106 cm 3 at
the altitude where the CO2 density is 2 1010 cm 3. Following
Krasnopolsky and Gladstone (2005), we adopt this altitude where
the CO2 density is 2 1010 cm 3 as the reference level for comparison of the various measurements. It is located near 145 km according
to the VTS3 empirical model (Hedin et al., 1983) and corresponds to a
vertical optical depth at 584 Å of 0.5. If a more realistic thermospheric temperature of 275 K than the 350 K used in the original
analysis is adopted, the density must be increased by a factor of 1.8
(von Zahn et al., 1983) and the Mariner 10 observations correspond to
a helium density of 3.6 106 cm 3 at the reference level. Further
observations of the HeI 584 Å emission were made with the EUV
spectrometers carried by the Venera 11 and Venera 12 spacecraft
during their Venus flyby in December 1978. A description of the
instrument and a preliminary analysis was given by Bertaux et al.
(1981). The observed intensity, following background subtraction,
reached a peak value of 280 R. A detailed analysis by Chassefie re et al.
(1986) based on the shape of the scans of the illuminated disk
concluded that the optical thickness of the 584 Å emission above the
CO2 absorption level was 3.571.5, independently of the instrumental
calibration. This optical depth corresponds to a helium density of
(2.671.2) 106 cm 3 at the reference level. Krasnopolsky and
Gladstone (2005) observed the 584 Å line with the Extreme
J.-C. Gérard et al. / Planetary and Space Science 59 (2011) 1524–1528
Table 1
Values of the helium density measured at the reference levela in the Venus lower
thermosphere.
Spacecraft
He density (in
106 cm 3)
Reference
Mariner 10
Mariner 10b
Venera 11 and 12
EUV Explorer
27 1
3.6
2.67 1.2
4.37 1.9
Pioneer Venus
bus
Cassini flybyc
7.0
Kumar and Broadfoot (1975)
von Zahn et al. (1983)
Chassefie re et al. (1986)
Krasnopolsky and Gladstone
(2005)
von Zahn et al. (1980)
8.07 4
This work
a
b
c
Where [CO2] ¼ 2 1010 cm 3.
With corrected temperature profile.
Value deduced using the EUVAC solar flux model.
Ultraviolet Explorer (EUVE) satellite during an accumulated 16,000 s
period in April 1998. During these observations, the sunlit fraction of
the Venus disk was 55.5% and the solar activity was moderate (F10.7
index¼113). The observed integrated disk brightness at 584 Å was
144732 R. They inferred a helium density of (4.371.9) 106 cm 3
at the reference level. Using a radiative transfer model, they calculated the He 584 Å brightness distribution over the Venus sunlit disk
for the conditions of the EUVE observations. It shows a significant
limb brightening along the subsolar meridian and a continuous
increase from the morning terminator to the subsolar meridian.
Chassefie re et al. (1986) analyzed how the presence of a limb
brightening depends on the interplay between the optical thickness
of the 584 Å emission and the level of CO2 absorption. They showed
that for a helium optical thickness of 1 above the CO2 absorption
level, a clear center to limb brightening is present. Instead, for
t ¼ 7.56, a limb darkening was predicted by their model. The
calculated 584 Å intensity increased everywhere on the disk as the
helium optical thickness increased.
In situ measurements of the helium density were also
performed during the Pioneer Venus mission. The bus carrying a
mass spectrometer entered the atmosphere on the morning side
where the solar zenith angle was close to 601 and measured the
chemical composition down to 130 km (von Zahn et al., 1980).
The measured helium density at the reference altitude was
6.8 106 cm 3. The neutral mass spectrometer (ONMS) on board
the Pioneer Venus orbiter provided He density measurement
down to 145 km at all solar times in a region confined to low
latitude (Niemann et al., 1980). The measurements showed little
variation as a function of local time on the dayside, but indicated
the presence of a pronounced bulge centered near 4:30 LT. This
bulge was interpreted as a consequence of the subsolar to
antisolar circulation and the presence of superrotation in the
upper atmosphere of the planet. The VTS3 model, which is mainly
driven by the ONMS measurements, predicts equatorial helium
density at the reference level varying from 1.8 107 cm 3 at
0600 LT to 8 106 cm 3 during daytime. These helium density
values derived from earlier measurements are summarized in
Table 1.
2. The observations
The Cassini spacecraft flew by Venus on 24 June 1999 on its
way to Saturn, reaching a closest altitude of 602 km. A series of
simultaneous FUV and EUV spectra were collected with the
UltraViolet Imaging Spectrograph (UVIS) (Esposito et al., 2004).
The UVIS line of sight was oriented nearly perpendicular to the
Sun-spacecraft line, so that the phase angle remained close to 901.
1525
During the observations, 55 records of 32 s each were obtained,
25 of which were recorded while the instrument scanned the
sunlit disk of Venus from the morning terminator to sunlit limb in
the vicinity of the subsolar point. The latitude of the UVIS slit
footprint on the planet varied from 241 north to 51 south. Figure
1 and Table 1 by Gérard et al. (2011) show the foot track
geometry and describe the variation of the solar zenith angle
(SZA) and emission angle (the angle between the line of sight and
local zenith at the altitude of airglow emission) at each record. At
this period, solar activity was rising, and reached a F10.7 solar
index of 212 at Earth distance, accounting for the difference in
solar longitude between the Earth and Venus. Observations in the
FUV and in the EUV down to 834 Å made with UVIS during the
Cassini flyby have been described and analyzed by Hubert et al.
(2010) and Gérard et al. (2011).
The bandpass of the UVIS EUV channel covers the range
563–1182 Å. The two-dimensional CODACON detector allows
spectral and one-dimensional spatial coverage as the footprint
of the instrument scans the planetary disk. The UVIS slit image on
the detector is composed of 1024 pixels in the dispersion direction and 64 pixels in the spatial direction. The slit was oriented
nearly perpendicular to the ecliptic plane. The full spectral
resolution has been used during the Venus observations, providing spectra at a resolution of 3.7 Å FWHM. The spatial direction
has been rebinned over 16 pixels, leaving a resolution of 4 pixels
along the spatial direction. Each record presented here is the sum
of the two central spatial pixels in order to increase the signal/
noise ratio. The EUV field of view along the slit is 59 mrad,
corresponding to 415 km projected on the planet surface from
an altitude of 7000 km. Consequently, only a small fraction of the
disk is seen during each record, causing negligible smoothing,
except in the vicinity of the planetary limb. The spacecraft moved
500 km during the 32 s integration period of each record.
The EUV channel was initially calibrated following the preflight laboratory measurements described by Esposito et al.
(2004). The primary standards used for determining the UVIS
absolute radiometric sensitivity were photodiodes provided by
the National Institute for Standards and Technology (NIST).
Measurements of the star Alpha Virginis (Spica) obtained during
cruise in January 1999 agreed with previous results to within 10%
for both EUV and FUV channels, validating the UVIS laboratory
results. An empirically derived background noise level of
4.5 10 4 count s 1 pixel 1 originating from the radioisotope
thermoelectric generators has been removed and a flat-field
correction derived from observations of Spica (Steffl et al., 2004)
has been applied. A contaminating source also affects wavelengths below 920 Å (known as ‘‘the mesa’’). It is caused by a
small light leak allowing undispersed interplanetary Lyman-a
photons to reach the portion of the EUV detector corresponding to
short wavelengths. The signal associated with the Lyman-a leak
smoothly rises from 0.063 count s 1 pixel 1 at 560 Å to
0.125 count s 1 pixel 1 at 920 Å, rapidly dropping to zero at
1020 Å. Consequently, the residual spectrum in this region is
quite noisy, preventing reliable detection of weak emissions
below 930 Å, with the exception of the bright HeI emission at
584 Å. This background signal has been manually subtracted from
each individual spectrum to determine the brightness of the
584 Å line across the disk. The average signal of the 584 Å
emission is about 60 counts/spectrum, to be compared with the
135 counts/spectrum from the Lyman-a leak. The count rate has
been converted into physical units using the latest UVIS calibration routines. Fig. 1 shows the intensity distribution of the 584 Å
line measured along the slit scan of the Venus disk during the
flyby. The statistical 1-s error bars are indicated for each record.
The main source of uncertainty comes from the subtraction of the
signal due to the Lyman-a leak.
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J.-C. Gérard et al. / Planetary and Space Science 59 (2011) 1524–1528
HeI 584−Å intensity (R)
1000
UVIS
EUVAC solar flux
Woods & Rottman solar flux
Sensitivity to [He]
800
[He] X e
[He] X 2
600
[He] / 2
400
[He] / e
200
0
97
97
94
88
79
73 64 57
SZA (deg)
47
39
26
11
00
Fig. 1. Measurements of the 584 Å airglow intensity along the UVIS foottrack on
the Venus disk (diamonds). The thin vertical bars show the statistical 1-s standard
deviation. The solid line shows the modeled intensity with the helium density
from the VTS3 model and the solar flux from the EUVAC model. The dotted lines
show the modeled intensities calculated by scaling the helium density by factors
of e 1, 0.5, 2 and e. The dashed line represents the calculated intensity distribution with the solar flux from Woods and Rottman (2002).
The observed emission brightness starts rising near the morning terminator, increasing steadily toward the sunlit limb where
it reaches a peak of about 800 R. The mean disk intensity,
excluding the last 4 points near the bright limb is 319 711 R.
These intensity values exceed any previous observation of the
584 Å Venus dayglow. They are larger than the 280 R maximum
value observed with Venera 11 and 12 and are about a factor of
2 higher than the 144 R average disk brightness observed with
EUVE. We note however that the Cassini flyby occurred during a
period of very high solar activity (F10.7¼212) in comparison with
the Venera (F10.7¼153 and 173) and EUVE (F10.7¼113) observations. In the next section we discuss how solar activity affects
the brightness of the He solar emission line.
3. Comparison with model
To determine the He abundance in the Venus upper atmosphere, we now compare the observed variation across the disk of
the 584.33 Å intensity to the calculated distribution based on the
VTS3 empirical model. The 1s2 1S–1s2p 1P transition is excited by
resonance scattering of HeI chromospheric solar line by atmospheric helium atoms. This transition is moderately optically
thick above the reference level. The oscillator strength of the
transition is taken equal to 0.276 (Wiese and Fuhr, 2009). Carbon
dioxide is the dominant absorber at 584 Å with a CO2 absorption
cross-section of 3.66 10 17 cm2 (Shaw et al., 1995), recommended by Huestis and Berkowitz (2010). The number densities
of CO2 and He for the baseline model are provided by the VTS3
empirical model (Hedin et al., 1983) mostly based on mass
spectrometer measurements made at low latitude during a period
of high solar activity. Multiple scattering is calculated using the
radiative transfer code described by Gladstone (1985) based on
the Feautrier method. In this study we use angle-averaged partial
frequency redistribution. 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. The role of spherical geometry becomes important for viewing and solar zenith angles larger than 701.
Although the radiative transfer code is plane-parallel, spherical
effects are correctly accounted for in calculating the initial source
function and the final intensities.
The absolute flux near the center of the solar He 584 Å line is a
key quantity controlling the airglow intensity. The Doppler width
at the temperature of the Venus upper dayside atmosphere
(TE 250 K) is much smaller than that of the chromospheric
emission line. A high-resolution measurement of the solar line
profile in the second order of the Solar Ultraviolet Measurements
of Emitted Radiation (SUMER) instrument on board the SOHO
satellite indicates that the line shape is slightly flattened Gaussian
with a full-width at half-maximum (FWHM) of about 155 mÅ
(Wilhelm et al., 1997). This result is in agreement with the earlier
observations such as the determination by Maloy et al. (1978)
based on measurements with helium-filled gas cells. Some
observations suggest that the width of the solar line varies to
some extent with the level of solar activity (Ogawa et al., 1984)
but the statistics are too poor to draw definite conclusions. The
Doppler FWHM of the atmospheric absorption line ( 3 mÅ) is
much less than the solar line so that the actual profile of the solar
line does not influence the vertical distribution of the 584 Å
airglow. The calculated intensity of the resonance line is, in any
observing geometry, proportional to the solar flux at the line
center. The integrated flux of the solar line may be obtained from
different empirical models setting up proxies relating the solar UV
flux and the F10.7 index. The EUVAC model (Richards et al., 1994)
is based on the PF10.7 Index combining the daily and the 80-day
average F10.7 indices. It predicts a flux of 1.85 109 ph cm 2 s 1.
The SOLARVUV model by Woods and Rottman (2002), which
uses the F10.7 cm index as a proxy, gives a value of 1.36 109 ph cm 2 s 1. The solar flux at 584 Å has been observed in the
second diffraction order with the SUMER spectrometer on board
SOHO at a solar activity level corresponding to F10.7¼ 67
(Wilhelm et al., 1997). The solar flux at the line center was
5.2 109 ph cm 2 s 1 Å 1, corresponding to an integrated flux of
0.85 109 ph cm 2 s 1. For comparison, the EUVAC model predicts a flux of 1.18 109 ph cm 2 s 1 for this solar activity level,
that is 39% higher than the SUMER value. Another estimate of the
solar flux can also be made based on observations made with the
Solar EUV Experiment (SEE) on board the Thermosphere, Ionosphere, Mesosphere, Energetics and Dynamics (TIMED) satellite.
Although these measurements only started in 2002, some SEE
measurements were found to correspond to values of the F10.7
and F10.7A indices identical to those of the Cassini Venus flyby.
The 584 Å fluxes are 1.4 109 ph cm 2 s 1, 19% higher than the
EUVAC value. We adopt the EUVAC model as our baseline value.
For larger values of the solar flux, the derived He density is less, as
will be discussed hereafter.
We now compare the observed intensity with the model
calculations for the geometry appropriate for each UVIS spectral
record. The radiative transfer model provides the intensity integrated along the line of sight for each UVIS record in the geometry
of the UVIS pointing during the flyby. However, we note that at
the limb the observed signal is averaged over the size of the
projected UVIS slit. This effect is not accounted for in the
comparisons presented here. Fig. 1 illustrates this comparison
based on the helium density profile from the VTS3 model and the
EUVAC solar flux. The agreement is excellent (correlation coefficient R ¼0.95), considering the various sources of uncertainties.
The shape of the brightness distribution across the disk is well
reproduced by the model. If the SOLARVUV solar intensity is used,
which is nearly identical to renormalizing the EUVAC flux to the
SUMER low solar activity observation, the calculated curve underestimates the observed brightness by 35%, the difference
between the EUVAC and the SOLARVUV solar flux values. In this
case, a larger helium density is required to best match the
observations. The vertical distribution of the primary resonance
J.-C. Gérard et al. / Planetary and Space Science 59 (2011) 1524–1528
scattering source calculated for the geometry of UVIS record 25
(solar zenith angle (SZA)¼ 64.21, emission angle (EMA)¼47.11) is
shown in Fig. 2 by the solid line. The dotted line shows the
optically thick source function, which accounts for multiple
scattering and represents the number of photons emitted
per second in a unit volume. The primary source is maximum at
250 km and drops at lower altitude as a result of the decrease of
the solar photon flux scattering away from the beam by the
overlying helium atoms. The optically thick source function
instead peaks near 145 km, the altitude where CO2 begins to
significantly absorb the 584 Å photons. The sensitivity of the
emerging intensity to the helium column has been investigated
by applying scaling factors to the VTS3 He density. Fig. 3 illustrates the intensity calculated for UVIS record 25 with the EUVAC
600
Altitude (km)
500
400
Source function
1527
solar flux and a range of helium densities obtained by multiplying
the VTS3 model by a factor ranging from e 1 to 10. Interestingly,
the dependence on the helium abundance is quasi linear for small
values, but departs from a linear dependence when the density
increases. This behavior is explained by relative altitude of the
unit optical depths of the HeI line and the CO2 absorption. As long
as the optical depth of the He column overlying the CO2 cutoff
altitude ( 145 km) is small enough, the t ¼1 level for the 584 Å
line is below the cutoff altitude, the emission is optically thin
and its emerging intensity remains roughly proportional to the
column abundance. However, once the column becomes optically
thick, the altitude of the t ¼1 level moves above the CO2 cutoff
and the integrated brightness becomes less sensitive to the
amount of helium, as observed in Fig. 3.
Fig. 1 also shows the 584 Å intensity calculated along the UVIS
slit track for a helium density multiplied by 0.5, e 1, 2 and e.
A multiplication (division) by e corresponds to rising (lowering)
the helium column by one scale height. If the SOLARVUV flux
value is used instead of EUVAC, the best overall agreement with
the UVIS observations is reached when the VTS3 density is
multiplied by a factor of 2, corresponding to a helium density of
1.6 107 cm 3 at the reference level. The main sources of
uncertainties in the comparison between the data and the model
include (a) the statistical error on the signal (shown by the 1-s
error bars in Fig. 1), (b) the absolute calibration of the UVIS
instrument, and (c) the absolute solar flux at the center of the
584 Å line (estimated as 20–30%).
300
4. Conclusions
200
Primary source
100
0.1
1.0
10.0
100.0
-3 -1
Emission rate (ph cm s )
Modeled HeI 584−Å intensity (R)
Fig. 2. Vertical distribution of the primary source (resonance scattering) of 584 Å
photons and optically thick source function for the baseline model (see text).
500
SZA = 64.2°
EMA = 47.1°
400
Observed intensity
300
200
0
2
4
6
[He]/[He]VTS3
8
10
Fig. 3. Modeled intensity of the 584 Å emission for the observing geometry of
record a 25 of the UVIS measurements as a function of the scaling factor applied
to the helium density profile from the VTS3 model. The observed intensity is
360 R and the gray zone represents the 1-s uncertainty. The dashed lines
indicate the central and the extreme He densities compatible with the intensity
measured for this record. The vertical dotted line indicates e times the VTS3 value.
The 584 Å emission is only moderately optically thick above
the altitude of significant CO2 absorption which is located near
140 km. The 584 Å unit optical depth for a vertical observation is
reached at 182 km. The 584 Å optical depth at the reference level
in our best match is t ¼ 3.1. This value is in close agreement with
the 3.5 71.5 derived from the shape of the disk brightness
intensity distribution observed from Venera 11 and 12.
The analysis of the Venus disk scan made with the UVIS
instrument at 584 Å during the Cassini flyby shows that the
intensity increases from the morning terminator to a maximum
value of about 700 R at the sunlit limb. Using the solar line
intensity from the EUVAC model, we find that the helium
densities providing the best-fit to observed brightness distribution across the disk are in full agreement with those given by the
VTS3 empirical model. Although the observed brightness is larger
than previously observed from space-borne platforms, it is well
reproduced by the model because the intensity of the solar
chromospheric line significantly varies with the level of solar
activity. If the somewhat lower solar line irradiance based on the
SUMER quiet sun spectrum (or the SOLARVUV model) is adopted,
the best agreement is reached for twice as much helium as given
by the empirical VTS3 model. The dayside helium density at the
reference altitude derived from the UVIS observations is close to
8 106 cm 3. Combining the uncertainty of the absolute instrumental calibration and the statistical error of the measurements
with the calculated sensitivity of the intensity to the helium
density, we estimate the error bar on the He density as 50%. As
shown in Table 1, the He density is somewhat larger than those
derived from the analysis of the Venera and EUV Explorer airglow
observations, but in good agreement with the densities measured
with the mass spectrometers on board the Pioneer Venus bus and
orbiter. It appears that the main sources of uncertainty for the
determination of the helium abundance in the thermosphere are
the absolute instrumental calibration and the solar flux emitted in
the HeI solar line at the time of the observations.
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J.-C. Gérard et al. / Planetary and Space Science 59 (2011) 1524–1528
Acknowledgments
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
provided by the PRODEX program of the European Space Agency
in collaboration with the Belgian Science Policy Office and by
FRFC Grant #2.4541.11.
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