A&A 386, 319-330 (2002)
DOI: 10.1051/0004-6361:20020097
K. Dennerl 1 - V. Burwitz 1 - J. Englhauser 1 - C. Lisse 2 - S. Wolk 3
1 - Max-Planck-Institut für extraterrestrische Physik,
Giessenbachstraße, 85748 Garching, Germany
2 -
University of Maryland,
Department of Astronomy, College Park, MD 20742, USA
3 -
Chandra X-Ray Center,
Harvard-Smithsonian Center for Astrophysics,
60 Garden Street, Cambridge, MA 02138, USA
Received 28 September 2001 / Accepted 17 January 2002
Abstract
On January 10 and 13, 2001, Venus was observed for the first time
with an X-ray astronomy satellite. The observation, performed with the
ACIS-I and LETG/ACIS-S instruments on Chandra, yielded data of
high spatial, spectral, and temporal resolution. Venus is clearly
detected as a half-lit crescent, with considerable brightening on the
sunward limb. The morphology agrees well with that expected from
fluorescent scattering of solar X-rays in the planetary atmosphere.
The radiation is observed at discrete energies, mainly at the
O-K energy of 0.53 keV. Fluorescent radiation is also
detected from C-K
at 0.28 keV and, marginally, from
N-K
at 0.40 keV. An additional emission line is indicated
at 0.29 keV, which might be the signature of the
C
transition in CO2 and CO.
Evidence for temporal variability of the
X-ray flux was found at the
level, with fluctuations by
factors of a few times indicated on time scales of minutes. All these
findings are fully consistent with fluorescent scattering of solar
X-rays. No other source of X-ray emission was detected, in
particular none from charge exchange interactions between highly
charged heavy solar wind ions and atmospheric neutrals, the dominant
process for the X-ray emission of comets. This is in agreement with
the sensitivity of the observation.
Key words: atomic processes - molecular processes - scattering - Sun: X-rays, gamma rays - planets and satellites: individual: Venus - X-rays: general
The first detection of X-ray emission from a planetary atmosphere came as a surprise, when unexpectedly high background radiation was observed in 1967 during a daytime stellar X-ray survey by a rocket (Grader et al. 1968). This radiation was correctly attributed to X-ray fluorescence of the Earth's atmosphere. Aikin (1970) estimated that the same process should also cause other planetary atmospheres to glow in X-rays, although the expected flux at Earth orbit would be extremely small and only detectable with sophisticated instrumentation. X-ray fluorescence of the Earth atmosphere, however, became a well-known component of the X-ray background of satellites in low-Earth orbit. Its properties and its impact on observations were studied in detail by Fink et al. (1988) and Snowden & Freyberg (1993).
The detection of unexpectedly bright X-ray emission from comets (Lisse et al. 1996; Dennerl et al. 1997; Mumma et al. 1997) has led to increased interest in X-ray studies of solar system objects. With its carbon and oxygen rich atmosphere, the absence of a strong magnetic field, and its proximity to the Sun, Venus represents a close planetary analog to a comet. Dissociative recombination of O2+ in the Venus ionosphere leads to a hot oxygen exosphere out to over 4000 km (Russell et al. 1985), resembling a cometary coma. To investigate the X-ray properties of Venus, we performed a pioneering observation with the X-ray astronomy satellite Chandra.
Orbiting the Sun at heliocentric distances of 0.718-0.728
astronomical units (AU), the angular separation of Venus from the Sun,
as seen from Earth, never exceeds 47.8 degrees. While the observing
window of imaging X-ray astronomy satellites is usually restricted to
solar elongations of at least
,
Chandra is the first such
satellite which is able to observe as close as
from the
limb of the Sun. Thus, with Chandra an observation of Venus with an
imaging X-ray astronomy satellite became possible for the first time.
The observation was scheduled to take place around the time of
greatest eastern elongation, when Venus was
away from the
Sun. At that time it appeared optically as a very bright (-4.4 mag),
approximately half-illuminated crescent with a diameter of
(Table 1).
obsid | date | time | exp time | instrument | r | ![]() |
phase | elong | diam |
2001 | [UT] | [s] | [AU] | [AU] | [![]() |
[![]() |
[
![]() |
||
2411 | Jan.10 | 19:32:47-21:11:55 | 5948 | LETG/ACIS-S | 0.722 | 0.734 | 85.0 | 47.0 | 22.7 |
2414 | Jan.10 | 21:24:30-23:00:26 | 5756 | LETG/ACIS-S | 0.722 | 0.734 | 85.0 | 47.0 | 22.8 |
583 | Jan.13 | 12:39:51-15:57:40 | 11869 | ACIS-I | 0.721 | 0.714 | 86.5 | 47.1 | 23.4 |
Obsid: Chandra observation identifier, exp time: exposure time,
r: distance from Sun, :
distance from Earth,
phase: angle Sun-Venus-Earth, elong: angle Sun-Earth-Venus,
diam: angular diameter.
![]() |
Figure 1: First X-ray image of Venus, obtained with Chandra ACIS-I on 13 January 2001. |
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Venus is the celestial object with the highest optical surface brightness after the Sun and a highly challenging target for an X-ray observation with Chandra, as the X-ray detectors there (CCDs and microchannel plates) are also sensitive to optical light. Suppression of optical light is achieved by optical blocking filters which, however, must not attenuate the X-rays significantly. The observation was originally planned to use the back-illuminated ACIS-S3 CCD, which has the highest sensitivity to soft (E < 1.4 keV) X-rays, for direct imaging of Venus, utilizing the intrinsic energy resolution for obtaining spectral information. Before the observation was scheduled, however, it turned out that the optical filter on this CCD would not be sufficient for blocking the extremely high optical flux from Venus. Therefore, half of the observation (obsid 583, cf.Table 1) was performed with the front-illuminated CCDs of the ACIS-I array (I1 and I3), which are less sensitive to soft X-rays, but which are also significantly less affected by optical light contamination.
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Figure 2:
Spatial distribution of photons around Venus in the ACIS-I
observation. a) All photons in the energy range
0.2-10.0 keV; those with
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In order to avoid any ambiguity in the X-ray spectrum due to residual
optical light, the other
half of the observation (obsid 2411 and 2414, cf.Table 1)
was performed with the low
energy transmission grating (LETG), where the optical and X-ray
spectra are completely separated by dispersion. To minimize the risk
of telemetry overload, the area around the zero order image
(contaminated by optical light) was not transmitted to Earth. The
telescope was offset by
from its nominal aimpoint, to
shift the most promising spectral region around the dispersed 0.53 keV
O-K
emission line well into the back-illuminated CCD S3
(Fig. 3). The combination of direct imaging and
spectroscopy with the transmission grating made it possible to obtain
complementary spatial and spectral information within the available
total observing time of 6.5 hours.
At the time of the observation, Venus was moving across the sky
with a proper motion of
/hour. As the CCDs were read out
every 3.2 s, there was no need for continuous tracking. The spacecraft
was oriented such that Venus would move parallel to one side of the
CCDs and perpendicular to the dispersion direction in the LETG
observation. To keep Venus well inside the
field of view (FOV)
of ACIS-S perpendicular to the dispersion direction, Chandra was
repointed at the middle of the LETG/ACIS-S observation
(Table 1). For ACIS-I with its larger
FOV, no repointing during the observation was necessary.
As the photons were recorded time-tagged, an individual post-facto
transformation into the rest frame of Venus is possible. This was done
with the geocentric ephemeris of Venus as computed with the JPL
ephemeris calculator
.
Correction for the parallax of Chandra was done with the orbit
ephemeris of the delivered data set.
For the whole analysis we used events with Chandra standard grades and
excluded bright pixels. The fact that all observations were performed
with CCDs with intrinsic energy resolution made it possible to suppress
the background with high efficiency.
In the X-ray image (Fig. 1) the crescent of Venus is clearly resolved and allows detailed comparisons with the optical appearance. An optical image (Fig. 8e), taken at the same phase angle, shows a sphere which is slightly more than half illuminated, closely resembling the geometric illumination, with the brightness maximum well inside the crescent. In the X-ray image the sphere appears to be less than half illuminated. The most striking difference, however, is the pronounced limb brightening, which is particularly obvious in the surface brightness profiles (Fig. 2b) and in the smoothed X-ray image (Fig. 8d). For a quantitative understanding of this brightening we performed a simulation of the appearance of Venus in soft X-rays, based on fluorescent scattering of solar X-ray radiation. This simulation will be described in Sect.4.
The ACIS-I observation is well suited to investigate the X-ray
environment of Venus outside the bright crescent, an area unaffected
by optical loading. In Fig. 2 the distribution of photons
within a
region around Venus is shown, together
with the evolution of the surface brightness with increasing distance
from Venus, for a soft (E=0.2-1.5 keV) and a hard
(E=1.5-10.0 keV) band. In the soft band there is some
indication that the surface brightness at
-
decreases with increasing distance from Venus, both at the day- and
nightside. The fact that the corresponding distributions in the hard
band are flat argues against the possibility that this drop is caused
by inhomogeneities in the overall sensitivity. If the drop is not
caused by other instrumental effects, e.g., the PSF of the telescope,
then this could be evidence for an extended X-ray halo around Venus.
This will be discussed further in Sect.5.
The ACIS-I data clearly show that the X-ray spectrum of Venus is very
soft: images accumulated from photons with energies
show no enhancement at the position of Venus
(cf.Fig. 2b). We determine a
upper limit of
to any
flux from Venus in the 1.5-2.0 keV energy range. The corresponding
value for E=2.0-8.0 keV is
.
Further spectral analysis of the ACIS-I data is complicated by the
presence of optical loading
.
Anticipating this possibility, and in order to fully utilize the
spectroscopic capabilities of Chandra, we performed also spectroscopic
observations with the LETG. Figure 3a shows what we
expected to see in the dispersed first order spectra: images of the
Venus crescent in the C-K,
N-K
and O-K
emission lines (the energies were taken from Bearden 1967).
Figure 3b shows the observed spectrum. The
photons from both LETG observations were transformed into the rest
frame of Venus and superimposed. Bright pixels and bright columns were
excluded, and the intrinsic energy resolution of ACIS-S was used to
suppress the background
.
This spectrum, which is completely uncontaminated by optical light, clearly
shows that most of the flux comes from O-K
fluorescence. As this
flux is monochromatic, images of the Venus crescent (illuminated from bottom)
show up along the dispersion direction. The intensity profile along this
direction is shown in Fig. 3c.
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Figure 3:
a) Expected LETG spectrum of Venus on the ACIS-S array. S1 and
S3 are back-illuminated CCDs with increased sensitivity at low
energies (underlined), while the others are front-illuminated. The
nominal aimpoint, in S3, is marked with a circle. The aimpoint was
shifted by 3.25' into S2, to get more of the fluorescent lines
covered by back-illuminated CCDs. During the two parts of the
observation, Venus was moving perpendicularly to the dispersion
direction. In order to avoid saturation of telemetry, the shaded area
around the zero order image in S2 was not transmitted. Energy and
wavelength scales are given along the dispersion direction. Images of
Venus are drawn at the position of the C, N, and O fluorescence lines, with
the correct size and orientation. The dashed rectangle indicates the
section of the observed spectrum shown below.
b) Observed LETG spectrum of Venus, smoothed with a Gaussian
function with
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grating | line | CCD | type | eff area | photons | photon flux | energy flux | energy |
LETG | C-K
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S1 | BI | 14 cm2 | ![]() |
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LETG | N-K![]() |
S1 | BI | 5 cm2 |
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LETG | O-K![]() |
S2 | FI | 7 cm2 | ![]() |
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LETG | O-K![]() |
S3 | BI | 21 cm2 | ![]() |
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none | (O-K![]() |
I1, 3 | FI | 135 cm2 | ![]() |
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For the determination of the line energies we proceeded as follows: from the
ACIS-I observation, we accumulated all photons with
within a circle of
around Venus along the solar direction, to get
the intensity profile perpendicular to the solar direction
(Fig. 2a), i.e., along the dispersion axis in the LETG observation
(Fig. 3). This profile, which shows a characteristic central dip
due to the effect of limb brightening, was then smoothed with a cubic spline
function and used as a template for the spectral fit. For each emission line,
the position and the normalization of the template were determined as free
parameters by
minimization. The dispersion was expressed in
arcseconds, and the width of the template was reduced by 3% (to take the
change of the apparent size into account; Table 1) and kept fixed.
Errors were determined by increasing the minimum
by 1.0, with the
normalization as free parameter. The results, converted into energies, are
listed in Table 2.
The observed line energies are higher than the values of
Bearden (1967), but in all cases the deviation is less than
.
A better agreement, with deviations of less than
,
is achieved with Aikin (1970), who quoted emission energies of 278.4,
392.9, and 526.0 eV for C, N, and O. Recent determinations of the dominant
emission line of atomic oxygen (for the 1s2s22p5(3P
)
configuration) yielded energies in the range 526.8-528.3 eV (McLaughlin & Kirby 1998, and references therein), which is in excellent
agreement with the
found in the LETG-S3
spectrum of Venus. If the oxygen atom is embedded in a molecule, this energy
is slightly shifted: to
527.3 eV for CO2 (Nordgren et al. 1997),
526.8 eV for O2 (McLaughlin & Kirby 1998), and
525.0-526.0 eV for CO (Skytt et al. 1997).
These shifts are too small to be discriminated with
the currently available spectral resolution, which is mainly limited by the
statistical uncertainty due to the low number of detected photons. The
situation is similar for carbon and nitrogen, where the following values were
recently found:
279.2 eV for CO (Skytt et al. 1997) and
393-394 eV for N2 (Nordgren et al. 1997).
Emission line spectra of molecules composed of C, N, and O atoms are
characterized by
an additional fluorescence line, caused by the
transition following core excitation. In CO, this line is much more pronounced
at the C than at the O fluorescence energy
(e.g. Skytt et al. 1997, Figs. 5 and 6).
The energies for the
transition are
287.4 eV for C in CO (Sodhi & Brion 1984),
290.7 eV for C in CO2 (Hitchcock & Ishi 1987),
401.1 eV for N2, 534.2 eV for O in CO (Sodhi & Brion 1984), and
535.4 eV for O in CO2 (McLaren et al. 1987).
We do not see such an additional line at N and O. The image of Venus at
the energy of carbon, however, appears elongated along the dispersion direction
(Fig. 3b), and the spectrum (Fig. 3c) does
indicate the presence of a secondary emission line at
,
consistent with the energy of the C
transition in
CO2 and CO.
Table 2 summarizes the number of detected
photons in the individual lines together with the derived
fluxes. The specified uncertainties are
errors from photon statistics; they do not include
uncertainties in the effective areas at
.
This could be the reason for the
deviation between the two O-K
fluxes observed with S2 and
S3, because both parts of the dispersed spectrum were recorded
simultaneously. For C-K
the number of recorded photons is
sufficient for detection, but not high enough for a reliable flux
determination. For N-K
a marginal detection is only
possible because the few photons were recorded exactly at the expected
position. The fact that no N-K
photons are detected at the
corresponding mirror site may be related to inhomogeneities in the
ACIS-S low energy response, in particular close to the CCD borders.
The last row in Table 2 contains the values for the
direct imaging observation with ACIS-I. Despite the lower
energy resolution and the problem with optical loading, it is very
likely that most of the flux came from O-K
.
The difference
between this flux and the (non-simultaneously obtained) O-K
fluxes from the LETG observations may be related to the X-ray
variability of Venus (Sect.3.3).
It is interesting to compare the total X-ray flux from Venus with
the optical flux. The visual magnitude
corresponds to an optical flux of
.
Adopting a total X-ray flux of
,
we get
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Figure 4:
Temporal behaviour of the soft X-ray flux from the Sun and
Venus. a) 1-8 Å (1.55-12.4 keV) solar flux in
10-3 erg cm-2 s-1 at 1.0 AU, as measured with GOES-8
and GOES-10. b) 1-500 Å (0.025-12.4 keV)
solar flux in 1010 photons cm-2 s-1 scaled to 1.0 AU,
as measured with SOHO/SEM. The times in a) and b)
were shifted by +240 s and +230 s for Jan. 10 and 13, to take the
light travel time delay between Sun![]() ![]() ![]() ![]() ![]() |
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While there is practically no variation of the optical flux from Venus
on time scales of hours and less, the X-ray flux shows indications
for pronounced variability on time scales of minutes
(Fig. 4c). A Kolmogorov-Smirnov test yields probabilities
of only 1% for both the observation with LETG/ACIS-S and ACIS-I
that the count rates are just statistical fluctuations around a
constant value. As variability of the 1-10 keV solar flux by a
factor of two on time scales of minutes is not uncommon
(e.g. Fig. 4a), we expect the scattered solar X-rays from
Venus to exhibit a similar variability. However, a direct comparison
with the solar flux monitored simultaneously with GOES-8 and GOES-10
(Fig. 4a) and SOHO/SEM (Fig. 4b) does not show
an obvious correlation. This may be related to the fact that solar
X-rays are predominantly emitted from localized regions and that
Venus saw a solar hemisphere which was rotated by
(LETG/ACIS-S) and
(ACIS-I) from the solar hemisphere facing
Earth.
Differences from the broad band solar X-ray flux (as measured with
GOES/SOHO) may also arise due to the fact that the X-ray flux from Venus
responds very sensitively to variability of the solar flux in a narrow
spectral range just above the K edges. As will be shown in the next section,
this is particularly the case for O-K
:
while
the C-K
emission increases by only 7% if the coronal temperature
rises by 8%, the O-K
emission increases by 33%.
Estimates on the X-ray luminosity of Venus due to scattering and fluorescence of solar x-rays have recently been made by Cravens & Maurellis (2001). We are, however, not aware of detailed predictions of how Venus would appear in X-rays. For comparison with the observed image we performed a numerical simulation of fluorescent scattering of solar X-rays in the atmosphere of Venus. We did not consider elastic scattering, as the corresponding luminosity is one order of magnitude below the fluorescence luminosity, according to Cravens & Maurellis (2001), and in agreement with the observed LETG spectrum (Fig. 3c).
The ingredients to the model are the composition and density structure
of the Venus atmosphere, the photoabsorption cross sections and
fluorescence efficiencies of the major atmospheric constituents, and
the incident solar spectrum.
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Figure 5:
a) Photoabsorption cross sections
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We adopted the Venus model atmosphere from Seiff (1983),
where the density in the lower and middle atmosphere, i.e., between
the surface and a height of 100 km, is tabulated in steps of 2 km for
different latitudes, while for the upper atmosphere, between 100 km
and 180 km, it is tabulated in steps of 4 km for two solar zenith
angles sza (subsolar:
and antisolar:
). For
we interpolated the density
exponentially. In order to calculate the number density of C, N, and O atoms,
we used the following values for the composition of the atmosphere:
65.2% oxygen, 32.6% carbon and 2.2% nitrogen. As the main
constituents, C and O, are contained in CO2, we assumed this
composition to be homogeneous throughout the atmosphere.
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Figure 6:
Optical depth
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The values for the photoabsorption cross sections were taken from Reilman & Manson (1979). We supplemented them with data from Chantler (1995) at energies close to the K edges. From these values and the C, N, and O contributions listed above, we computed the effective photoabsorption cross section of the Venus atmosphere (Fig. 5a). This, together with the atmospheric density structure, yielded the optical depth of the Venus atmosphere, as seen from outside (Fig. 6).
There is quite some discrepancy in the literature about the K-shell binding energies in C, N, and O atoms. These energies are affected by the outer electrons and thus depend on whether the element is in an atomic, molecular, or solid state. The values 283.8, 401.6, 532.0 eV for C, N, O, which Chantler (1995) computed for isolated atoms, are in good agreement with the values 283.84, 400, and 531.7 eV found by Henke et al. (1982) and used by, e.g., Morrison & McCammon (1983). However, Gould & Jung (1991) found significantly higher K-threshold energies for isolated C, N, and O atoms: 297.37 eV, 412.36 eV, and 546.02 eV. According to Snowden & Freyberg (1993), the values 400 eV and 532 eV of Henke et al. (1982) refer to molecular nitrogen and oxygen, while Ma et al. (1991) quote an ionization potential of 409.938 eV for N2 (and 296.080 eV for C in CO). A compilation by Sevier (1979) lists calculated values of 296.94 eV for atomic carbon, 410.7 eV and 411.88 eV for atomic nitrogen, 411.2 eV for N2, 545.37 eV for atomic oxygen, and 542.2 eV for O2. An accurate treatment of the K-edge is further complicated by the presence of considerable fine-structure: detailed calculations of the inner-shell photoabsorption of oxygen by McLaughlin & Kirby (1998) show that already the atomic state contains an impressive amount of resonance structure around the K-edge.
In order to estimate the consequences of all these uncertainties, we ran our
simulation also with the following, alternative set of K-edge energies:
In both cases we assumed that the energy will be emitted at
279.2, 393.5 and 527.3 eV, according to recent determinations and in
agreement with the observed LETG spectrum (Sect.3.2).
We found no significant differences in the results (Table 3).
The solar spectra for 2001 January 10 and 13 were derived from
SOLAR2000
(Tobiska et al. 2000).
These spectra, best estimate daily average values, do not
differ much between both dates. To improve the coverage towards
energies
,
we computed synthetic spectra with
the model of Mewe et al. (1985) and aligned them with the
SOLAR2000 spectra in the range 50-500 eV, by adjusting the
temperature and intensity. We derived a coronal temperature of
.
The spectrum, scaled to the heliocentric
distance of Venus, is shown in Fig. 5b (upper curve),
with a bin size of 1 eV, which we used in order to preserve the
spectral details.
The high dynamic range in the optical depth of the Venus atmosphere
(Fig. 6) requires a model with high spatial resolution. For
the simulation we choose a right-handed coordinate system (x,y,z)with the center of Venus at (0, 0, 0) and the Sun at
.
We
sample the irradiated part of the Venus atmosphere with a grid of
volume elements of size
.
Figure 6 shows that the atmosphere becomes optically thick for
X-rays with
already at heights above 100 km. This
means that solar X-rays do not reach the atmosphere below 100 km,
where the density shows some latitudinal dependence. Above 100 km,
however, the density of the model atmosphere depends only on the height
and the solar zenith angle. This simplifies the calculation: instead
of computing the solar irradiation of the volume elements
V(xi,yj,zk), it is only necessary to do this for volume elements
V(rij,zk), with
rij=(xi2+yj2)1/2.
The whole information about the irradiation of the atmosphere can thus
be computed and stored in volume elements
Vik=V(ri,zki),
i=1...n and
ki=1...mi, with
With this grid the calculation is performed in two steps: in the first step the solar radiation absorbed in each volume element is calculated and stored. In the second step an image of Venus is accumulated for a particular phase angle by integrating the emission and subsequent absorption of the corresponding volume elements along the line of sight.
The first step is performed by propagating the irradiation for each
column i from the top of the atmosphere along the -z direction,
i.e., away from the Sun. For the center of the corresponding volume
element Vik, the mass densitity is calculated from the height
above the surface and the solar zenith angle, by exponential
interpolation of the nearest tabulated grid points in the Venus model
atmosphere, and converted into a number density nik of the sum of
C, N, and O atoms. From the column densities
As the fluorescence photons are emitted at an energy
which
is below the corresponding K edge
(see Sects. 3.2 and 4.2), they are
not subject to subsequent K shell absorption by the same element. K shell absorption by lighter elements, however, is possible. All
photons are subject to elastic scattering, but as this process is
nearly isotropic, only weakly energy dependent and characterized
by cross sections which are several orders of magnitude smaller than
for photoabsorption, it should not much affect the distribution of
photons along the line of sight. We treat the attenuation of the
reemitted photon flux due to subsequent photoabsorption along the line
of sight in a similar way as we did for the attenuation of the
incident solar flux, but this time only at the three discrete energies
.
By sampling the radiation in the volume elements along the
line of sight, starting from the volume element which is farthest away
from the observer, we can then accumulate images of Venus in the three
energies
in, e.g., orthographic projection, for any phase
angle.
![]() Photon flux in units of 10-4 ph cm-2 s-1. Energy flux in units of 10-14 erg cm-2 s-1. |
The simulated images of Venus at the K
fluorescence
energies C, N, and O are shown in Figs. 8a-c.
They agree well with the observed X-ray image (Fig. 8d),
while the optical image (Fig. 8e) is characterized by
a different brightness distribution. In X-rays, Venus exhibits
significant brightening at the sunward limb, accompanied by reduced
brightness at the terminator, which causes it to appear less than half
illuminated. This is a consequence of the fact that the volume
emissivity extends into the tenuous, optically thin parts of the
thermosphere and exosphere (Fig. 7). From there, the volume
emissivities are accumulated along the line of sight without
considerable absorption, so that the observed brightness is mainly
determined by the extent of the atmospheric column along the line of
sight. Detailed comparison of the images shows that the amount of limb
brightening is different for the three energies. This can be
understood in the following way.
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Figure 7:
Volume emissivities of C, N, and O K![]() ![]() |
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Figure 8:
a-c) Simulated X-ray images of Venus at
C-K![]() ![]() ![]() ![]() ![]() ![]() ![]() ![]() ![]() ![]() ![]() |
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If the incident solar spectrum consisted only of photons above the
O-K
edge, then the peak of volume emissivity would occur
at the same height for all fluorescent lines, and this height would be
determined by the spectral hardness of the incident solar flux.
Differences in the height of the volume emissivity peak between C, N,
and O occur due to the presence of photons with energies between the
individual K
edges. Photons with energies between
N-K
and O-K
,
for example, influence the
atmospheric height of maximum N-K
emission, but do not
affect the O-K
peak. Due to the presence of such photons
in the incident solar spectrum (Fig. 5b), which are
affected by less photoabsorption (Fig. 5a) and
penetrate deeper into the atmosphere, the nitrogen volume emissivity
peak occurs at the lowest atmospheric heights (Fig. 7). In
a similar way, the carbon emissivity peak lies just below that of oxygen.
Although the difference in the atmospheric heights of the individual
peaks is only a few kilometers, this has consequences for the
appearance of Venus in the individual fluorescence lines. At heights
of
,
the density doubles every 3 km with
decreasing height. The deeper in the atmosphere the emission occurs,
the more absorbing layers are above. This effect is particularly
important at the limb, where the column of absorbing material along
the line of sight reaches a maximum, thus reducing the amount of limb
brightening. Another factor which determines the amount of limb
brightening is the photoabsorption cross section at the fluorescence
energy. This energy is just below the corresponding K-edge
(see Sects. 3.2 and 4.2).
Figure 5a shows that for the chemical composition of the
Venus atmosphere the photoabsorption cross section for nitrogen
K
fluorescence photons is
about twice as large as that for carbon and oxygen
K
fluorescence photons, causing an additional attenuation
of the limb brightness in the nitrogen image.
The simulations show that the limb brightening depends sensitively on
the density and chemical composition of the Venus atmosphere. Thus,
precise measurements of this brightening will provide direct
information about the atmospheric stucture in the thermosphere and
exosphere. With ACIS-I a brightening of
was observed
(Fig. 2b). From the computed images, smoothed with a
Gaussian function with
,
we determine corresponding limb
brightenings of 2.0 for C-K
,
1.7 for N-K
,
and 2.2 for O-K
.
The simulated images can also be used to derive the flux from the
whole visible side of Venus in the three energies. Table 3
shows that the derived flux values highly depend on the coronal
temperature, in particular for O-K
.
We derive from the
simulation a total flux of
for all three lines.
The corresponding value obtained from the LETG/ACIS-S observation
(Table 2) is
or
,
depending on whether we take the O-K
flux from the S2 or
S3 CCD. Considering all the uncertainties, these values are in good agreement
with each other.
In order to study the angular distribution of the scattered photons, we
computed X-ray images of Venus for different phase angles (Fig. 9a). We scanned the full range of
from
to
with a step size of
,
and determined the corresponding
X-ray intensities for the three emission lines by integrating the observed
flux from the images. Figure 9b shows the result. It is evident
that the intensity declines first very slowly, staying above half of its
maximum value for
.
At
,
the intensity has
dropped to
.
This illustrates that the solar X-rays are preferentially
scattered back towards the Sun. For larger phase angles the decline becomes
faster. Between
and
,
the intensity drops by a factor of
two, and the X-ray crescent starts to evolve into a thin ring
around the dark planet, which is seen fully developed at
.
This ring might be observable with sufficiently sensitive solar X-ray
observatories immediately before and after the upcoming Venus transits on 8
June 2004 and 5/6 June 2012. However, such observations would be extremely
challenging, as the intensity of the ring will be only 0.3% of the fully
illuminated Venus.
By spherically integrating the X-ray intensity for the three energies
(Fig. 9a) over phase angle, we determined the luminosities
listed in Table 3. The total X-ray luminosity of Venus,
55
14 MW, agrees well with the prediction of Cravens & Maurellis (2001), who estimated a luminosity of 35 MW with an uncertainty factor
of about two.
![]() |
Figure 9:
X-ray intensity of Venus as a function of phase angle, in the
fluorescence lines of C, N, and O. The images at top, all displayed in the
same intensity coding, illustrate the appearence of Venus at O-K![]() |
Open with DEXTER |
The Chandra data are fully consistent with fluorescent scattering of
solar X-rays in the Venus atmosphere. This is an especially
interesting result when compared with the X-ray emission of comets,
where the dominant process for the X-ray emission is charge exchange
between highly charged heavy ions in the solar wind and cometary
neutrals. The LETG/ACIS-S spectrum (Fig. 3)
definitively rules out that a similar process dominates the X-ray flux
from the atmosphere of Venus at heights below
.
The LETG/ACIS-S spectrum, however, does not exclude charge
exchange interactions in the outer exosphere of Venus, as they would be
too faint to be detected in the dispersed spectrum. A more sensitive
method for finding charge exchange signatures there is to look for
enhancements of the surface brightness in the environment of Venus. In
fact, the ACIS-I data do show indications for a decrease of the
surface brightness with increasing distance from Venus from
to
,
in the energy range 0.2-1.5 keV (Fig. 2). We
find the brightness at
to exceed that at
by
on the dayside and by
on the nightside. Both values are consistent with each other
and yield a mean excess of
.
Observations of comets in the ROSAT all-sky survey 1990-1991,
also performed at solar maximum, show that the peak surface X-ray
brightness which can be reached by charge exchange is
at
for an average composition and
density of the low-latitude solar wind (Dennerl et al. 1997);
it scales with r-2. This result is in good agreement with the
theoretical estimate by Cravens (1997). Charge exchange
produces a spectrum consisting of many narrow emission lines. The
overall properties, however, can be approximated by 0.2 keV thermal
bremsstrahlung emission (Wegmann et al. 1998). By applying
this approximation to the ACIS-I observation, we obtain a maximum
countrate due to charge exchange of
at
.
This maximum value is only slightly (by
)
larger than the
excess observed at
,
which implies that the exosphere
should be almost collisionally thick 4300 km above the surface.
At this height, however, both the hydrogen and the hot oxygen densities
are
(Bertaux et al. 1982; Nagy & Cravens 1988)
and thus orders of magnitude too low. Furthermore, the ACIS-I
spectrum of all events within
and
radius around Venus
(not affected by optical loading) shows no evidence for the spectral
signatures observed in the ACIS-S spectrum of Comet C/1999 S4
(LINEAR), which were attributed to charge exchange interactions
(Lisse et al. 2001). We conclude that the observed excess in
the surface brightness is either spurious or produced by other effects.
At heights of 155-180 km, however, the exosphere of Venus does
become collisionally thick due to the large cross section of charge
transfer interactions (Fig. 6). This implies that if the
flux of highly charged heavy solar wind ions reached these atmospheric
layers, we would indeed observe the maximum flux estimated above. But
even then not more than about 3 photons would have been detected due to
charge exchange from the area of the crescent during the ACIS-I
exposure. Taking the presence of an ionosphere into account, which
shields the lower parts from the solar wind, then even this estimate
appears to be too high. By integrating the X-ray production rate over
altitude, starting at 500 km, the approximate location of the
ionopause, where the density is dominated by the hot oxygen corona,
Cravens (2000) estimated a luminosity of
for the total X-ray luminosity of Venus due to charge exchange. With
the 0.2 keV thermal bremsstrahlung approximation this would result in
a total ACIS-I countrate of
,
or only 0.2 counts accumulated during the whole observation. We
conclude that the observation was not sensitive enough for detecting
charge exchange signatures.
It is interesting to compare the X-ray properties of Venus with those of comets, where the X-ray emission is dominated by the charge exchange process, while fluorescent scattering of solar X-rays is negligible. This opposite behaviour is a direct consequence of the different cross sections of both processes and the way the gas is distributed.
The cross sections for charge exchange typically exceed
and are thus at least three orders of magnitude
greater than for fluorescent emission, which are
and less in the energy range of interest (Fig. 5a). The
gas in a cometary coma is distributed over a much larger volume and
solid angle than in a planetary atmosphere. The particle density in a
coma is too low to reach a column density for efficiently scattering
solar X-rays, but high enough to provide a sufficient number of target
electrons for charge exchange.
The atmosphere of Venus, on the other hand, is so dense that it is
optically thick even to fluorescent scattering (Fig. 6). As
the solar wind ions become discharged already in the outermost parts,
only a tiny fraction of the atmospheric electrons can participate in
the charge exchange process. Additionally, the flux of incident solar
wind ions is reduced by the presence of an ionosphere. Even in the
absence of an ionosphere, the peak X-ray surface flux due to charge
exchange would not exceed that of a comet (with a sufficiently dense
coma), when exposed to the same solar wind conditions. But while the
X-ray bright area is confined to less than one arcminute in diameter in
the case of Venus, the bright part of the cometary X-ray emission can
extend over tens of arcminutes, thus increasing the total amount of
charge exchange induced X-ray photons by two orders of magnitude or more.
The Chandra observation clearly shows that Venus is an X-ray source. From the
X-ray spectrum and morphology we conclude that fluorescent scattering of
solar X-rays is the main process for this radiation, which is
dominated by the K
emission lines from C, N, and O, plus some
possible contribution from the C
transition in CO2 and
CO. By modeling the X-ray appearance of Venus due to fluorescence, we have
demonstrated that the amount of limb brightening depends sensitively on the
properties of the Venus atmosphere at heights above 110 km. Thus, information
about the chemical composition and density structure of the Venus thermosphere
and exosphere can be obtained by measuring the X-ray brightness distribution
across the planet at the individual K
fluorescence lines. This opens
the possibility of using X-ray observations for remotely monitoring the
properties of regions in the Venus atmosphere which are difficult to
investigate otherwise, and their response to solar activity.
Acknowledgements
SOLAR2000 Research Grade v1.15 historical irradiances are provided courtesy of W. Kent Tobiska and SpaceWx.com. These historical irradiances have been developed with funding from the NASA UARS, TIMED, and SOHO missions. The SOHO CELIAS/SEM data were provided by the USC Space Sciences Center. SOHO is a joint European Space Agency, United States National Aeronautics and Space Administration mission.