Astronomy & Astrophysics manuscript no. ms c© ESO 2021
August 26, 2021
Spectral Analysis of the Chandra Comet Survey
D. Bodewits1, D. J. Christian2, M. Torney3, M. Dryer4, C. M. Lisse5, K. Dennerl6, T. H. Zurbuchen7, S. J. Wolk8,
A. G. G. M. Tielens9, and R. Hoekstra1
1 kvi atomic physics, University of Groningen, Zernikelaan 25, NL-9747 AA Groningen, The Netherlands
e-mail: bodewits@kvi.nl, hoekstra@kvi.nl
2 Queen’s University Belfast, Department of Physics and Astronomy, Belfast, BT7 1NN, UK
e-mail: d.christian@qub.ac.uk
3 Atoms Beams and Plasma Group, University of Strathclyde, Glasgow, G4 0NG, UK
e-mail: torney@phys.strath.ac.uk
4 noaa Space Environment Center, 325 Broadway, Boulder, CO 80305, USA
e-mail: murray.dryer@noaa.gov
5 Planetary Exploration Group, Space Department, Johns Hopkins University Applied Physics Laboratory, 11100 Johns Hopkins Rd,
Laurel, MD 20723, USA
e-mail: carey.lisse@jhuapl.edu
6 Max-Planck-Institut für extraterrestrische Physik, Giessenbachstrasse, 85748 Garching, Germany
e-mail: kod@mpe.mpg.de
7 The University of Michigan, Department of Atmospheric, Oceanic and Space Sciences, Space Research Building, Ann Arbor, MI
48109-2143, USA
e-mail: thomasz@umich.edu
8 Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA
e-mail: swolk@head.cfa.harvard.edu
9 nasa Ames Research Center, MS 245-3, Moffett Field, CA 9435-1000, USA
e-mail: tielens@astro.rug.nl
Received August 26, 2021
ABSTRACT
Aims. We present results of the analysis of cometary X-ray spectra with an extended version of our charge exchange emission
model (Bodewits et al. 2006). We have applied this model to the sample of 8 comets thus far observed with the Chandra X-ray
observatory and acis spectrometer in the 300–1000 eV range. The surveyed comets are C/1999 S4 (linear), C/1999 T1 (McNaught–
Hartley), C/2000 WM1 (linear), 153P/2002 (Ikeya–Zhang), 2P/2003 (Encke), C/2001 Q4 (neat), 9P/2005 (Tempel 1) and 73P/2006-
B (Schwassmann–Wachmann 3) and the observations include a broad variety of comets, solar wind environments and observational
conditions.
Methods. The interaction model is based on state selective, velocity dependent charge exchange cross sections and is used to explore
how cometary X-ray emission depend on cometary, observational and solar wind characteristics. It is further demonstrated that
cometary X-ray spectra mainly reflect the state of the local solar wind. The current sample of Chandra observations was fit using
the constrains of the charge exchange model, and relative solar wind abundances were derived from the X-ray spectra.
Results. Our analysis showed that spectral differences can be ascribed to different solar wind states, as such identifying comets
interacting with (I) fast, cold wind, (II), slow, warm wind and (III) disturbed, fast, hot winds associated with interplanetary coronal
mass ejections. We furthermore predict the existence of a fourth spectral class, associated with the cool, fast high latitude wind.
Key words. Surveys, atomic processes, molecular processes, Sun: solar wind, coronal mass ejections (cmes), X-rays: solar system,
Comets: general Comets: individual: C/1999 S4 (linear), C/1999 T1 (McNaught–Hartley), C/2000 WM1, 153P/2002 (Ikeya–Zhang),
2P/2003 (Encke), C/2001 Q4 (neat), 9P/2005 (Tempel 1) and 73/P-B 2006 (Schwassmann–Wachmann 3B)
1. Introduction
When highly charged ions from the solar wind collide on a
neutral gas, the ions get partially neutralized by capturing elec-
trons into an excited state. These ions subsequently decay to the
ground state by the emission of one or more photons. This pho-
ton emission is called charge exchange emission (cxe) and it has
been observed from comets, planets and the interstellar medium
in X-rays and the Far-UV Lisse et al. (1996); Krasnopolsky
(1997); Snowden et al. (2004); Dennerl (2002). The spec-
tral shape of the cxe depends on properties of both the neutral
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gas and the solar wind and the subsequent emission can there-
fore be regarded as a fingerprint of the underlying interactions
Cravens et al. (1997); Kharchenko and Dalgarno (2000, 2001);
Beiersdorfer et al. (2003); Bodewits et al. (2004a, 2006).
Since the first observations of cometary X-ray emission,
more than 20 comets have been observed with various X-ray
and Far-UV observatories Lisse et al. (2004); Krasnopolsky et
al. (2004). This observational sample contains a broad variety of
comets, solar wind environments and observational conditions.
The observations clearly demonstrate the diagnostics available
from cometary charge exchange emission.
First of all, the emission morphology is a tomography of
the distribution of neutral gas around the nucleus Wegmann et
2 D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey
al. (2004). Gaseous structures in the collisionally thin parts
of the coma brighten, such as the jets in 2P/Encke Lisse et al.
(2005), the Deep Impact triggered plume in 9P/Tempel 1 Lisse
et al. (2007) and the unusual morphology of comet 6P/d’Arrest
Mumma et al (1997). In other comets, the X-ray emission
clearly mapped a spherical gas distribution. This resulted in a
characteristic crescent shape for larger and hence collisionally
thick comets observed at phase angles of roughly 90 degrees
(e.g. Hyakutake - Lisse et al. (1996), linear S4 - Lisse et al.
(2001)). Macroscopic features of the plasma interaction such as
the bowshock are observable, too Wegmann & Dennerl (2005).
Secondly, by observing the temporal behavior of the
comets X-ray emission, the activity of the solar wind and
comet can be monitored. This was first shown for comet
C/1996 B2 (Hyakutake) Neugebauer et al. (2000) and re-
cently in great detail by long term observations of comet
9P/2005 (Tempel 1) Willingale et al. (2006); Lisse et al.
(2007) and 73P/2006 (Schwassmann–Wachmann 3C) Brown et
al. (2007), where cometary X-ray flares could be assigned to
either cometary outbursts and/or solar wind enhancements.
Thirdly, cometary spectra reflect the physical characteristics
of the solar wind; e.g. spectra resulting from either fast, cold
(polar) wind and slow, warm equatorial solar wind should be
clearly different Schwadron and Cravens (2000); Kharchenko
and Dalgarno (2001); Bodewits et al. (2004a). Several attempts
were made to extract ionic abundances from the X-ray spectra.
The first generation spectral models have all made strong
assumptions when modelling the X-ray spectra Haeberli et al
(1997); Wegmann et al. (1998); Kharchenko and Dalgarno
(2000); Schwadron and Cravens (2000); Lisse et al. (2001);
Kharchenko and Dalgarno (2001); Krasnopolsky et al. (2002);
Beiersdorfer et al. (2003); Wegmann et al. (2004); Bodewits et
al. (2004a); Krasnopolsky (2004); Lisse et al. (2005). Here,
we present a more elaborate and sophisticated procedure to an-
alyze cometary X-ray spectra based on atomic physics input,
which for the first time allows for a comparative study of all
existing cometary X-ray spectra. In Section 2, our comet-wind
interaction model is briefly introduced. In Section 3, it is demon-
strated how cometary spectra are affected by the velocity and
target dependencies of charge exchange reactions. In Section 4,
the various existing observations performed with the Chandra
X-ray Observatory, as well as the solar wind data available are
introduced. Based upon our modelling, we construct an analyt-
ical method of which the details and results are presented in
Section 5. In Section 6, we discuss our results in terms of comet
and solar wind characteristics. Lastly, in Section 7 we summa-
rize our findings. Details of the individual Chandra comet ob-
servations are given in Appendix A.
2. Charge Exchange Model
2.1. Atomic structure of He-like ions
Electron capture by highly charged ions populates highly excited
states, which subsequently decay to the ground state. These cas-
cading pathways follow ionic branching ratio statistics. Because
decay schemes work as a funnel, the lowest transitions (n = 2→
1) are the strongest emission lines in cxe spectra. For helium-like
ions, these are the forbidden line (z: 1s2 1S0–1s2s 3S1), the inter-
combination lines (y, x: 1s2 1S0–1s2p 3P1,2), and the resonance
line (w: 1s2 1S0–1s2p 1P1), see Figure 1.
The apparent branching ratio, Beff , for the intercombination
transitions is determined by weighting branching ratios (B j) de-
rived from theoretical transition rates compiled by Porquet et al.
Fig. 1. Part of the decay scheme of a helium–like ion. The 1S0 decays to
the ground state via two-photon processes (not indicated).
Table 1. Apparent effective branching ratios (Beff) for the relaxation of
the 23P-state of He-like carbon, nitrogen, oxygen and neon.
transition C v N vi O vii Ne ix
1s2 (1S0)–1s2p (3P1,2) 0.11 0.22 0.30 0.34
1s2s (3S1)–1s2p (3P0,1,2) 0.89 0.78 0.70 0.66
(2000, 2001), by an assumed statistical population of the triplet
P-term:
Beff =
(2 j + 1)
(2L + 1)(2S + 1)
· B j (1)
The resulting effective branching ratios are given in Table 1.
These ratios can only be observed at conditions where the
metastable state is not destroyed (e.g. by UV flux or collisions)
before it decays. In contrast to many other astrophysical X-
ray sources, this condition is fulfilled in cometary atmospheres,
making the forbidden lines strong markers of cxe emission.
2.2. Emission Cross Sections
To obtain line emission cross sections we start with an initial
state population based on state selective electron capture cross
sections and then track the relaxation pathways defined by the
ion’s branching ratios.
Electron capture reactions can be strongly dependent on
target effects. An important difference between reactions with
atomic hydrogen and the other species is the presence of multi-
ple electrons, hence allowing for multiple (mostly double) elec-
tron transfer. It has been demonstrated both experimentally and
theoretically that double electron capture can be an important
reaction channel in multi-electron targets and that after autoion-
ization to an excited state it may contribute to the X-ray emis-
sion Ali et al. (2005); Hoekstra et al. (1989); Beiersdorfer
et al. (2003); Otranto et al (2006); Bodewits et al. (2006).
Unfortunately, experimental data on reactions with species typ-
ical for cometary atmospheres, such as H2O, atomic O and CO
are at best scarcely available. Because the first ionization po-
tentials of these species are all close to that of atomic H, using
state selective one electron capture cross sections for bare ions
charge exchanging with atomic hydrogen from theory is a rea-
sonable assumption, which is also confirmed by experimental
studies Greenwood et al. (2000, 2001); Bodewits et al. (2006).
Here, we will use the working hypothesis that effective one elec-
tron cross sections for multi-electron targets present in cometary
atmospheres are at least roughly comparable to cross sections for
D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey 3
Table 2. Compilation of theoretical, velocity dependent emission cross sections for collisions between bare- and H-like solar wind ions and atomic
hydrogen, in units of 10−16 cm2. See text for details. We estimate uncertainties to be ca. 20%. The ion column contains the resulting ion, not the
original solar wind ion. Line energies compiled from Garcia & Mack (1965); Vainshtein & Safronova (1985); Drake (1988); Savukov et al.
(2003) and the chianti database Dere et al. (1997); Landi et al. (2006).
E (eV) Ion Transition 200 km s−1 400 km s−1 600 km s−1 800 km s−1 1000 km s−1
299.0 C v z 8.7 12 16 18 20
304.4 C v x,y 0.65 1.0 1.5 1.7 1.8
307.9 C v w 1.8 3.0 4.1 4.8 5.2
354.5 C v 1s3p-1s2 0.55 0.71 0.81 1.0 1.3
367.5 C v 1s4p-1s2 0.70 0.66 0.76 0.74 0.72
367.5 C vi 2p-1s 15 26 30 33 34
378.9 C v 1s5p-1s2 0.00 0.02 0.05 0.04 0.04
419.8 N vi z 13 23 28 29 29
426.3 N vi x,y 2.7 4.3 5.3 5.7 6.0
430.7 N vi w 3.8 6.0 7.4 8.1 8.5
435.5 C vi 3p-1s 1.6 4.0 4.7 4.7 4.8
459.4 C vi 4p-1s 2.9 5.9 7.0 6.4 6.0
471.4 C vi 5p-1s 0.55 1.0 1.3 0.85 0.54
497.9 N vi 1s3p-1s2 0.43 0.99 1.3 1.3 1.3
500.3 N vii 2p-1s 40 45 44 42 42
523.0 N vi 1s4p-1s2 0.81 1.6 1.9 1.8 1.7
534.1 N vi 1s5p-1s2 0.14 0.31 0.33 0.21 0.14
561.1 O vii z 37 34 33 32 31
568.6 O vii x,y 10 10 10 9.9 9.7
574.0 O vii w 9.9 11 11 11 10
592.9 N vii 3p-1s 6.3 4.9 4.8 4.5 4.3
625.3 N vii 4p-1s 2.9 2.9 3.7 4.3 4.6
640.4 N vii 5p-1s 11 5.2 3.7 2.7 2.2
650.2 N vii 6p-1s 0.00 0.21 0.13 0.09 0.08
653.5 O viii 2p-1s 27 40 48 51 53
665.6 O vii 1s3p-1s2 1.7 1.3 1.3 1.2 1.2
697.8 O vii 1s4p-1s2 0.81 0.79 1.0 1.2 1.3
712.8 O vii 1s5p-1s2 2.8 1.3 0.92 0.68 0.54
722.7 O vii 1s6p-1s2 0.00 0.06 0.04 0.02 0.02
774.6 O viii 3p-1s 2.6 4.7 5.6 5.3 5.0
817.0 O viii 4p-1s 1.0 1.6 2.0 2.2 2.3
836.5 O viii 5p-1s 2.4 4.0 4.6 4.1 3.7
849.1 O viii 6p-1s 1.6 1.6 1.5 1.1 0.67
one electron capture from H. Based on this hypothesis, we will
use our comet-wind interaction model to evaluate the contribu-
tion of the different species.
For our calculations, we use a compilation of theoretical state
selective, velocity dependent cross sections for collisions with
atomic hydrogen Errea et al. (2004); Fritsch and Lin (1984);
Green et al. (1982); Shipsey et al. (1983). We furthermore
assume that capture by H-like ions leads to a statistical triplet
to singlet ratio of 3:1, based on measurements by Suraud et al.
(1991); Bliek et al. (1998). We will first focus on the strongest
emission features, which are the n = 2 → 1 transitions, i.e.,
the Ly-α transition (H-like ions) or the forbidden, resonance and
intercombination lines (He-like ions).
In Fig. 2, the emission cross sections of the Ly-α or the sum
of the emission cross sections of the forbidden, resonance and
intercombination lines of different ions (C, N, O) are shown as
a function of collision velocity, for one electron capture reac-
tions with atomic hydrogen. This figure sets the stage for solar
wind velocity induced effects in cometary X-ray spectra. Most
important is the effect of the velocity on the two carbon emis-
sion features; their prime emission features increase by a factor
of almost two when going from typical ‘slow’ to typical ‘fast’
solar wind velocities. The O viii Ly-α emission cross section can
be seen to drop steeply below ca. 300 km s−1. The N vi K-α dis-
plays a similar, though somewhat less strong behavior.
Fig. 2. Velocity dependence of Ly-α or the sum of the forbid-
den/resonance/intercombination emission cross sections of different so-
lar wind ions: O viii (dashed, grey line), O vii (solid, black line), N vii
(dotted, black line), N vi (solid, grey line), C vi (dashed, black line) and
C v (dash-dotted, black line).
The relative intensity of the emission lines (per species) is
governed by the state selective electron capture cross sections
4 D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey
Fig. 3. Velocity dependence of the hardness ratio of different solar wind
ions: O viii (solid line), O vii (dashed line) N vii (dashed line) and C vi
(dash-dotted line). Also shown are two experimentally obtained hard-
ness ratios by Beiersdorfer et al. (2001) and Greenwood et al. (2000)
for O8+ colliding on CO2 and H2O, respectively (see text).
of the charge exchange reaction and the branching ratios of the
resulting ion. A measure of these intensities is the hardness ra-
tio (Beiersdorfer et al. 2001), which is defined as the ratio be-
tween the emission cross sections of the higher order terms of
the Lyman-series and Ly-α (or between the higher order K-series
and K-α in case of He-like ions):∑
n>2 σem(Ly−n)
σem(Ly−α)
For electron capture by H-like ions, we will use the ratio be-
tween the sum of the resonance-, intercombination and forbid-
den emission lines and the rest of the K-series as the hardness
ratio. Fig. 3 shows the hardness ratios of cxe from abundant so-
lar wind ions. The figure shows that most hardness ratios are
constant at typical solar wind velocities (above 300 km s−1) but
it also clearly demonstrates the suggestion made by Beiersdorfer
et al. (2001) that hardness ratios are good candidates for studies
of velocimetry deep within the coma when the solar wind has
slowed down by mass loading.
2.3. Interaction Model
Cometary high-energy emission depends upon certain properties
of both the comet (gas production rate, composition, distance to
the Sun) and the solar wind (speed, composition). Recently, we
developed a model that takes each of these effects into account
Bodewits et al. (2006), which we will briefly describe here.
The neutral gas model is based on the Haser-equation, which
assumes that a comet has a spherically expanding neutral coma
Haser (1957); Festou (1981). The lifetime of neutrals in the
solar radiation field varies greatly amongst species typical for
cometary atmospheres Huebner et al. (1992). The dissociation
and ionization scale lengths also depend on absolute UV fluxes,
and therefore on the distance to the Sun. The coma interacts
with solar wind ions, penetrating from the sunward side follow-
ing straight line trajectories. The charge exchange processes be-
tween solar wind ions and coma neutrals are explicitly followed
both in the change of the ionization state of the solar wind ions
Fig. 4. Modeled charge state distribution along the comet-Sun line, as-
suming an equatorial 300 km s−1 wind interacting with a comet with
outgassing rate Q=1029 molecules s−1 at 1 AU from the Sun. A compo-
sition typical for the slow, equatorial wind was assumed.
and in the relaxation cascade of the excited ions (as discussed
above).
Due to its interaction with the cometary atmosphere, the so-
lar wind is both decelerated and heated in the bow shock. This
bow shock does not affect the ionic charge state distribution. The
bow shock lowers the drift velocity of the wind but at the same
time increases its temperature and the net collision velocity of
the ions is ca. 77% of the initial velocity v(∞) throughout the
interaction zone. We use a rule of thumb derived by Wegmann
et al. (2004) to estimate the stand-off distance Rbs of the bow
shock.
Deep within the coma, the solar wind finally cools down as
the hot wind ions, neutralized by charge exchange, are replaced
by cooler cometary ions. For simplicity however, we shall as-
sume that the wind keeps a constant velocity and temperature
after crossing the bow shock.
Initially, the charge state distribution depends on the solar
wind state. For most simulation purposes, we will assume the
‘average’ ionic composition for the slow, equatorial solar wind
as given by Schwadron and Cravens (2000). Using our compi-
lation of charge changing cross sections, we can solve the differ-
ential equations that describe the charge state distribution in the
coma in the 2D-geometry fixed by the comet-Sun axis. Figure 4
shows the charge state distribution for a 300 km s−1 equato-
rial wind interacting with a comet with an outgassing rate Q of
= 1029 molecules s−1 comet. From this charge state distribution,
it can be seen that along the comet-Sun axis, the comet becomes
collisionally thick between 3500 km (O8+) to 2000 km (C6+),
depending on the cross section of the ions. A maximum in the
C5+ abundance can be seen around 2,000 km, which is due to the
relatively large initial C6+ population and the small cross section
of C5+ charge exchange.
A 3D integration assuming cylindrical symmetry around the
comet-Sun axis finally yields the absolute intensity of the emis-
sion lines. Effects due to the observational geometry (i.e. field of
view and phase angle) are included at this step in the model.
D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey 5
Fig. 5. Relative contribution of target species to the total intensity of
O vii 570 eV emission complex with increasing field of view, for an
active Q= 1029 molecules s−1 comet, interacting with a 300 km s−1
solar wind at 1 AU from the Sun. The shaded area indicates the range
of apertures used to obtain spectra discussed within this survey.
3. Model Results
3.1. Relative Contribution of Target Species
Figure 5 shows the dominant collisions which underly the X-ray
emission of comets. Shown is the total intensity projected on the
sky, with increasing field of view. Within 104 km around the nu-
cleus, water is the dominant collision partner. Farther outward
(≥ 2 × 105 km), the atomic dissociation products of water take
over, and atomic oxygen becomes the most important collision
partner. When the field of view exceeds 107 km, atomic hydro-
gen becomes the sole collision partner. Note that collisions with
water never account for 100% of the emission, even with very
small apertures, due to the contribution of collisions with atomic
hydrogen, OH and oxygen in the line of sight towards the nu-
cleus.
The comets observed with Chandra are all observed with an
aperture of ca. 7.5′ centered on the nucleus. This corresponds to
a range of 1.6−22×104 km (as indicated in Figure 5). Our model
predicts that the emission from nearby comets will be dominated
by cxe from water, but that for comets observed with a larger
field of view, up to 60% of the emission can come from cxe
interactions with the water dissociation products atomic oxygen
and OH, and 10% from interactions with atomic hydrogen.
3.2. Solar Wind Velocity
To illustrate solar wind velocity induced variations in charge ex-
change spectra, we simulated charge exchange spectra follow-
ing solar wind interactions between an equatorial wind and a
Q = 1029 molecules s−1 comet, and assumed the same solar
wind composition in all cases. In Fig. 6, spectra resulting from
collisional velocities of 300 km s−1 and 700 km s−1 are shown.
In the spectrum from the faster wind, the C vi 367 eV and O vii
570 eV emission features are roughly equally strong, whereas at
300 km s−1, the oxygen feature is clearly stronger. Assuming the
wind’s composition remains the same, within the range of typ-
ical solar wind velocities (300–700 km s−1), the cross sectional
dependence on solar wind velocity does not affect cometary X-
ray spectra by more than a factor 1.5. In practice, the composi-
tional differences between slow and fast wind will induce much
stronger spectral changes.
Fig. 6. Simulated X-ray spectra for a 1029 molecules s−1 comet interact-
ing with an equatorial wind with velocities of 300 km s−1 (solid grey
line) and 700 km s−1 (dashed black line). The spectra are convolved
with Gaussians with a width of σ = 50 eV to simulate the Chandra
spectral resolution. To indicate the different lines, also the 700 km s−1
σ = 1 eV spectrum is indicated (not to scale). A field of view of 105 km
and ‘typical’ slow wind composition were used.
3.3. Collisional Opacity
Many of the 20+ comets that have been observed in X-ray dis-
play a typical crescent shape as the solar wind ion content
is depleted via charge exchange. Comets with low outgassing
rates around 1028 molecules s−1, such as 2P/2002 (Encke) and
9P/2005 (Tempel 1), did not display this emission morphology
Lisse et al. (2005, 2007). Whether or not the crescent shape
can be resolved depends mainly on properties of the comet (out-
gassing rate), but, to a minor extent, also on the solar wind
(velocity dependence of cross sections). Other parameters (sec-
ondary, but important), are the spatial resolution of the instru-
ment and the distance of the comet to the observer.
In a collisionally thin environment, the ratio between emis-
sion features is the product of the ion abundance ratios and the
ratio between the relevant emission cross sections:
rthin =
n(Aq+)
n(Bq+)
em (v)
em (v)
The flux ratio for a collisionally thick system depends on the
charge states considered. In case of a bare ion A and a hydro-
genic ion B, the ratio between the photon fluxes from A and B
is given by the abundance ratio weighted by efficiency factors µ
and η:
rthick =
n(Aq+)
n(B(r−1)+) + µ(Br+)n(Br+)
η(Aq+)
η(B(r−1)+)
The efficiency factor µ is a measure of how much B(r−1)+ is pro-
duced by charge exchange reactions by Bq+:
σr,r−1(v)
σr(v)
where σr is the total charge exchange cross section and σr,r−1
the one electron charge changing cross section. The efficiency
factor η describes the emission yield per reaction and is given
by the ratio between the relevant emission cross section σem and
the total charge changing cross section σr:
σem(v)
σr(v)
6 D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey
Fig. 7. Collisional opacity effects on flux ratios within the field of view.
The outer bounds of the fields of view within this survey were between
104 − 105 km, as indicated by the shaded area. We considered a 500 km
s−1 equatorial wind interacting with comets with different activities:
Q = 1028 molecules s−1 (dashed lines) and Q = 1029 molecules s−1
(solid lines). All flux ratios are normalized to 1 at infinity.
To explore the effect of collisional opacity on spectra, we
simulated two comets at 1 AU from the Sun, with gas pro-
duction rates of 1028 and 1029 molecules s−1, interacting with
a solar wind with a velocity of 500 km s−1 and an averaged
slow wind composition Schwadron and Cravens (2000). The
results are summarized in Figure 7 where different flux ratios
are shown. The behavior of these ratios as a function of aper-
ture is important because they can be used to derive relative
ionic abundances. All ratios are normalized to 1 at infinite dis-
Fig. 8. Simulated X-ray spectra for a 1029 molecules s−1 comet inter-
acting with an equatorial wind with a velocity of 300 km s−1 for fields
of view decreasing from 105 km (solid line), 104 km (dashed line) and
103 km (dotted line).
tance from the comet’s nucleus. For low activity comets with
Q ≤ 1028 molecules s−1, the collisional opacity does not affect
the comet’s X-ray spectrum. Within typical field of views all line
flux ratios are close to the collisionally thin value. For more ac-
tive comets (Q = 1029 molecules s−1), collisional opacity can
become important within the field of view. Observed flux ratios
involving C v should be treated with care, see e.g. C v/O vii and
C vi/C v, because the flux ratios within the field of view can be
affected by almost 50% and 35%, respectively. The effect is the
strongest in these cases because of the large relative abundance
of C6+, that contributes to the C v emission via sequential elec-
tron capture reactions in the collisionally thick zones. For N vii
and O viii, a small field of view of 104 km could affect the ob-
served ionic ratios by some 20%.
To further illustrate these results, we show the result-
ing X-ray spectra in Fig. 8. There, we consider a Q =
1029 molecules s−1 comet interacting with a 300 km s−1 wind
and show the effect of slowly zooming from the collisionally
thin to the collisionally thick zone around the nucleus. The field
of view decreases from 105 to 103 km. At 105 km, the spectrum
is not affected by collisionally thick emission, whereas the emis-
sion within an aperture of 1000 km is almost purely from the
interactions within the collisionally thick zones of the comet,
which can be most clearly seen by the strong enhancement of
the C v emission around 300 eV.
The results of our model efforts demonstrate that cometary
X-ray spectra reflect characteristics of the comet, the solar wind
and the observational conditions. Firstly, charge exchange cross
sections depend on the velocity of the solar wind, but its effects
are the strongest at velocities below regular solar wind velocities.
Secondly, collisional opacity can affect cometary X-ray spectra
but mainly when an active comet (Q = 1029 molecules s−1) is
observed with a small field of view (≤ 5×104 km). The dominant
factor however to explain differences in cometary CXE spectra
is therefore the state and hence composition of the solar wind.
This implies that the spectral analysis of cometary X-ray spectra
can be used as a direct, remote quantitative and qualitative probe
of the solar wind.
D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey 7
8 D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey
Fig. 9. Chandra comet observations during the descending phase of so-
lar cycle # 23. Monthly sunspot numbers (grey line) and smoothed
monthly sunspot number (black lines) from the Solar Influences Data
Analysis Center of the Department of Solar Physics, Royal Observatory
of Belgium (http://sidc.oma.be/). Letters refer to the chronological or-
der of observation.
4. Observations
In this section, we will briefly introduce the different comet ob-
servations performed with Chandra. A summary of comet and
solar wind parameters is given in Table 3. More observational
details on the comet and a summary of the state of the solar wind
at the location of the comet during the X-ray observations can be
found in Appendix A.
4.1. Solar Wind Data
Our survey spans the whole period between solar maximum (mid
2000) and solar minimum (mid 2006), see Fig. 9. During solar
minimum, the solar wind can be classified in polar- and equa-
torial streams, where the polar can be found at latitudes larger
than 30◦ and the equatorial wind within 15◦ of the helioequator.
Polar streams are fast (ca. 700 km s−1) and show only small vari-
ations in time, in contrast to the irregular equatorial wind. Cold,
fast wind is also ejected from coronal holes around the equa-
tor, and when these streams interact with the slower background
wind corotating interaction regions (cirs) are formed. As was
illustrated by Schwadron and Cravens (2000), different wind
types vary greatly in their compositions, with the cooler, fast
wind consisting of on average lower charged ions than the hot-
ter equatorial wind. This clear distinction disappears during solar
maximum, when at all latitudes the equatorial type of wind dom-
inates. In addition, coronal mass ejections are far more common
around solar maximum.
There is a strong variability of heavy ion densities due to
variations in the solar source regions and dynamic changes in
the solar wind itself Zurbuchen & Richardson (2006). The vari-
ations mainly concern the charge state of the wind as elemental
variations are only on the order of a factor of 2 (Von Steiger et
al. (2000), and references therein).
We obtained solar wind data from the online data archives
of ace (proton velocities and densities from the swepam instru-
ment, heavy ion fluxes from the swics and swims instruments1)
and soho (proton fluxes from the Proton Monitor Instrument2).
Both ace and soho are located near Earth, at its Lagrangian
1 http://www.srl.caltech.edu/ace/ASC/level2/index.html
2 http://umtof.umd.edu/pm/crn/
point L1. In order to map the solar wind from L1 to the posi-
tion of the comets, we used the time shift procedure described
by Neugebauer et al. (2000). The calculations are based on
the comet ephemeris, the location of L1 and the measured wind
speed. With this procedure, the time delay between an element
of the corotating solar wind arriving at L1 and the comet can be
predicted. A disadvantage of this procedure is that it cannot ac-
count for latitudinal structures in the wind or the magnetohydro-
dynamical behavior of the wind (i.e., the propagation of shocks
and cmes). These shortcomings imply that especially for comets
that have large longitudinal, latitudinal and/or radial separations
from Earth, the solar wind data is at best an estimate of the local
wind conditions. The resulting proton velocities at the comets
near the time of the Chandra observations are shown in Fig. 10.
Parallel to this helioradial and heliolongitudinal mapping, we
compared our comet survey to a 3D mhd time–dependent so-
lar wind model that was employed during most of Solar Cycle
23 (1997 - 2006) on a continuous basis when significant solar
flares were observed. The model (reported by Fry et al. (2003);
McKenna-Lawlor et al. (2006) and Z.K. Smith, private com-
munication, for, respectively, the ascending, maximum, and de-
scending phases) treats solar flare observations and maps the
progress of interplanetary shocks and cmes. The papers men-
tioned above provide an rms error for ”hits” of ±11 hours Smith
et al. (2000); McKenna-Lawlor et al. (2006). cir fast forward
shocks were also taken into account in order to differentiate be-
tween the co-rotating ”quiet” and transient structures. It was im-
portant, in this differentiating analysis, to examine (as we have
done here) the ecliptic plane plots of both of these structures as
simulated by the deforming interplanetary magnetic field lines
(see, for example, Lisse et al. (2005, 2007) for several of the
comets discussed here.) Therefore, the various comet locations
(Table 3) were used to estimate the probability of their X-ray
emission during the observations being influenced by either of
these heliospheric situations.
4.2. X-ray Observations
After its launch in 1999, 8 comets have been observed with
the Chandra X-ray Observatory and Advanced ccd Imaging
Spectrometer (acis). Here, we have mainly considered obser-
vations made with the acis-S3 chip, which has the most sensi-
tive low energy response and for which the majority of comets
were centered. The Chandra’s acis-S instrument provides mod-
erate energy resolution (σ ≈ 50 eV) in the 300 to 1500 eV en-
ergy range, the primary range for the relatively soft cometary
emission. All comets in our sample were re-mapped into comet-
centered coordinates using the standard Chandra Interactive
Analysis of Observations (ciao v3.4) software ‘sso freeze’ al-
gorithm.
Comet source spectra were extracted from the S3 chip with
a circular aperture with a diameter of 7.5′, centered on the
cometary emission. The exception was comet C/2001 Q4, which
filled the chip and a 50% larger aperture was used. acis’ re-
sponse matrices were used to model the instrument’s effective
area and energy dependent sensitivity matrices were created for
each comet separately using the standard ciao tools.
Due to the large extent of cometary X-ray emission, and
Chandra’s relatively narrow field of view, it is not trivial to ob-
tain a background uncontaminated by the comet and sufficiently
close in time and viewing direction. We extracted background
spectra using several techniques: spectra from the S3 chip in
an outer region generally > 8′, an available acis S3 blank sky
observation, and backgrounds extracted from the S1 ccd. For
D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey 9
Fig. 10. Solar wind proton velocities estimated from ace and soho data. For all comets, the time of the observations is indicated with a dotted line.
Letters refer to the chronological order of observation.
several comets there are still a significant number of cometary
counts in the outer region of the S3 ccd. Background spec-
tra taken from the S1 chip have the advantage of having been
taken simultaneous with the S3 observation and thus having
the same space environment as the S3 observation. In general
the background spectra were extracted with the same 7.5′ aper-
ture as the source spectra but centered on the S1 chip. For
comet Encke, where the S1 chip was off during the observa-
tion the background from the outer region of the S3 chip was
used. Comet C/2000 WM1 (linear) was observed with the Low-
Energy Transmission Grating (letg) and acis-S array. For the
latter, we analyzed the zero-th order spectrum, and used a back-
ground extracted from the outer region of the S3 chip. It is possi-
ble that the proportion of incident X-rays diffracted onto the S3
chip will vary with photon energy. Background-subtracted spec-
tra generally have a signal-to-noise at 561 eV of at least 10, and
over 50 for 153P/2002 C1 (Ikeya–Zhang).
5. Spectroscopy
The observed spectra are shown in Figure 11. The spectra
suggest a classification based upon three competing emission
features, i.e. the combined carbon and nitrogen emission (be-
low 500 eV), O vii emission around 565 eV and O viii emis-
sion at 654 eV. Firstly, the C+N emission (<500 eV) seems to
be anti-correlated with the oxygen emission. This clearly sets
the spectra of 73P/2006 S.–W.3B and 2P/2003 (Encke) apart, as
for those two comets the C+N features are roughly as strong as
the O vii emission. In the spectra of the remaining five comets,
oxygen emission dominates over the carbon and nitrogen emis-
sion below 500 eV. The O viii/O vii ratio can be seen to increase
continuously, culminating in the spectrum of 153P/2002 (Ikeya–
Zhang) where the spectrum is completely dominated by oxygen
emission with almost comparable O viii and O vii emission fea-
tures. From our modelling, we expect that the separate classes
reflect different states of the solar wind, which imply different
ionic abundances. To explore the obtained spectra more quanti-
tatively, we will use a spectral fitting technique based on our cxe
model to extract X-ray line fluxes.
5.1. Spectral Fitting
The charge exchange mechanism implies that cometary X-ray
spectra result from a set of solar wind ions, which produce at
least 35 emission lines in the regime visible with Chandra. As
comets are extended sources, these lines cannot all be resolved.
All spectra were therefore fit using the 6 groups of fixed lines of
our cxe model (see Table 2) and spectral parameters were de-
rived using the least squares fitting procedure with the xspec
package. The relative strengths from all lines were fixed per
ionic species, according to their velocity dependent emission
cross sections. Thus, the free parameters were the relative fluxes
of the C, N and O ions contained in our model.
Two additional Ne lines at 907 eV (Ne ix) and 1024 eV
(Ne x) were also included, giving a total of 8 free parameters.
All line widths were fixed at the acis-S3 instrument resolution.
The spectra were fit in the 300 to 1000 eV range. This pro-
vided 49 spectral bins, and thus 41 degrees of freedom. acis
spectra below 300 eV are discarded because of the rising back-
ground contributions, calibration problems and a decreased ef-
fective area near the instrument’s carbon edge.
As a more detailed example of the cxe model and compari-
son to the data, we show in Fig 12 the acis-S3 data for C/1999
S4 (linear). The figure shows the background subtracted source
spectrum over-plotted with the background spectrum, the differ-
ence between the model and data, and the model spectrum and
10 D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey
Fig. 11. Observed spectrum and fit of all 8 comets observed with
Chandra, grouped by their spectral shape (see text). The histogram lines
indicate the CXE model fit.
data to indicate to contribution of the different ions. Only the
emission lines with >3% strength of the strongest line in their
species are shown for ease of presentation.
The fluxes obtained by our fitting are converted into relative
ionic abundances by weighting them by their velocity dependent
emission cross sections. For comets observed near the ecliptic
plane (< 15◦), solar wind conditions mapped to the comet were
used (Section 4.1). For comets observed at higher latitudes, these
data are most likely not applicable and a solar wind velocity of
500 km s−1 was assumed.
Fig. 12. Details of the cxe fit for the spectrum of comet 1999/S4
(linear). Top panel: Comet (filled triangles) and background (open
squares) spectrum. Middle panel: Residuals of cxe fit Bottom panel:
cxemodel and observed spectrum indicating the different lines and their
strengths. Carbon - red; nitrogen - orange; oxygen - blue; neon - green.
The unfolded model is scaled above the emission lines for the ease of
presentation.
Fig. 13. Parameter sensitivity for the major emission features in the fit
of C/1999 S4 (linear), with respect to the O vii 561 eV feature. All units
are 10−4 photons cm−2 s−1. The contours indicate a χ2R of 9.2 (or 99%
confidence, largest, green contour), a χ2R of 4.6 (90%, red contour) and
a χ2R of 2.3 (68%, smallest, blue contour).
5.2. Spectroscopic Results
The fits to all cometary spectra are shown in Fig. 11 and the
results of the fits are given in Table 4. For the majority of the
comets, the model is a good fit to the data within a 95% con-
fidence limit (χ2R ≈ 1.4). Results for comet 153P/2002 (Ikeya–
Zhang) are presented in Table 5 with an additional systematic
error to account for its brightness and any uncertainties in the
response.
The spectra for all comets are well reproduced in the 300 to
1000 eV range. The nitrogen contribution is statistically signifi-
cant for all comets except the fainter ones, 2P/2003 (Encke) and
D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey 11
12 D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey
Table 6. Solar wind abundance relative to O7+, obtained for comet linear S4. References: Bei ’03 – Beiersdorfer et al. (2003), Kra ’04 –
Krasnopolsky (2004), Kra ’06 – Krasnopolsky (2006), Otr ’06 – Otranto et al (2006) and S&C ’00 – Schwadron and Cravens (2000). Dots
indicate that an ion was included in the fitting, but no abundances were derived; dash means that an ion was not included in the fitting. Otranto et
al (2006) did not fit the observed spectrum, but used a combination of ace-data and solar wind averages from Schwadron and Cravens (2000) to
compute a syntectic spectrum of the comet. Solar wind averages are given for comparison Schwadron and Cravens (2000)
Ion this work Bei 03 Kra 04 Kra 06 Otr 06
O8+ 0.32 ± 0.03 0.13 ± 0.03 0.13 ± 0.05 0.15 ± 0.03 0.35 0.35
C6+ 1.4 ± 0.4 0.9 ± 0.3 0.7 ± 0.2 0.7 ± 0.2 1.02 1.59
C5+ 12 ± 4.0 11 ± 9 . . . 1.7 ± 0.7 1.05 1.05
N7+ 0.07 ± 0.06 0.06 ± 0.02 – – 0.03 0.03
N6+ 0.63 ± 0.21 0.5 ± 0.3 – – 0.29 0.29
Ne10+ 0.02 ± 0.01 – – – – –
Ne9+ . . . – (15 ± 6) × 10−3 (20 ± 7) × 10−3 – –
73P/2006 (S.-W.3B). For example, removing the nitrogen com-
ponents from linear S4’s cxe model and re-fitting, increases χ2R
to over 7.
χ2 contours for C/1999 S4 (linear) are presented in Fig 13.
The line strengths for each ionic species are generally well con-
strained, except where spectral features overlap. This can be
readily seen when comparing the contours for the N vii 500 eV
and O vii 561 eV features where a strong anti-correlation exists
(Figure 12). Due to the limited resolution of acis an increase in
the N vii feature will decrease the O vii strength. Similar anti-
correlations exist between the nitrogen N vi or N vii and C v
299 eV lines. Since the line strength for the main line in each
ionic species is linked to weaker lines, a range of energies can
contribute and better constrain its strength. However with O vii
as the strongest spectral feature the nitrogen and carbon compo-
nents may be artificially lower as a result of the aforementioned
anti-correlations. The lack of effective area due to the carbon
edge in the acis response also may over-estimate the C v line
flux. The neon features were well constrained for the brighter
comets, but this is a region of lower signal and some caution
must be taken when treating the neon line strengths and they are
included here largely for completeness.
In the case of 153P/Ikeya–Zhang, the χ2R > 1.4. The main
discrepancy is that the model produces not enough flux in the
700 to 850 eV range compared to the observed spectrum. This
may reflect an underestimation of higher O viii transitions or the
presence of species not (yet) included in the model, such as Fe.
This will be discussed further in the last section of this paper and
in a separate paper dedicated to the observations of this comet
(K. Dennerl, private communication).
One of the best studied comets is C/1999 S4 (linear), be-
cause of its good signal-to-noise ratio. To discuss our results, we
will compare our findings with earlier studies of this comet. In
general, the spectra analyzed here have more counts than ear-
lier analyzes, because of improvements in the Chandra process-
ing software and because we took special care to use a back-
ground that is as comet-free as possible. Previous studies appear
to have removed true comet signal when the background subtrac-
tion was performed. In particular, both the Krasnopolsky (2004)
and Lisse et al. (2001) studies used background regions from
the outer part of the S3 chip and this may have still had true
cometary emission. Krasnopolsky (2004) subtracted over 70%
of the total signal as background. We find that using the S1-chip,
the background contributes only 20% of the total counts.
Different attempts to derive relative ionic abundances from
C1999/S4’s X-ray spectrum are compared in Table 6. Our atomic
physics based spectral analysis combines the benefits of ear-
lier analytical approaches by Kharchenko and Dalgarno (2000,
2001); Beiersdorfer et al. (2003). These methods were all ap-
plied to just one or two comets. Beiersdorfer et al. (2003) inter-
pret C1999/S4’s X-ray spectrum by fitting it with 6 experimental
spectra obtained with their ebit setup. The resulting abundances
are very similar to ours. The advantage of their method is that
it includes multiple electron capture, but in order to observe the
forbidden line emission, the spectra were obtained with trapped
ions colliding at CO2, at collision energies of 200 to 300 eV or
ca. 30 km s−1. As was shown in Fig. 3, the cxe hardness ratio
may change rapidly below 300 km s−1, implying an overesti-
mation of the higher order lines compared to the n = 2 → 1
transition, which for O vii overlap with the O viii emission. We
therefore find higher abundances of O8+.
Krasnopolsky (2004, 2006) obtained fluxes and ionic abun-
dances by fitting the spectrum with 10 lines of which the energies
were semi-free. Their analysis thus does not take the contamina-
tion of unresolved emission into account, and N vi and N vii are
not included in the fit. The line energies were attributed to cxe
lines of mainly solar wind C and O but also to ions of Mg and Ne.
The inclusion of the resulting low energy emission (near 300 eV)
results in lower C5+ fluxes (see also Otranto et al (2006)).
There are several factors that may contribute to the unexpect-
edly low C vi/C v ratios: 1) There may be a small contribution to
the C v line from other ions in the 250-300 eV range (e.g. Si,
Mg, Ne) that are currently not included in the model. Including
these species in the model would lower the C v flux, but proba-
bly only with a small amount. 2) The low acis effective area in
the 250-300 eV region allows the C v flux to be unconstrained,
and this increases the uncertainty in the C v flux. We estimate
that the uncertainty in the effective area, introduced by the car-
bon edge, can account for an uncertainty as large as a factor of
10 in the observed C v/C vi ratios.
We will not compare our results with measured ace/swics
ionic data. As discussed in section 4, the solar wind is highly
variable in time and its composition can change dramatically
over the course of less than a day. Variations in the solar wind’s
ionic composition are often more than 50% during the course of
an observation. Data on N, Ne, and O8+ ions have not been well
documented as the errors of these abundances are dominated by
counting statistics. As discussed above, latitudinal and corota-
tional separations imply large inaccuracies in any solar wind
mapping procedure. These conditions clearly disfavor modelling
based on either average solar wind data or ace/swics data.
6. Comparative Results
As noted in Section 5, spectral differences show up in the be-
havior of the low energy C+N emission (< 500 eV), the O vii
D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey 13
Table 7. Correlation between classification according to spectral shape and comet/solar wind characteristics during the observations. Comet
families from Marsden & Williams (2005). Phase refers to where in the solar cycle the comet was observed, where 1 is the solar maximum and 0
the solar minimum of cycle #23’s descending phase. For other references, see Table 3.
Class # Comet Comet Q Latitude Wind Type
Family (1028 mol. s−1)
cold H 73P/2006 (S.-W.3B) Jupiter 2 0.5 cir
E 2P/2003 (Encke) Jupiter 0.7 11.4 Flare/PS
warm F C/2001 Q4 (neat) unknown 10 -3 Quiet
G 9P/2005 (Tempel 1) Jupiter 0.9 0.8 Quiet
hot C C/2000 WM1 (linear) unknown 3-9 -34 PS
A C/1999 S4 (linear) unknown 3 24 icme
B C/1999 T1 (McNaught–Hartley) unknown 6-20 15 Flare/cir
D C/2002 C1 (Ikeya–Zhang) Oort 20 26 icme
Fig. 14. Flux ratios of all observed comets. The low energy C+N feature
is anti-correlated to the oxygen ionic ratio. Letters refer to the chrono-
logical order of observation.
Fig. 15. Ion ratios of all observed comets. The C+N ionic abundantie is
anti-correlated to the oxygen ionic ratio. Letters refer to the chronolog-
ical order of observation.
emission at 561 eV and the O viii emission at 653 eV. Figure 14
shows a color plot of the fluxes of these three emission features,
and Figure 15 the corresponding ionic abundances. There is a
clear separation between the two comets with a large C+N con-
tribution and the other ‘oxygen-dominated’ comets, which on
their turn show a gradual increase in the oxygen ionic ratio. This
sample of comet observations suggest that we can distinguish
two or three spectral classes.
Table 7 surveys the comet parameters for the different spec-
tral classes. The outgassing rate, heliocentric- or geocentric dis-
tance and comet family do not correlate to the different classes,
in accordance with our model findings. The data does suggest
a correlation between latitude and wind conditions during the
observations. At first sight, the apparent correlation between lat-
itude and oxygen ratio seems paradoxical. According to the bi-
modal structure of the solar wind the fast, cold wind dominates
at latitudes > 15◦, implying less O viii emission. In Figure 9,
the comet observations are shown with respect to the phase of
the last solar cycle. Interestingly, we note that all comets that
were observed at higher latitudes were observed around solar
maximum. The solar wind is highly chaotic during solar maxi-
mum and the frequency of impulsive events like CMEs is much
higher than during the descending and minimum phase of the
cycle. This explains both why the comets observed in the period
2000–2002 encountered a disturbed solar wind and why our sur-
vey does not contain a sample of the cool fast wind from polar
coronal holes.
The observed classification can therefore be fully ascribed
to solar wind states. The first class is associated with cold, fast
winds with lower average ionization. These winds are found in
cirs and behind flare related shocks. The spectra due to these
winds are dominated by the low energy x-rays, because of the
low abundances of highly charged oxygen. At the relevant tem-
peratures, most of the solar wind oxygen is He-like O6+, which
does not produce any emission visible in the 300–1000 eV
regime accessible with Chandra. Secondly, there is an interme-
diate class with two comets that were all observed during periods
of quiet solar wind. These comets interacted with the equatorial,
warm slow wind. The third class then comprises comets that in-
teracted with a fast, hot, disturbed wind associated with icmes
or flares. From the solar wind data, Ikeya–Zhang was proba-
bly the most extreme example of this case. This comet had 10
times more signal than any other comet in our sample and small
discrepancies in the response may be important at this level.
Extending into the 1-2 keV regime, a preliminary analysis indi-
cates the presence of bare and H-like Si, Mg and Fe xv-xx ions, in
accordance with ace measurements of icme compositions Lepri
& Zurbuchen (2004).
The variability and complex nature of the solar wind allows
for many intermediate states in between these three categories
Zurbuchen et al. (2002), which explain the gradual increase
of the O viii/O vii ratio that we observed in the cometary spec-
tra. As the solar wind is a collisionless plasma, the charge state
14 D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey
Fig. 16. Spectrum derived ionic oxygen ratios and corresponding
freezing-in temperatures from Mazotta et al. (1998). The shaded area
indicates the typical range of slow wind associated with streamers.
Letters refer to the chronological order of observation.
distribution in the solar wind is linked to the temperature in its
source region. Ionic temperatures are therefore a good indicator
of the state of the wind encountered by a comet. The ratio be-
tween O7+ and O6+ ionic abundances has been demonstrated to
be a good probe of solar wind states. Zurbuchen et al. (2002)
observed that slow, warm wind associated with streamers typi-
cally lies within 0.1 < O7+/O6+ < 1.0, corresponding to freez-
ing in temperatures of 1.3–2.1 MK. The corresponding temper-
ature range is indicated in the Figure 16. In the figure, we show
the observed O8+ to O7+ ratios and the corresponding freezing-
in temperatures from the ionizational/recombination equilibrium
model by Mazotta et al. (1998). Most observations are within or
near to the streamer-associated range of oxygen freezing in tem-
peratures. Four comets interacted with a wind significantly hot-
ter than typical streamer winds, and in all four cases we found
evidence in solar wind archives that the comets most likely en-
countered a disturbed wind.
7. Conclusions
Cometary X-ray emission arises from collisions between bare-
and H-like ions (such as C, N, O, Ne, . . . ) with mainly water and
its dissociation products OH, O and H. The manifold of depen-
dencies of the cxe mechanism on characteristics of both comet
and wind offers many diagnostic opportunities, which are ex-
plored in the first part of this paper. Charge exchange cross sec-
tions are strongly dependent on the velocity of the solar wind,
and these effects are strongest at velocities below the regular
wind conditions. This dependency might be used as a remote
plasma diagnostics in future observations. Ruling out collisional
opacity effects, we used our model to demonstrate that the spec-
tral shape of cometary cxe emission is in the first place deter-
mined by local solar wind conditions. Cometary X-ray spectra
hence reflect the state of the solar wind.
Based on atomic physic modelling of cometary charge ex-
change emission, we developed an analytical method to study
cometary X-ray spectra. First, the data of 8 comets observed
with Chandra were carefully reprocessed to avoid the subtrac-
tion of cometary signal as background. The spectra were then
fit using an extensive data set of velocity dependent emission
cross sections for eight different solar wind species. Although
the limited observational resolution currently available hampers
the interpretation of cometary X-ray spectra to some degree, our
spectral analysis allows for the unravelling of cometary X-ray
spectra and allowed us to derive relative solar wind abundances
from the spectra.
Because the solar wind is a collisionless plasma, local ionic
charge states reflect conditions of its source regions. Comparing
the fluxes of the C+N emission below 500 eV, the O vii emission
and the O viii emission yields a quantitative probe of the state
of the wind. In accordance with our modelling, we found that
spectral differences amongst the comets in our survey could be
very well understood in terms of solar wind conditions. We are
able to distinguish interactions with three different wind types,
being the cold, fast wind (I), the warm, slow wind (II); and the
hot, fast, disturbed wind (III). Based on our findings, we pre-
dict the existence of even cooler cometary X-ray spectra when a
comet interacts with the fast, cool high latitude wind from polar
coronal holes. The upcoming solar minimum offers the perfect
opportunity for such an observation.
Acknowledgements. DB and RH acknowledge support within the framework of
the fom–euratom association agreement and by the Netherlands Organization
for Scientific Research (nwo). MD thanks the noaa Space Environment Center
for its post-retirement hospitality. We are grateful for the cometary ephemerides
of D. K. Yeomans published at the jpl/horizons website. Proton velocities used
here are courtesy of the soho/celias/pm team. soho is a mission of international
cooperation between esa and nasa. Chianti is a collaborative project involving
the nrl (USA), ral (UK), mssl (UK), the Universities of Florence (Italy) and
Cambridge (UK), and George Mason University (USA).
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Appendix A: Observations within this Survey
This Appendix presents the observational details of the Chandra
data and the corresponding solar wind state. The prefix ’FF’
(fearless forecast) used in this appendix refers to the real time
forecasting of coronal mass ejection shocks arrivals at Earth. The
numbers were so-named for flare/coronal shock events during
solar cycle #23.
A.1. C/1999 S4 (linear)
X-rays. The first Chandra cometary observation was of comet
C/1999 S4 (linear) Lisse et al. (2001), with observations being
made both before and after the breakup of the nucleus. Due to the
low signal-to-noise ratio of the second detection, only the July
14th 2000 pre-breakup observation is discussed here. Summing
the 8 pointings of the satellite gave a total time interval of 9390 s.
In this period, the acis-S3 ccd collected a total of 11 710 photons
were detected in the range 300–1000 eV. Detections out side this
range or on other acis-ccds were not attributed to the comet. As
a result, data from the S1-ccd (which is configured identically to
S3) may be used as an indicator of the local X-ray background.
The morphology can be described by a crescent shape, with
the maximum brightness point 24 000 km from the nucleus
on the Sun-facing side. The brightness dims to 10% of the
maximum level at 110 000 km from the nucleus.
Solar wind. A large velocity jump can be seen around DoY 199,
which was due to the famous ”Bastille Day” flare on 14 July
(FF#153, Dryer et al (2001); Fry et al. (2003)). This flare
reached the comet only after the first observation. At July 12,
2017UT a solar flare started at N17W65 (FF#152), which was
nicely placed to hit this comet with a very high probability dur-
ing the first observations Fry et al. (2003). As for the second
observation, there was another flare on July 28, S17E24, at 1713
UT (FF#164) and there was a high probability that its shock’s
weaker flank hit the comet.
A.2. C/1999 T1 (McNaught–Hartley)
X-rays. The allocated observing time of comet McNaught–
Hartley was partitioned into 5 one-hour-slots between January
8th and January 15th, 2001 Krasnopolsky et al. (2002). The
strongest observing period was on January 8th, when ∆ = 1.37
AU and rh = 1.26 AU.
There were 15 000 total counts observed by the acis-S3 ccd
between 300 and 1000 eV. The emission region can be described
by a crescent, with the peak brightness is at 29 000 km from
the nucleus. The brightness dims to 10% of the maximum at a
cometocentric distance of 260 000 km. Again, the acis-S1 ccd
may be used to indicate the local background signal.
Solar wind. The comet was not within the heliospheric cur-
rent/plasma sheet (HCS/HPS). Two corotating cirs are probably
associated with the first two observations. Two flares (FF#233
and #234) took place; however, another corotating cir more
likely arrived before the flare’s transient shock’s effects did
McKenna-Lawlor et al. (2006).
A.3. C/2000 WM1 (linear)
X-rays. The only attempt to use the high-resolution grat-
ing capability of the acis-S array was made with comet
C/2000 WM1 (linear). Here, the Low-Energy Transmission
Grating (letg) was used. The dimness of the observed X-rays,
and the extended nature of the emitting atmosphere meant that
the grated spectra did not yield significant results. It is still
16 D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey
possible to extract a spectrum based on the pulse-heights gener-
ated by each X-ray detection on the acis-S3 chip, although the
morphology is not recorded. 6300 total counts were recorded for
the pulse-height spectrum of the S3 chip in the 300 to 1000 eV
range.
Solar wind. Comet WM1 was observed at the highest latitude
available within this survey, and at a latitude of 34 degrees, it was
far outside the hcs. During the observations, this comet might
have experienced the southerly flank of the shock of a strong
X3.4 flare at S20E97 and its icme and shock on December 28,
2001 (FF#359) McKenna-Lawlor et al. (2006).
A.4. 153P/2002 (Ikeya–Zhang)
X-rays. The brightest X-ray comet in the Chandra archive is
153P/2002 (Ikeya–Zhang). The heliographic latitude, geocentric
distance and heliocentric distance were comparable to those for
comet C/1999 S4 (linear), with a latitude of 26◦, ∆ = 0.457 AU
and rh = 0.8 AU. Rather than periodically re-point the detector
to track the comet, the pointing direction was fixed and the
comet was monitored as it passed through the field of view, thus
increasing the effective FoV. There were two observing periods
on April 15th 2002, each lasting for approximately 3 hours and
15 minutes. In both periods, a strong cometary signal is detected
on all of the activated acis-ccds. Consequently, a background
signal cannot be taken from the observation. A crescent shape
on the Sun side of the comet is observed over all of the ccd
array. Over 200 000 total counts were observed from the S3
chip in the 300 to 1000 eV range. The time intervals for each
observing period are 11 570 and 11 813 seconds.
Solar wind. Like C/2000 WM1, this comet was observed at a
relatively high heliographic latitude. Solar wind data obtained
in the ecliptic plane can therefore not be used to determine the
wind state at the comet. 153P/2002 (Ikeya–Zhang) was well-
positioned during the first observation on 15 April 2002 for a
flare at N16E05 (FF#388) on 12 April 2002. During the second
observation on 16 April, there was an earlier flare on 14 April at
N14W57, but this flare was probably too far to the west to be ef-
fective McKenna-Lawlor et al. (2006). The comet was observed
at a high latitude, and hence ace solar wind data is most likely
not applicable.
A.5. 2P/2003 (Encke)
X-rays. The Chandra observation of Encke took place on the
24th of November 2003 Lisse et al. (2005), when the comet
had a heliocentric distance of rh = 0.891 AU and a geocentric
distance of ∆ = 0.275 AU and a heliographic latitude of 11.4
degrees. The comet was continuously tracked for over 15 hours,
resulting in a useful exposure of 44 000 seconds. The acis-S3 ccd
counted 6140 X-rays in the range 300–1000 eV.
The brightest point was offset from the nucleus by
11 000 km, dimming to 10% of this value at a distance of
60 000 km.
The acis-S1 ccd was not activated in this observation. The
low quantum efficiency of the other activated ccds below 0.5 keV
makes them unsuitable as background references.
Solar wind. The proton velocity decreased during observations
from 600 km s−1 to 500 km s−1. A flare on 20 November
2003, at N01W08 (FF#525), was well-positioned to affect the
observations on 23 November (data from work in progress by
Z.K. Smith et al.). The comet most likely interacted with the
overexpanded, rarified plasma flow that followed the earlier hot
shocked and compressed flow behind the flare’s shock.
A.6. C/2001 Q4 (neat)
X-rays. A short observation of comet C/2001 Q4 was made on
May 12 2004, when the geocentric and heliocentric distances
were ∆ = 0.362 AU and rh = 0.964 AU respectively. With a
heliographic latitude of 3 degrees, the comet was almost in
the ecliptic plane. From 3 pointings, the useful exposure was
10 328 seconds. The acis-S3 chip detected 6540 X-rays in be-
tween 300 and 1000 eV. The acis-S1 was used as a background
signal.
Solar wind. There was no significant solar activity during the
observations (Z.K. Smith et al., ibid.). From solar wind data, the
comet interacted with a quiet, slow 352 km s−1 wind.
A.7. 9P/2005 (Tempel 1)
X-rays. The observation of comet 9P/2005 (Tempel 1) was de-
signed to coincide with the Deep Impact mission Lisse et al.
(2007). The allocated observation time of 291.6 ks was split
into 7 periods, starting on June 30th, July 4th (encompassing the
Deep Impact collision), July 5th, July 8th, July 10th, July 13th
and July 24th. The brightest observing periods were June 30th
and July 8th. The focus here is on the June 30th observation. On
this date, rh = 1.507 AU and ∆ = 0.872 AU.
The useful exposure was 50 059 seconds, with a total of 7300
counts, 4000 from the June 30th flare alone, were detected in the
energy range of 300–1000 eV.
The brightest point for the June 30th observation was
located 11 000 km from the nucleus. The morphology appears
to be more spherical than in other comet observations.
Solar wind. Observations were taken over a long time span cov-
ering different solar wind environments. There was no signifi-
cant solar activity during the 30 June 2005 observations (Z.K.
Smith et al., ibid. Lisse et al. (2007)). From the ace data, it can
be seen that at June 30, the comet most likely interacted with a
quiet, slow solar wind.
A.8. 73P/2006 (Schwassmann–Wachmann 3B)
X-rays. The close approach of comet 73P/2006 (Schwassmann–
Wachmann 3B) in May 2005 (∆ = 0.106 AU, rh = 0.965 AU)
provided an opportunity to examine cometary X-rays in high
spatial resolution. Chandra was one of several X-ray missions to
focus on one of the large fragments of the comet. Between 300
and 1000 eV, 6285 counts were obtained in a useful exposure of
20 600 seconds.
Solar wind. There was a weak flare on 22 May 2006 (FF#655,
Z.K. Smith, priv. comm.). A sequence of three high speed coro-
nal hole streams passed the comet in the period around the ob-
servations and a corotating cir might have reached the comet in
association with the observations on 23 May, which is confirmed
by the mapped solar wind data.

  We present results of the analysis of cometary X-ray spectra with an extended
version of our charge exchange emission model (Bodewits et al. 2006). We have
applied this model to the sample of 8 comets thus far observed with the Chandra
X-ray observatory and ACIS spectrometer in the 300-1000 eV range. The surveyed
comets are C/1999 S4 (LINEAR), C/1999 T1 (McNaught-Hartley), C/2000 WM1
(LINEAR), 153P/2002 (Ikeya-Zhang), 2P/2003 (Encke), C/2001 Q4 (NEAT), 9P/2005
(Tempel 1) and 73P/2006-B (Schwassmann-Wachmann 3) and the observations include
a broad variety of comets, solar wind environments and observational
conditions. The interaction model is based on state selective, velocity
dependent charge exchange cross sections and is used to explore how cometary
X-ray emission depend on cometary, observational and solar wind
characteristics. It is further demonstrated that cometary X-ray spectra mainly
reflect the state of the local solar wind. The current sample of Chandra
observations was fit using the constrains of the charge exchange model, and
relative solar wind abundances were derived from the X-ray spectra. Our
analysis showed that spectral differences can be ascribed to different solar
wind states, as such identifying comets interacting with (I) fast, cold wind,
(II), slow, warm wind and (III) disturbed, fast, hot winds associated with
interplanetary coronal mass ejections. We furthermore predict the existence of
a fourth spectral class, associated with the cool, fast high latitude wind.

Introduction
When highly charged ions from the solar wind collide on a
neutral gas, the ions get partially neutralized by capturing elec-
trons into an excited state. These ions subsequently decay to the
ground state by the emission of one or more photons. This pho-
ton emission is called charge exchange emission (cxe) and it has
been observed from comets, planets and the interstellar medium
in X-rays and the Far-UV Lisse et al. (1996); Krasnopolsky
(1997); Snowden et al. (2004); Dennerl (2002). The spec-
tral shape of the cxe depends on properties of both the neutral
Send offprint requests to: D. Bodewits
gas and the solar wind and the subsequent emission can there-
fore be regarded as a fingerprint of the underlying interactions
Cravens et al. (1997); Kharchenko and Dalgarno (2000, 2001);
Beiersdorfer et al. (2003); Bodewits et al. (2004a, 2006).
Since the first observations of cometary X-ray emission,
more than 20 comets have been observed with various X-ray
and Far-UV observatories Lisse et al. (2004); Krasnopolsky et
al. (2004). This observational sample contains a broad variety of
comets, solar wind environments and observational conditions.
The observations clearly demonstrate the diagnostics available
from cometary charge exchange emission.
First of all, the emission morphology is a tomography of
the distribution of neutral gas around the nucleus Wegmann et
2 D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey
al. (2004). Gaseous structures in the collisionally thin parts
of the coma brighten, such as the jets in 2P/Encke Lisse et al.
(2005), the Deep Impact triggered plume in 9P/Tempel 1 Lisse
et al. (2007) and the unusual morphology of comet 6P/d’Arrest
Mumma et al (1997). In other comets, the X-ray emission
clearly mapped a spherical gas distribution. This resulted in a
characteristic crescent shape for larger and hence collisionally
thick comets observed at phase angles of roughly 90 degrees
(e.g. Hyakutake - Lisse et al. (1996), linear S4 - Lisse et al.
(2001)). Macroscopic features of the plasma interaction such as
the bowshock are observable, too Wegmann & Dennerl (2005).
Secondly, by observing the temporal behavior of the
comets X-ray emission, the activity of the solar wind and
comet can be monitored. This was first shown for comet
C/1996 B2 (Hyakutake) Neugebauer et al. (2000) and re-
cently in great detail by long term observations of comet
9P/2005 (Tempel 1) Willingale et al. (2006); Lisse et al.
(2007) and 73P/2006 (Schwassmann–Wachmann 3C) Brown et
al. (2007), where cometary X-ray flares could be assigned to
either cometary outbursts and/or solar wind enhancements.
Thirdly, cometary spectra reflect the physical characteristics
of the solar wind; e.g. spectra resulting from either fast, cold
(polar) wind and slow, warm equatorial solar wind should be
clearly different Schwadron and Cravens (2000); Kharchenko
and Dalgarno (2001); Bodewits et al. (2004a). Several attempts
were made to extract ionic abundances from the X-ray spectra.
The first generation spectral models have all made strong
assumptions when modelling the X-ray spectra Haeberli et al
(1997); Wegmann et al. (1998); Kharchenko and Dalgarno
(2000); Schwadron and Cravens (2000); Lisse et al. (2001);
Kharchenko and Dalgarno (2001); Krasnopolsky et al. (2002);
Beiersdorfer et al. (2003); Wegmann et al. (2004); Bodewits et
al. (2004a); Krasnopolsky (2004); Lisse et al. (2005). Here,
we present a more elaborate and sophisticated procedure to an-
alyze cometary X-ray spectra based on atomic physics input,
which for the first time allows for a comparative study of all
existing cometary X-ray spectra. In Section 2, our comet-wind
interaction model is briefly introduced. In Section 3, it is demon-
strated how cometary spectra are affected by the velocity and
target dependencies of charge exchange reactions. In Section 4,
the various existing observations performed with the Chandra
X-ray Observatory, as well as the solar wind data available are
introduced. Based upon our modelling, we construct an analyt-
ical method of which the details and results are presented in
Section 5. In Section 6, we discuss our results in terms of comet
and solar wind characteristics. Lastly, in Section 7 we summa-
rize our findings. Details of the individual Chandra comet ob-
servations are given in Appendix A.
2. Charge Exchange Model
2.1. Atomic structure of He-like ions
Electron capture by highly charged ions populates highly excited
states, which subsequently decay to the ground state. These cas-
cading pathways follow ionic branching ratio statistics. Because
decay schemes work as a funnel, the lowest transitions (n = 2→
1) are the strongest emission lines in cxe spectra. For helium-like
ions, these are the forbidden line (z: 1s2 1S0–1s2s 3S1), the inter-
combination lines (y, x: 1s2 1S0–1s2p 3P1,2), and the resonance
line (w: 1s2 1S0–1s2p 1P1), see Figure 1.
The apparent branching ratio, Beff , for the intercombination
transitions is determined by weighting branching ratios (B j) de-
rived from theoretical transition rates compiled by Porquet et al.
Fig. 1. Part of the decay scheme of a helium–like ion. The 1S0 decays to
the ground state via two-photon processes (not indicated).
Table 1. Apparent effective branching ratios (Beff) for the relaxation of
the 23P-state of He-like carbon, nitrogen, oxygen and neon.
transition C v N vi O vii Ne ix
1s2 (1S0)–1s2p (3P1,2) 0.11 0.22 0.30 0.34
1s2s (3S1)–1s2p (3P0,1,2) 0.89 0.78 0.70 0.66
(2000, 2001), by an assumed statistical population of the triplet
P-term:
Beff =
(2 j + 1)
(2L + 1)(2S + 1)
· B j (1)
The resulting effective branching ratios are given in Table 1.
These ratios can only be observed at conditions where the
metastable state is not destroyed (e.g. by UV flux or collisions)
before it decays. In contrast to many other astrophysical X-
ray sources, this condition is fulfilled in cometary atmospheres,
making the forbidden lines strong markers of cxe emission.
2.2. Emission Cross Sections
To obtain line emission cross sections we start with an initial
state population based on state selective electron capture cross
sections and then track the relaxation pathways defined by the
ion’s branching ratios.
Electron capture reactions can be strongly dependent on
target effects. An important difference between reactions with
atomic hydrogen and the other species is the presence of multi-
ple electrons, hence allowing for multiple (mostly double) elec-
tron transfer. It has been demonstrated both experimentally and
theoretically that double electron capture can be an important
reaction channel in multi-electron targets and that after autoion-
ization to an excited state it may contribute to the X-ray emis-
sion Ali et al. (2005); Hoekstra et al. (1989); Beiersdorfer
et al. (2003); Otranto et al (2006); Bodewits et al. (2006).
Unfortunately, experimental data on reactions with species typ-
ical for cometary atmospheres, such as H2O, atomic O and CO
are at best scarcely available. Because the first ionization po-
tentials of these species are all close to that of atomic H, using
state selective one electron capture cross sections for bare ions
charge exchanging with atomic hydrogen from theory is a rea-
sonable assumption, which is also confirmed by experimental
studies Greenwood et al. (2000, 2001); Bodewits et al. (2006).
Here, we will use the working hypothesis that effective one elec-
tron cross sections for multi-electron targets present in cometary
atmospheres are at least roughly comparable to cross sections for
D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey 3
Table 2. Compilation of theoretical, velocity dependent emission cross sections for collisions between bare- and H-like solar wind ions and atomic
hydrogen, in units of 10−16 cm2. See text for details. We estimate uncertainties to be ca. 20%. The ion column contains the resulting ion, not the
original solar wind ion. Line energies compiled from Garcia & Mack (1965); Vainshtein & Safronova (1985); Drake (1988); Savukov et al.
(2003) and the chianti database Dere et al. (1997); Landi et al. (2006).
E (eV) Ion Transition 200 km s−1 400 km s−1 600 km s−1 800 km s−1 1000 km s−1
299.0 C v z 8.7 12 16 18 20
304.4 C v x,y 0.65 1.0 1.5 1.7 1.8
307.9 C v w 1.8 3.0 4.1 4.8 5.2
354.5 C v 1s3p-1s2 0.55 0.71 0.81 1.0 1.3
367.5 C v 1s4p-1s2 0.70 0.66 0.76 0.74 0.72
367.5 C vi 2p-1s 15 26 30 33 34
378.9 C v 1s5p-1s2 0.00 0.02 0.05 0.04 0.04
419.8 N vi z 13 23 28 29 29
426.3 N vi x,y 2.7 4.3 5.3 5.7 6.0
430.7 N vi w 3.8 6.0 7.4 8.1 8.5
435.5 C vi 3p-1s 1.6 4.0 4.7 4.7 4.8
459.4 C vi 4p-1s 2.9 5.9 7.0 6.4 6.0
471.4 C vi 5p-1s 0.55 1.0 1.3 0.85 0.54
497.9 N vi 1s3p-1s2 0.43 0.99 1.3 1.3 1.3
500.3 N vii 2p-1s 40 45 44 42 42
523.0 N vi 1s4p-1s2 0.81 1.6 1.9 1.8 1.7
534.1 N vi 1s5p-1s2 0.14 0.31 0.33 0.21 0.14
561.1 O vii z 37 34 33 32 31
568.6 O vii x,y 10 10 10 9.9 9.7
574.0 O vii w 9.9 11 11 11 10
592.9 N vii 3p-1s 6.3 4.9 4.8 4.5 4.3
625.3 N vii 4p-1s 2.9 2.9 3.7 4.3 4.6
640.4 N vii 5p-1s 11 5.2 3.7 2.7 2.2
650.2 N vii 6p-1s 0.00 0.21 0.13 0.09 0.08
653.5 O viii 2p-1s 27 40 48 51 53
665.6 O vii 1s3p-1s2 1.7 1.3 1.3 1.2 1.2
697.8 O vii 1s4p-1s2 0.81 0.79 1.0 1.2 1.3
712.8 O vii 1s5p-1s2 2.8 1.3 0.92 0.68 0.54
722.7 O vii 1s6p-1s2 0.00 0.06 0.04 0.02 0.02
774.6 O viii 3p-1s 2.6 4.7 5.6 5.3 5.0
817.0 O viii 4p-1s 1.0 1.6 2.0 2.2 2.3
836.5 O viii 5p-1s 2.4 4.0 4.6 4.1 3.7
849.1 O viii 6p-1s 1.6 1.6 1.5 1.1 0.67
one electron capture from H. Based on this hypothesis, we will
use our comet-wind interaction model to evaluate the contribu-
tion of the different species.
For our calculations, we use a compilation of theoretical state
selective, velocity dependent cross sections for collisions with
atomic hydrogen Errea et al. (2004); Fritsch and Lin (1984);
Green et al. (1982); Shipsey et al. (1983). We furthermore
assume that capture by H-like ions leads to a statistical triplet
to singlet ratio of 3:1, based on measurements by Suraud et al.
(1991); Bliek et al. (1998). We will first focus on the strongest
emission features, which are the n = 2 → 1 transitions, i.e.,
the Ly-α transition (H-like ions) or the forbidden, resonance and
intercombination lines (He-like ions).
In Fig. 2, the emission cross sections of the Ly-α or the sum
of the emission cross sections of the forbidden, resonance and
intercombination lines of different ions (C, N, O) are shown as
a function of collision velocity, for one electron capture reac-
tions with atomic hydrogen. This figure sets the stage for solar
wind velocity induced effects in cometary X-ray spectra. Most
important is the effect of the velocity on the two carbon emis-
sion features; their prime emission features increase by a factor
of almost two when going from typical ‘slow’ to typical ‘fast’
solar wind velocities. The O viii Ly-α emission cross section can
be seen to drop steeply below ca. 300 km s−1. The N vi K-α dis-
plays a similar, though somewhat less strong behavior.
Fig. 2. Velocity dependence of Ly-α or the sum of the forbid-
den/resonance/intercombination emission cross sections of different so-
lar wind ions: O viii (dashed, grey line), O vii (solid, black line), N vii
(dotted, black line), N vi (solid, grey line), C vi (dashed, black line) and
C v (dash-dotted, black line).
The relative intensity of the emission lines (per species) is
governed by the state selective electron capture cross sections
4 D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey
Fig. 3. Velocity dependence of the hardness ratio of different solar wind
ions: O viii (solid line), O vii (dashed line) N vii (dashed line) and C vi
(dash-dotted line). Also shown are two experimentally obtained hard-
ness ratios by Beiersdorfer et al. (2001) and Greenwood et al. (2000)
for O8+ colliding on CO2 and H2O, respectively (see text).
of the charge exchange reaction and the branching ratios of the
resulting ion. A measure of these intensities is the hardness ra-
tio (Beiersdorfer et al. 2001), which is defined as the ratio be-
tween the emission cross sections of the higher order terms of
the Lyman-series and Ly-α (or between the higher order K-series
and K-α in case of He-like ions):∑
n>2 σem(Ly−n)
σem(Ly−α)
For electron capture by H-like ions, we will use the ratio be-
tween the sum of the resonance-, intercombination and forbid-
den emission lines and the rest of the K-series as the hardness
ratio. Fig. 3 shows the hardness ratios of cxe from abundant so-
lar wind ions. The figure shows that most hardness ratios are
constant at typical solar wind velocities (above 300 km s−1) but
it also clearly demonstrates the suggestion made by Beiersdorfer
et al. (2001) that hardness ratios are good candidates for studies
of velocimetry deep within the coma when the solar wind has
slowed down by mass loading.
2.3. Interaction Model
Cometary high-energy emission depends upon certain properties
of both the comet (gas production rate, composition, distance to
the Sun) and the solar wind (speed, composition). Recently, we
developed a model that takes each of these effects into account
Bodewits et al. (2006), which we will briefly describe here.
The neutral gas model is based on the Haser-equation, which
assumes that a comet has a spherically expanding neutral coma
Haser (1957); Festou (1981). The lifetime of neutrals in the
solar radiation field varies greatly amongst species typical for
cometary atmospheres Huebner et al. (1992). The dissociation
and ionization scale lengths also depend on absolute UV fluxes,
and therefore on the distance to the Sun. The coma interacts
with solar wind ions, penetrating from the sunward side follow-
ing straight line trajectories. The charge exchange processes be-
tween solar wind ions and coma neutrals are explicitly followed
both in the change of the ionization state of the solar wind ions
Fig. 4. Modeled charge state distribution along the comet-Sun line, as-
suming an equatorial 300 km s−1 wind interacting with a comet with
outgassing rate Q=1029 molecules s−1 at 1 AU from the Sun. A compo-
sition typical for the slow, equatorial wind was assumed.
and in the relaxation cascade of the excited ions (as discussed
above).
Due to its interaction with the cometary atmosphere, the so-
lar wind is both decelerated and heated in the bow shock. This
bow shock does not affect the ionic charge state distribution. The
bow shock lowers the drift velocity of the wind but at the same
time increases its temperature and the net collision velocity of
the ions is ca. 77% of the initial velocity v(∞) throughout the
interaction zone. We use a rule of thumb derived by Wegmann
et al. (2004) to estimate the stand-off distance Rbs of the bow
shock.
Deep within the coma, the solar wind finally cools down as
the hot wind ions, neutralized by charge exchange, are replaced
by cooler cometary ions. For simplicity however, we shall as-
sume that the wind keeps a constant velocity and temperature
after crossing the bow shock.
Initially, the charge state distribution depends on the solar
wind state. For most simulation purposes, we will assume the
‘average’ ionic composition for the slow, equatorial solar wind
as given by Schwadron and Cravens (2000). Using our compi-
lation of charge changing cross sections, we can solve the differ-
ential equations that describe the charge state distribution in the
coma in the 2D-geometry fixed by the comet-Sun axis. Figure 4
shows the charge state distribution for a 300 km s−1 equato-
rial wind interacting with a comet with an outgassing rate Q of
= 1029 molecules s−1 comet. From this charge state distribution,
it can be seen that along the comet-Sun axis, the comet becomes
collisionally thick between 3500 km (O8+) to 2000 km (C6+),
depending on the cross section of the ions. A maximum in the
C5+ abundance can be seen around 2,000 km, which is due to the
relatively large initial C6+ population and the small cross section
of C5+ charge exchange.
A 3D integration assuming cylindrical symmetry around the
comet-Sun axis finally yields the absolute intensity of the emis-
sion lines. Effects due to the observational geometry (i.e. field of
view and phase angle) are included at this step in the model.
D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey 5
Fig. 5. Relative contribution of target species to the total intensity of
O vii 570 eV emission complex with increasing field of view, for an
active Q= 1029 molecules s−1 comet, interacting with a 300 km s−1
solar wind at 1 AU from the Sun. The shaded area indicates the range
of apertures used to obtain spectra discussed within this survey.
3. Model Results
3.1. Relative Contribution of Target Species
Figure 5 shows the dominant collisions which underly the X-ray
emission of comets. Shown is the total intensity projected on the
sky, with increasing field of view. Within 104 km around the nu-
cleus, water is the dominant collision partner. Farther outward
(≥ 2 × 105 km), the atomic dissociation products of water take
over, and atomic oxygen becomes the most important collision
partner. When the field of view exceeds 107 km, atomic hydro-
gen becomes the sole collision partner. Note that collisions with
water never account for 100% of the emission, even with very
small apertures, due to the contribution of collisions with atomic
hydrogen, OH and oxygen in the line of sight towards the nu-
cleus.
The comets observed with Chandra are all observed with an
aperture of ca. 7.5′ centered on the nucleus. This corresponds to
a range of 1.6−22×104 km (as indicated in Figure 5). Our model
predicts that the emission from nearby comets will be dominated
by cxe from water, but that for comets observed with a larger
field of view, up to 60% of the emission can come from cxe
interactions with the water dissociation products atomic oxygen
and OH, and 10% from interactions with atomic hydrogen.
3.2. Solar Wind Velocity
To illustrate solar wind velocity induced variations in charge ex-
change spectra, we simulated charge exchange spectra follow-
ing solar wind interactions between an equatorial wind and a
Q = 1029 molecules s−1 comet, and assumed the same solar
wind composition in all cases. In Fig. 6, spectra resulting from
collisional velocities of 300 km s−1 and 700 km s−1 are shown.
In the spectrum from the faster wind, the C vi 367 eV and O vii
570 eV emission features are roughly equally strong, whereas at
300 km s−1, the oxygen feature is clearly stronger. Assuming the
wind’s composition remains the same, within the range of typ-
ical solar wind velocities (300–700 km s−1), the cross sectional
dependence on solar wind velocity does not affect cometary X-
ray spectra by more than a factor 1.5. In practice, the composi-
tional differences between slow and fast wind will induce much
stronger spectral changes.
Fig. 6. Simulated X-ray spectra for a 1029 molecules s−1 comet interact-
ing with an equatorial wind with velocities of 300 km s−1 (solid grey
line) and 700 km s−1 (dashed black line). The spectra are convolved
with Gaussians with a width of σ = 50 eV to simulate the Chandra
spectral resolution. To indicate the different lines, also the 700 km s−1
σ = 1 eV spectrum is indicated (not to scale). A field of view of 105 km
and ‘typical’ slow wind composition were used.
3.3. Collisional Opacity
Many of the 20+ comets that have been observed in X-ray dis-
play a typical crescent shape as the solar wind ion content
is depleted via charge exchange. Comets with low outgassing
rates around 1028 molecules s−1, such as 2P/2002 (Encke) and
9P/2005 (Tempel 1), did not display this emission morphology
Lisse et al. (2005, 2007). Whether or not the crescent shape
can be resolved depends mainly on properties of the comet (out-
gassing rate), but, to a minor extent, also on the solar wind
(velocity dependence of cross sections). Other parameters (sec-
ondary, but important), are the spatial resolution of the instru-
ment and the distance of the comet to the observer.
In a collisionally thin environment, the ratio between emis-
sion features is the product of the ion abundance ratios and the
ratio between the relevant emission cross sections:
rthin =
n(Aq+)
n(Bq+)
em (v)
em (v)
The flux ratio for a collisionally thick system depends on the
charge states considered. In case of a bare ion A and a hydro-
genic ion B, the ratio between the photon fluxes from A and B
is given by the abundance ratio weighted by efficiency factors µ
and η:
rthick =
n(Aq+)
n(B(r−1)+) + µ(Br+)n(Br+)
η(Aq+)
η(B(r−1)+)
The efficiency factor µ is a measure of how much B(r−1)+ is pro-
duced by charge exchange reactions by Bq+:
σr,r−1(v)
σr(v)
where σr is the total charge exchange cross section and σr,r−1
the one electron charge changing cross section. The efficiency
factor η describes the emission yield per reaction and is given
by the ratio between the relevant emission cross section σem and
the total charge changing cross section σr:
σem(v)
σr(v)
6 D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey
Fig. 7. Collisional opacity effects on flux ratios within the field of view.
The outer bounds of the fields of view within this survey were between
104 − 105 km, as indicated by the shaded area. We considered a 500 km
s−1 equatorial wind interacting with comets with different activities:
Q = 1028 molecules s−1 (dashed lines) and Q = 1029 molecules s−1
(solid lines). All flux ratios are normalized to 1 at infinity.
To explore the effect of collisional opacity on spectra, we
simulated two comets at 1 AU from the Sun, with gas pro-
duction rates of 1028 and 1029 molecules s−1, interacting with
a solar wind with a velocity of 500 km s−1 and an averaged
slow wind composition Schwadron and Cravens (2000). The
results are summarized in Figure 7 where different flux ratios
are shown. The behavior of these ratios as a function of aper-
ture is important because they can be used to derive relative
ionic abundances. All ratios are normalized to 1 at infinite dis-
Fig. 8. Simulated X-ray spectra for a 1029 molecules s−1 comet inter-
acting with an equatorial wind with a velocity of 300 km s−1 for fields
of view decreasing from 105 km (solid line), 104 km (dashed line) and
103 km (dotted line).
tance from the comet’s nucleus. For low activity comets with
Q ≤ 1028 molecules s−1, the collisional opacity does not affect
the comet’s X-ray spectrum. Within typical field of views all line
flux ratios are close to the collisionally thin value. For more ac-
tive comets (Q = 1029 molecules s−1), collisional opacity can
become important within the field of view. Observed flux ratios
involving C v should be treated with care, see e.g. C v/O vii and
C vi/C v, because the flux ratios within the field of view can be
affected by almost 50% and 35%, respectively. The effect is the
strongest in these cases because of the large relative abundance
of C6+, that contributes to the C v emission via sequential elec-
tron capture reactions in the collisionally thick zones. For N vii
and O viii, a small field of view of 104 km could affect the ob-
served ionic ratios by some 20%.
To further illustrate these results, we show the result-
ing X-ray spectra in Fig. 8. There, we consider a Q =
1029 molecules s−1 comet interacting with a 300 km s−1 wind
and show the effect of slowly zooming from the collisionally
thin to the collisionally thick zone around the nucleus. The field
of view decreases from 105 to 103 km. At 105 km, the spectrum
is not affected by collisionally thick emission, whereas the emis-
sion within an aperture of 1000 km is almost purely from the
interactions within the collisionally thick zones of the comet,
which can be most clearly seen by the strong enhancement of
the C v emission around 300 eV.
The results of our model efforts demonstrate that cometary
X-ray spectra reflect characteristics of the comet, the solar wind
and the observational conditions. Firstly, charge exchange cross
sections depend on the velocity of the solar wind, but its effects
are the strongest at velocities below regular solar wind velocities.
Secondly, collisional opacity can affect cometary X-ray spectra
but mainly when an active comet (Q = 1029 molecules s−1) is
observed with a small field of view (≤ 5×104 km). The dominant
factor however to explain differences in cometary CXE spectra
is therefore the state and hence composition of the solar wind.
This implies that the spectral analysis of cometary X-ray spectra
can be used as a direct, remote quantitative and qualitative probe
of the solar wind.
D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey 7
8 D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey
Fig. 9. Chandra comet observations during the descending phase of so-
lar cycle # 23. Monthly sunspot numbers (grey line) and smoothed
monthly sunspot number (black lines) from the Solar Influences Data
Analysis Center of the Department of Solar Physics, Royal Observatory
of Belgium (http://sidc.oma.be/). Letters refer to the chronological or-
der of observation.
4. Observations
In this section, we will briefly introduce the different comet ob-
servations performed with Chandra. A summary of comet and
solar wind parameters is given in Table 3. More observational
details on the comet and a summary of the state of the solar wind
at the location of the comet during the X-ray observations can be
found in Appendix A.
4.1. Solar Wind Data
Our survey spans the whole period between solar maximum (mid
2000) and solar minimum (mid 2006), see Fig. 9. During solar
minimum, the solar wind can be classified in polar- and equa-
torial streams, where the polar can be found at latitudes larger
than 30◦ and the equatorial wind within 15◦ of the helioequator.
Polar streams are fast (ca. 700 km s−1) and show only small vari-
ations in time, in contrast to the irregular equatorial wind. Cold,
fast wind is also ejected from coronal holes around the equa-
tor, and when these streams interact with the slower background
wind corotating interaction regions (cirs) are formed. As was
illustrated by Schwadron and Cravens (2000), different wind
types vary greatly in their compositions, with the cooler, fast
wind consisting of on average lower charged ions than the hot-
ter equatorial wind. This clear distinction disappears during solar
maximum, when at all latitudes the equatorial type of wind dom-
inates. In addition, coronal mass ejections are far more common
around solar maximum.
There is a strong variability of heavy ion densities due to
variations in the solar source regions and dynamic changes in
the solar wind itself Zurbuchen & Richardson (2006). The vari-
ations mainly concern the charge state of the wind as elemental
variations are only on the order of a factor of 2 (Von Steiger et
al. (2000), and references therein).
We obtained solar wind data from the online data archives
of ace (proton velocities and densities from the swepam instru-
ment, heavy ion fluxes from the swics and swims instruments1)
and soho (proton fluxes from the Proton Monitor Instrument2).
Both ace and soho are located near Earth, at its Lagrangian
1 http://www.srl.caltech.edu/ace/ASC/level2/index.html
2 http://umtof.umd.edu/pm/crn/
point L1. In order to map the solar wind from L1 to the posi-
tion of the comets, we used the time shift procedure described
by Neugebauer et al. (2000). The calculations are based on
the comet ephemeris, the location of L1 and the measured wind
speed. With this procedure, the time delay between an element
of the corotating solar wind arriving at L1 and the comet can be
predicted. A disadvantage of this procedure is that it cannot ac-
count for latitudinal structures in the wind or the magnetohydro-
dynamical behavior of the wind (i.e., the propagation of shocks
and cmes). These shortcomings imply that especially for comets
that have large longitudinal, latitudinal and/or radial separations
from Earth, the solar wind data is at best an estimate of the local
wind conditions. The resulting proton velocities at the comets
near the time of the Chandra observations are shown in Fig. 10.
Parallel to this helioradial and heliolongitudinal mapping, we
compared our comet survey to a 3D mhd time–dependent so-
lar wind model that was employed during most of Solar Cycle
23 (1997 - 2006) on a continuous basis when significant solar
flares were observed. The model (reported by Fry et al. (2003);
McKenna-Lawlor et al. (2006) and Z.K. Smith, private com-
munication, for, respectively, the ascending, maximum, and de-
scending phases) treats solar flare observations and maps the
progress of interplanetary shocks and cmes. The papers men-
tioned above provide an rms error for ”hits” of ±11 hours Smith
et al. (2000); McKenna-Lawlor et al. (2006). cir fast forward
shocks were also taken into account in order to differentiate be-
tween the co-rotating ”quiet” and transient structures. It was im-
portant, in this differentiating analysis, to examine (as we have
done here) the ecliptic plane plots of both of these structures as
simulated by the deforming interplanetary magnetic field lines
(see, for example, Lisse et al. (2005, 2007) for several of the
comets discussed here.) Therefore, the various comet locations
(Table 3) were used to estimate the probability of their X-ray
emission during the observations being influenced by either of
these heliospheric situations.
4.2. X-ray Observations
After its launch in 1999, 8 comets have been observed with
the Chandra X-ray Observatory and Advanced ccd Imaging
Spectrometer (acis). Here, we have mainly considered obser-
vations made with the acis-S3 chip, which has the most sensi-
tive low energy response and for which the majority of comets
were centered. The Chandra’s acis-S instrument provides mod-
erate energy resolution (σ ≈ 50 eV) in the 300 to 1500 eV en-
ergy range, the primary range for the relatively soft cometary
emission. All comets in our sample were re-mapped into comet-
centered coordinates using the standard Chandra Interactive
Analysis of Observations (ciao v3.4) software ‘sso freeze’ al-
gorithm.
Comet source spectra were extracted from the S3 chip with
a circular aperture with a diameter of 7.5′, centered on the
cometary emission. The exception was comet C/2001 Q4, which
filled the chip and a 50% larger aperture was used. acis’ re-
sponse matrices were used to model the instrument’s effective
area and energy dependent sensitivity matrices were created for
each comet separately using the standard ciao tools.
Due to the large extent of cometary X-ray emission, and
Chandra’s relatively narrow field of view, it is not trivial to ob-
tain a background uncontaminated by the comet and sufficiently
close in time and viewing direction. We extracted background
spectra using several techniques: spectra from the S3 chip in
an outer region generally > 8′, an available acis S3 blank sky
observation, and backgrounds extracted from the S1 ccd. For
D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey 9
Fig. 10. Solar wind proton velocities estimated from ace and soho data. For all comets, the time of the observations is indicated with a dotted line.
Letters refer to the chronological order of observation.
several comets there are still a significant number of cometary
counts in the outer region of the S3 ccd. Background spec-
tra taken from the S1 chip have the advantage of having been
taken simultaneous with the S3 observation and thus having
the same space environment as the S3 observation. In general
the background spectra were extracted with the same 7.5′ aper-
ture as the source spectra but centered on the S1 chip. For
comet Encke, where the S1 chip was off during the observa-
tion the background from the outer region of the S3 chip was
used. Comet C/2000 WM1 (linear) was observed with the Low-
Energy Transmission Grating (letg) and acis-S array. For the
latter, we analyzed the zero-th order spectrum, and used a back-
ground extracted from the outer region of the S3 chip. It is possi-
ble that the proportion of incident X-rays diffracted onto the S3
chip will vary with photon energy. Background-subtracted spec-
tra generally have a signal-to-noise at 561 eV of at least 10, and
over 50 for 153P/2002 C1 (Ikeya–Zhang).
5. Spectroscopy
The observed spectra are shown in Figure 11. The spectra
suggest a classification based upon three competing emission
features, i.e. the combined carbon and nitrogen emission (be-
low 500 eV), O vii emission around 565 eV and O viii emis-
sion at 654 eV. Firstly, the C+N emission (<500 eV) seems to
be anti-correlated with the oxygen emission. This clearly sets
the spectra of 73P/2006 S.–W.3B and 2P/2003 (Encke) apart, as
for those two comets the C+N features are roughly as strong as
the O vii emission. In the spectra of the remaining five comets,
oxygen emission dominates over the carbon and nitrogen emis-
sion below 500 eV. The O viii/O vii ratio can be seen to increase
continuously, culminating in the spectrum of 153P/2002 (Ikeya–
Zhang) where the spectrum is completely dominated by oxygen
emission with almost comparable O viii and O vii emission fea-
tures. From our modelling, we expect that the separate classes
reflect different states of the solar wind, which imply different
ionic abundances. To explore the obtained spectra more quanti-
tatively, we will use a spectral fitting technique based on our cxe
model to extract X-ray line fluxes.
5.1. Spectral Fitting
The charge exchange mechanism implies that cometary X-ray
spectra result from a set of solar wind ions, which produce at
least 35 emission lines in the regime visible with Chandra. As
comets are extended sources, these lines cannot all be resolved.
All spectra were therefore fit using the 6 groups of fixed lines of
our cxe model (see Table 2) and spectral parameters were de-
rived using the least squares fitting procedure with the xspec
package. The relative strengths from all lines were fixed per
ionic species, according to their velocity dependent emission
cross sections. Thus, the free parameters were the relative fluxes
of the C, N and O ions contained in our model.
Two additional Ne lines at 907 eV (Ne ix) and 1024 eV
(Ne x) were also included, giving a total of 8 free parameters.
All line widths were fixed at the acis-S3 instrument resolution.
The spectra were fit in the 300 to 1000 eV range. This pro-
vided 49 spectral bins, and thus 41 degrees of freedom. acis
spectra below 300 eV are discarded because of the rising back-
ground contributions, calibration problems and a decreased ef-
fective area near the instrument’s carbon edge.
As a more detailed example of the cxe model and compari-
son to the data, we show in Fig 12 the acis-S3 data for C/1999
S4 (linear). The figure shows the background subtracted source
spectrum over-plotted with the background spectrum, the differ-
ence between the model and data, and the model spectrum and
10 D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey
Fig. 11. Observed spectrum and fit of all 8 comets observed with
Chandra, grouped by their spectral shape (see text). The histogram lines
indicate the CXE model fit.
data to indicate to contribution of the different ions. Only the
emission lines with >3% strength of the strongest line in their
species are shown for ease of presentation.
The fluxes obtained by our fitting are converted into relative
ionic abundances by weighting them by their velocity dependent
emission cross sections. For comets observed near the ecliptic
plane (< 15◦), solar wind conditions mapped to the comet were
used (Section 4.1). For comets observed at higher latitudes, these
data are most likely not applicable and a solar wind velocity of
500 km s−1 was assumed.
Fig. 12. Details of the cxe fit for the spectrum of comet 1999/S4
(linear). Top panel: Comet (filled triangles) and background (open
squares) spectrum. Middle panel: Residuals of cxe fit Bottom panel:
cxemodel and observed spectrum indicating the different lines and their
strengths. Carbon - red; nitrogen - orange; oxygen - blue; neon - green.
The unfolded model is scaled above the emission lines for the ease of
presentation.
Fig. 13. Parameter sensitivity for the major emission features in the fit
of C/1999 S4 (linear), with respect to the O vii 561 eV feature. All units
are 10−4 photons cm−2 s−1. The contours indicate a χ2R of 9.2 (or 99%
confidence, largest, green contour), a χ2R of 4.6 (90%, red contour) and
a χ2R of 2.3 (68%, smallest, blue contour).
5.2. Spectroscopic Results
The fits to all cometary spectra are shown in Fig. 11 and the
results of the fits are given in Table 4. For the majority of the
comets, the model is a good fit to the data within a 95% con-
fidence limit (χ2R ≈ 1.4). Results for comet 153P/2002 (Ikeya–
Zhang) are presented in Table 5 with an additional systematic
error to account for its brightness and any uncertainties in the
response.
The spectra for all comets are well reproduced in the 300 to
1000 eV range. The nitrogen contribution is statistically signifi-
cant for all comets except the fainter ones, 2P/2003 (Encke) and
D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey 11
12 D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey
Table 6. Solar wind abundance relative to O7+, obtained for comet linear S4. References: Bei ’03 – Beiersdorfer et al. (2003), Kra ’04 –
Krasnopolsky (2004), Kra ’06 – Krasnopolsky (2006), Otr ’06 – Otranto et al (2006) and S&C ’00 – Schwadron and Cravens (2000). Dots
indicate that an ion was included in the fitting, but no abundances were derived; dash means that an ion was not included in the fitting. Otranto et
al (2006) did not fit the observed spectrum, but used a combination of ace-data and solar wind averages from Schwadron and Cravens (2000) to
compute a syntectic spectrum of the comet. Solar wind averages are given for comparison Schwadron and Cravens (2000)
Ion this work Bei 03 Kra 04 Kra 06 Otr 06
O8+ 0.32 ± 0.03 0.13 ± 0.03 0.13 ± 0.05 0.15 ± 0.03 0.35 0.35
C6+ 1.4 ± 0.4 0.9 ± 0.3 0.7 ± 0.2 0.7 ± 0.2 1.02 1.59
C5+ 12 ± 4.0 11 ± 9 . . . 1.7 ± 0.7 1.05 1.05
N7+ 0.07 ± 0.06 0.06 ± 0.02 – – 0.03 0.03
N6+ 0.63 ± 0.21 0.5 ± 0.3 – – 0.29 0.29
Ne10+ 0.02 ± 0.01 – – – – –
Ne9+ . . . – (15 ± 6) × 10−3 (20 ± 7) × 10−3 – –
73P/2006 (S.-W.3B). For example, removing the nitrogen com-
ponents from linear S4’s cxe model and re-fitting, increases χ2R
to over 7.
χ2 contours for C/1999 S4 (linear) are presented in Fig 13.
The line strengths for each ionic species are generally well con-
strained, except where spectral features overlap. This can be
readily seen when comparing the contours for the N vii 500 eV
and O vii 561 eV features where a strong anti-correlation exists
(Figure 12). Due to the limited resolution of acis an increase in
the N vii feature will decrease the O vii strength. Similar anti-
correlations exist between the nitrogen N vi or N vii and C v
299 eV lines. Since the line strength for the main line in each
ionic species is linked to weaker lines, a range of energies can
contribute and better constrain its strength. However with O vii
as the strongest spectral feature the nitrogen and carbon compo-
nents may be artificially lower as a result of the aforementioned
anti-correlations. The lack of effective area due to the carbon
edge in the acis response also may over-estimate the C v line
flux. The neon features were well constrained for the brighter
comets, but this is a region of lower signal and some caution
must be taken when treating the neon line strengths and they are
included here largely for completeness.
In the case of 153P/Ikeya–Zhang, the χ2R > 1.4. The main
discrepancy is that the model produces not enough flux in the
700 to 850 eV range compared to the observed spectrum. This
may reflect an underestimation of higher O viii transitions or the
presence of species not (yet) included in the model, such as Fe.
This will be discussed further in the last section of this paper and
in a separate paper dedicated to the observations of this comet
(K. Dennerl, private communication).
One of the best studied comets is C/1999 S4 (linear), be-
cause of its good signal-to-noise ratio. To discuss our results, we
will compare our findings with earlier studies of this comet. In
general, the spectra analyzed here have more counts than ear-
lier analyzes, because of improvements in the Chandra process-
ing software and because we took special care to use a back-
ground that is as comet-free as possible. Previous studies appear
to have removed true comet signal when the background subtrac-
tion was performed. In particular, both the Krasnopolsky (2004)
and Lisse et al. (2001) studies used background regions from
the outer part of the S3 chip and this may have still had true
cometary emission. Krasnopolsky (2004) subtracted over 70%
of the total signal as background. We find that using the S1-chip,
the background contributes only 20% of the total counts.
Different attempts to derive relative ionic abundances from
C1999/S4’s X-ray spectrum are compared in Table 6. Our atomic
physics based spectral analysis combines the benefits of ear-
lier analytical approaches by Kharchenko and Dalgarno (2000,
2001); Beiersdorfer et al. (2003). These methods were all ap-
plied to just one or two comets. Beiersdorfer et al. (2003) inter-
pret C1999/S4’s X-ray spectrum by fitting it with 6 experimental
spectra obtained with their ebit setup. The resulting abundances
are very similar to ours. The advantage of their method is that
it includes multiple electron capture, but in order to observe the
forbidden line emission, the spectra were obtained with trapped
ions colliding at CO2, at collision energies of 200 to 300 eV or
ca. 30 km s−1. As was shown in Fig. 3, the cxe hardness ratio
may change rapidly below 300 km s−1, implying an overesti-
mation of the higher order lines compared to the n = 2 → 1
transition, which for O vii overlap with the O viii emission. We
therefore find higher abundances of O8+.
Krasnopolsky (2004, 2006) obtained fluxes and ionic abun-
dances by fitting the spectrum with 10 lines of which the energies
were semi-free. Their analysis thus does not take the contamina-
tion of unresolved emission into account, and N vi and N vii are
not included in the fit. The line energies were attributed to cxe
lines of mainly solar wind C and O but also to ions of Mg and Ne.
The inclusion of the resulting low energy emission (near 300 eV)
results in lower C5+ fluxes (see also Otranto et al (2006)).
There are several factors that may contribute to the unexpect-
edly low C vi/C v ratios: 1) There may be a small contribution to
the C v line from other ions in the 250-300 eV range (e.g. Si,
Mg, Ne) that are currently not included in the model. Including
these species in the model would lower the C v flux, but proba-
bly only with a small amount. 2) The low acis effective area in
the 250-300 eV region allows the C v flux to be unconstrained,
and this increases the uncertainty in the C v flux. We estimate
that the uncertainty in the effective area, introduced by the car-
bon edge, can account for an uncertainty as large as a factor of
10 in the observed C v/C vi ratios.
We will not compare our results with measured ace/swics
ionic data. As discussed in section 4, the solar wind is highly
variable in time and its composition can change dramatically
over the course of less than a day. Variations in the solar wind’s
ionic composition are often more than 50% during the course of
an observation. Data on N, Ne, and O8+ ions have not been well
documented as the errors of these abundances are dominated by
counting statistics. As discussed above, latitudinal and corota-
tional separations imply large inaccuracies in any solar wind
mapping procedure. These conditions clearly disfavor modelling
based on either average solar wind data or ace/swics data.
6. Comparative Results
As noted in Section 5, spectral differences show up in the be-
havior of the low energy C+N emission (< 500 eV), the O vii
D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey 13
Table 7. Correlation between classification according to spectral shape and comet/solar wind characteristics during the observations. Comet
families from Marsden & Williams (2005). Phase refers to where in the solar cycle the comet was observed, where 1 is the solar maximum and 0
the solar minimum of cycle #23’s descending phase. For other references, see Table 3.
Class # Comet Comet Q Latitude Wind Type
Family (1028 mol. s−1)
cold H 73P/2006 (S.-W.3B) Jupiter 2 0.5 cir
E 2P/2003 (Encke) Jupiter 0.7 11.4 Flare/PS
warm F C/2001 Q4 (neat) unknown 10 -3 Quiet
G 9P/2005 (Tempel 1) Jupiter 0.9 0.8 Quiet
hot C C/2000 WM1 (linear) unknown 3-9 -34 PS
A C/1999 S4 (linear) unknown 3 24 icme
B C/1999 T1 (McNaught–Hartley) unknown 6-20 15 Flare/cir
D C/2002 C1 (Ikeya–Zhang) Oort 20 26 icme
Fig. 14. Flux ratios of all observed comets. The low energy C+N feature
is anti-correlated to the oxygen ionic ratio. Letters refer to the chrono-
logical order of observation.
Fig. 15. Ion ratios of all observed comets. The C+N ionic abundantie is
anti-correlated to the oxygen ionic ratio. Letters refer to the chronolog-
ical order of observation.
emission at 561 eV and the O viii emission at 653 eV. Figure 14
shows a color plot of the fluxes of these three emission features,
and Figure 15 the corresponding ionic abundances. There is a
clear separation between the two comets with a large C+N con-
tribution and the other ‘oxygen-dominated’ comets, which on
their turn show a gradual increase in the oxygen ionic ratio. This
sample of comet observations suggest that we can distinguish
two or three spectral classes.
Table 7 surveys the comet parameters for the different spec-
tral classes. The outgassing rate, heliocentric- or geocentric dis-
tance and comet family do not correlate to the different classes,
in accordance with our model findings. The data does suggest
a correlation between latitude and wind conditions during the
observations. At first sight, the apparent correlation between lat-
itude and oxygen ratio seems paradoxical. According to the bi-
modal structure of the solar wind the fast, cold wind dominates
at latitudes > 15◦, implying less O viii emission. In Figure 9,
the comet observations are shown with respect to the phase of
the last solar cycle. Interestingly, we note that all comets that
were observed at higher latitudes were observed around solar
maximum. The solar wind is highly chaotic during solar maxi-
mum and the frequency of impulsive events like CMEs is much
higher than during the descending and minimum phase of the
cycle. This explains both why the comets observed in the period
2000–2002 encountered a disturbed solar wind and why our sur-
vey does not contain a sample of the cool fast wind from polar
coronal holes.
The observed classification can therefore be fully ascribed
to solar wind states. The first class is associated with cold, fast
winds with lower average ionization. These winds are found in
cirs and behind flare related shocks. The spectra due to these
winds are dominated by the low energy x-rays, because of the
low abundances of highly charged oxygen. At the relevant tem-
peratures, most of the solar wind oxygen is He-like O6+, which
does not produce any emission visible in the 300–1000 eV
regime accessible with Chandra. Secondly, there is an interme-
diate class with two comets that were all observed during periods
of quiet solar wind. These comets interacted with the equatorial,
warm slow wind. The third class then comprises comets that in-
teracted with a fast, hot, disturbed wind associated with icmes
or flares. From the solar wind data, Ikeya–Zhang was proba-
bly the most extreme example of this case. This comet had 10
times more signal than any other comet in our sample and small
discrepancies in the response may be important at this level.
Extending into the 1-2 keV regime, a preliminary analysis indi-
cates the presence of bare and H-like Si, Mg and Fe xv-xx ions, in
accordance with ace measurements of icme compositions Lepri
& Zurbuchen (2004).
The variability and complex nature of the solar wind allows
for many intermediate states in between these three categories
Zurbuchen et al. (2002), which explain the gradual increase
of the O viii/O vii ratio that we observed in the cometary spec-
tra. As the solar wind is a collisionless plasma, the charge state
14 D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey
Fig. 16. Spectrum derived ionic oxygen ratios and corresponding
freezing-in temperatures from Mazotta et al. (1998). The shaded area
indicates the typical range of slow wind associated with streamers.
Letters refer to the chronological order of observation.
distribution in the solar wind is linked to the temperature in its
source region. Ionic temperatures are therefore a good indicator
of the state of the wind encountered by a comet. The ratio be-
tween O7+ and O6+ ionic abundances has been demonstrated to
be a good probe of solar wind states. Zurbuchen et al. (2002)
observed that slow, warm wind associated with streamers typi-
cally lies within 0.1 < O7+/O6+ < 1.0, corresponding to freez-
ing in temperatures of 1.3–2.1 MK. The corresponding temper-
ature range is indicated in the Figure 16. In the figure, we show
the observed O8+ to O7+ ratios and the corresponding freezing-
in temperatures from the ionizational/recombination equilibrium
model by Mazotta et al. (1998). Most observations are within or
near to the streamer-associated range of oxygen freezing in tem-
peratures. Four comets interacted with a wind significantly hot-
ter than typical streamer winds, and in all four cases we found
evidence in solar wind archives that the comets most likely en-
countered a disturbed wind.
7. Conclusions
Cometary X-ray emission arises from collisions between bare-
and H-like ions (such as C, N, O, Ne, . . . ) with mainly water and
its dissociation products OH, O and H. The manifold of depen-
dencies of the cxe mechanism on characteristics of both comet
and wind offers many diagnostic opportunities, which are ex-
plored in the first part of this paper. Charge exchange cross sec-
tions are strongly dependent on the velocity of the solar wind,
and these effects are strongest at velocities below the regular
wind conditions. This dependency might be used as a remote
plasma diagnostics in future observations. Ruling out collisional
opacity effects, we used our model to demonstrate that the spec-
tral shape of cometary cxe emission is in the first place deter-
mined by local solar wind conditions. Cometary X-ray spectra
hence reflect the state of the solar wind.
Based on atomic physic modelling of cometary charge ex-
change emission, we developed an analytical method to study
cometary X-ray spectra. First, the data of 8 comets observed
with Chandra were carefully reprocessed to avoid the subtrac-
tion of cometary signal as background. The spectra were then
fit using an extensive data set of velocity dependent emission
cross sections for eight different solar wind species. Although
the limited observational resolution currently available hampers
the interpretation of cometary X-ray spectra to some degree, our
spectral analysis allows for the unravelling of cometary X-ray
spectra and allowed us to derive relative solar wind abundances
from the spectra.
Because the solar wind is a collisionless plasma, local ionic
charge states reflect conditions of its source regions. Comparing
the fluxes of the C+N emission below 500 eV, the O vii emission
and the O viii emission yields a quantitative probe of the state
of the wind. In accordance with our modelling, we found that
spectral differences amongst the comets in our survey could be
very well understood in terms of solar wind conditions. We are
able to distinguish interactions with three different wind types,
being the cold, fast wind (I), the warm, slow wind (II); and the
hot, fast, disturbed wind (III). Based on our findings, we pre-
dict the existence of even cooler cometary X-ray spectra when a
comet interacts with the fast, cool high latitude wind from polar
coronal holes. The upcoming solar minimum offers the perfect
opportunity for such an observation.
Acknowledgements. DB and RH acknowledge support within the framework of
the fom–euratom association agreement and by the Netherlands Organization
for Scientific Research (nwo). MD thanks the noaa Space Environment Center
for its post-retirement hospitality. We are grateful for the cometary ephemerides
of D. K. Yeomans published at the jpl/horizons website. Proton velocities used
here are courtesy of the soho/celias/pm team. soho is a mission of international
cooperation between esa and nasa. Chianti is a collaborative project involving
the nrl (USA), ral (UK), mssl (UK), the Universities of Florence (Italy) and
Cambridge (UK), and George Mason University (USA).
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Appendix A: Observations within this Survey
This Appendix presents the observational details of the Chandra
data and the corresponding solar wind state. The prefix ’FF’
(fearless forecast) used in this appendix refers to the real time
forecasting of coronal mass ejection shocks arrivals at Earth. The
numbers were so-named for flare/coronal shock events during
solar cycle #23.
A.1. C/1999 S4 (linear)
X-rays. The first Chandra cometary observation was of comet
C/1999 S4 (linear) Lisse et al. (2001), with observations being
made both before and after the breakup of the nucleus. Due to the
low signal-to-noise ratio of the second detection, only the July
14th 2000 pre-breakup observation is discussed here. Summing
the 8 pointings of the satellite gave a total time interval of 9390 s.
In this period, the acis-S3 ccd collected a total of 11 710 photons
were detected in the range 300–1000 eV. Detections out side this
range or on other acis-ccds were not attributed to the comet. As
a result, data from the S1-ccd (which is configured identically to
S3) may be used as an indicator of the local X-ray background.
The morphology can be described by a crescent shape, with
the maximum brightness point 24 000 km from the nucleus
on the Sun-facing side. The brightness dims to 10% of the
maximum level at 110 000 km from the nucleus.
Solar wind. A large velocity jump can be seen around DoY 199,
which was due to the famous ”Bastille Day” flare on 14 July
(FF#153, Dryer et al (2001); Fry et al. (2003)). This flare
reached the comet only after the first observation. At July 12,
2017UT a solar flare started at N17W65 (FF#152), which was
nicely placed to hit this comet with a very high probability dur-
ing the first observations Fry et al. (2003). As for the second
observation, there was another flare on July 28, S17E24, at 1713
UT (FF#164) and there was a high probability that its shock’s
weaker flank hit the comet.
A.2. C/1999 T1 (McNaught–Hartley)
X-rays. The allocated observing time of comet McNaught–
Hartley was partitioned into 5 one-hour-slots between January
8th and January 15th, 2001 Krasnopolsky et al. (2002). The
strongest observing period was on January 8th, when ∆ = 1.37
AU and rh = 1.26 AU.
There were 15 000 total counts observed by the acis-S3 ccd
between 300 and 1000 eV. The emission region can be described
by a crescent, with the peak brightness is at 29 000 km from
the nucleus. The brightness dims to 10% of the maximum at a
cometocentric distance of 260 000 km. Again, the acis-S1 ccd
may be used to indicate the local background signal.
Solar wind. The comet was not within the heliospheric cur-
rent/plasma sheet (HCS/HPS). Two corotating cirs are probably
associated with the first two observations. Two flares (FF#233
and #234) took place; however, another corotating cir more
likely arrived before the flare’s transient shock’s effects did
McKenna-Lawlor et al. (2006).
A.3. C/2000 WM1 (linear)
X-rays. The only attempt to use the high-resolution grat-
ing capability of the acis-S array was made with comet
C/2000 WM1 (linear). Here, the Low-Energy Transmission
Grating (letg) was used. The dimness of the observed X-rays,
and the extended nature of the emitting atmosphere meant that
the grated spectra did not yield significant results. It is still
16 D. Bodewits et al.: Spectral Analysis of the Chandra Comet Survey
possible to extract a spectrum based on the pulse-heights gener-
ated by each X-ray detection on the acis-S3 chip, although the
morphology is not recorded. 6300 total counts were recorded for
the pulse-height spectrum of the S3 chip in the 300 to 1000 eV
range.
Solar wind. Comet WM1 was observed at the highest latitude
available within this survey, and at a latitude of 34 degrees, it was
far outside the hcs. During the observations, this comet might
have experienced the southerly flank of the shock of a strong
X3.4 flare at S20E97 and its icme and shock on December 28,
2001 (FF#359) McKenna-Lawlor et al. (2006).
A.4. 153P/2002 (Ikeya–Zhang)
X-rays. The brightest X-ray comet in the Chandra archive is
153P/2002 (Ikeya–Zhang). The heliographic latitude, geocentric
distance and heliocentric distance were comparable to those for
comet C/1999 S4 (linear), with a latitude of 26◦, ∆ = 0.457 AU
and rh = 0.8 AU. Rather than periodically re-point the detector
to track the comet, the pointing direction was fixed and the
comet was monitored as it passed through the field of view, thus
increasing the effective FoV. There were two observing periods
on April 15th 2002, each lasting for approximately 3 hours and
15 minutes. In both periods, a strong cometary signal is detected
on all of the activated acis-ccds. Consequently, a background
signal cannot be taken from the observation. A crescent shape
on the Sun side of the comet is observed over all of the ccd
array. Over 200 000 total counts were observed from the S3
chip in the 300 to 1000 eV range. The time intervals for each
observing period are 11 570 and 11 813 seconds.
Solar wind. Like C/2000 WM1, this comet was observed at a
relatively high heliographic latitude. Solar wind data obtained
in the ecliptic plane can therefore not be used to determine the
wind state at the comet. 153P/2002 (Ikeya–Zhang) was well-
positioned during the first observation on 15 April 2002 for a
flare at N16E05 (FF#388) on 12 April 2002. During the second
observation on 16 April, there was an earlier flare on 14 April at
N14W57, but this flare was probably too far to the west to be ef-
fective McKenna-Lawlor et al. (2006). The comet was observed
at a high latitude, and hence ace solar wind data is most likely
not applicable.
A.5. 2P/2003 (Encke)
X-rays. The Chandra observation of Encke took place on the
24th of November 2003 Lisse et al. (2005), when the comet
had a heliocentric distance of rh = 0.891 AU and a geocentric
distance of ∆ = 0.275 AU and a heliographic latitude of 11.4
degrees. The comet was continuously tracked for over 15 hours,
resulting in a useful exposure of 44 000 seconds. The acis-S3 ccd
counted 6140 X-rays in the range 300–1000 eV.
The brightest point was offset from the nucleus by
11 000 km, dimming to 10% of this value at a distance of
60 000 km.
The acis-S1 ccd was not activated in this observation. The
low quantum efficiency of the other activated ccds below 0.5 keV
makes them unsuitable as background references.
Solar wind. The proton velocity decreased during observations
from 600 km s−1 to 500 km s−1. A flare on 20 November
2003, at N01W08 (FF#525), was well-positioned to affect the
observations on 23 November (data from work in progress by
Z.K. Smith et al.). The comet most likely interacted with the
overexpanded, rarified plasma flow that followed the earlier hot
shocked and compressed flow behind the flare’s shock.
A.6. C/2001 Q4 (neat)
X-rays. A short observation of comet C/2001 Q4 was made on
May 12 2004, when the geocentric and heliocentric distances
were ∆ = 0.362 AU and rh = 0.964 AU respectively. With a
heliographic latitude of 3 degrees, the comet was almost in
the ecliptic plane. From 3 pointings, the useful exposure was
10 328 seconds. The acis-S3 chip detected 6540 X-rays in be-
tween 300 and 1000 eV. The acis-S1 was used as a background
signal.
Solar wind. There was no significant solar activity during the
observations (Z.K. Smith et al., ibid.). From solar wind data, the
comet interacted with a quiet, slow 352 km s−1 wind.
A.7. 9P/2005 (Tempel 1)
X-rays. The observation of comet 9P/2005 (Tempel 1) was de-
signed to coincide with the Deep Impact mission Lisse et al.
(2007). The allocated observation time of 291.6 ks was split
into 7 periods, starting on June 30th, July 4th (encompassing the
Deep Impact collision), July 5th, July 8th, July 10th, July 13th
and July 24th. The brightest observing periods were June 30th
and July 8th. The focus here is on the June 30th observation. On
this date, rh = 1.507 AU and ∆ = 0.872 AU.
The useful exposure was 50 059 seconds, with a total of 7300
counts, 4000 from the June 30th flare alone, were detected in the
energy range of 300–1000 eV.
The brightest point for the June 30th observation was
located 11 000 km from the nucleus. The morphology appears
to be more spherical than in other comet observations.
Solar wind. Observations were taken over a long time span cov-
ering different solar wind environments. There was no signifi-
cant solar activity during the 30 June 2005 observations (Z.K.
Smith et al., ibid. Lisse et al. (2007)). From the ace data, it can
be seen that at June 30, the comet most likely interacted with a
quiet, slow solar wind.
A.8. 73P/2006 (Schwassmann–Wachmann 3B)
X-rays. The close approach of comet 73P/2006 (Schwassmann–
Wachmann 3B) in May 2005 (∆ = 0.106 AU, rh = 0.965 AU)
provided an opportunity to examine cometary X-rays in high
spatial resolution. Chandra was one of several X-ray missions to
focus on one of the large fragments of the comet. Between 300
and 1000 eV, 6285 counts were obtained in a useful exposure of
20 600 seconds.
Solar wind. There was a weak flare on 22 May 2006 (FF#655,
Z.K. Smith, priv. comm.). A sequence of three high speed coro-
nal hole streams passed the comet in the period around the ob-
servations and a corotating cir might have reached the comet in
association with the observations on 23 May, which is confirmed
by the mapped solar wind data.
