Draft version October 29, 2018
Preprint typeset using LATEX style emulateapj v. 10/09/06
WHY DO SOME INTERMEDIATE POLARS SHOW SOFT X-RAY EMISSION? A SURVEY OF XMM-Newton
SPECTRA.
P.A. Evans
and Coel Hellier
Astrophysics Group, Keele University, Staffordshire, ST5 5BG, UK
Draft version October 29, 2018
ABSTRACT
We make a systematic analysis of the XMM-Newton X-ray spectra of intermediate polars (IPs) and
find that, contrary to the traditional picture, most show a soft blackbody component. We compare the
results with those from AM Her stars and deduce that the blackbody emission arises from reprocessing
of hard X-rays, rather than from the blobby accretion sometimes seen in AM Hers. Whether an IP
shows a blackbody component appears to depend primarily on geometric factors: a blackbody is not
seen in those that have accretion footprints that are always obscured by accretion curtains or are only
visible when foreshortened on the white-dwarf limb. Thus we argue against previous suggestions that
the blackbody emission characterises a separate sub-group of IPs which are more akin to AM Hers,
and develop a unified picture of the blackbody emission in these stars.
Subject headings: accretion, accretion discs – novae, cataclysmic variables – X-rays: binaries.
1. INTRODUCTION
Intermediate polars (IPs) – interacting binaries with a
magnetic white-dwarf primary – have traditionally been
noted for their hard X-ray emission. This arises as the
magnetic field of the white dwarf disrupts the accre-
tion disc and channels material towards the magnetic
polecaps. This material forms stand-off shocks, below
which it cools via free-free interactions, producing hard
X-rays. However, a growing number of systems have
been shown to emit a distinct blackbody component
in softer X-rays (e.g. Mason et al. (1992); Haberl et al.
(1994); de Martino et al. (2004)), reminiscent of the soft
component prominent in the X-ray spectra of many AM
Her stars (also known as polars). These systems are
similar to IPs but the white-dwarf has a magnetic field
strong enough to prevent an accretion disc from forming
at all. In these systems, the soft blackbody component
is though to arise from a heated polecap surrounding the
accretion column (See Warner (1985); Hellier (2001) for
a review of these objects).
Currently it is unclear why the blackbody component
is seen in some IPs and not others. Haberl & Motch
(1995) suggested that there are two distinct classes of IP,
with the ‘soft’ systems being evolutionary progenitors of
polars. They argued that the ‘hard IPs’ may have larger
and cooler polecaps, pushing the soft emission into the
EUV and explaining the difference in the spectra.
We present here a study of XMM-Newton X-ray data
of 12 IPs, aimed at discovering why some IPs show a
blackbody component while others don’t. Our method
is similar to that of Ramsay & Cropper (2004) (hereafter
RC04) who analysed the XMM-Newton data of twenty-
one polars, which enables us to compare the IPs with the
polars.
2. OBSERVATIONS AND DATA ANALYSIS
The XMM-Newton observatory (Jansen et al. 2001)
was launched in 1999, and we have obtained observa-
1 Current address: Department of Physics and Astronomy, Uni-
versity of Leicester, Leicester, LE1 7RH, UK
TABLE 1
The XMM observations of IPs analysed in this
paper.
Object ObsID Date References
AO Psc 0009650101 2001-06-09 1,2
EX Hya 0111020101 2000-07-01 1,2
0111020201 2000-07-01 1,2
FO Aqr 0009650201 2001-05-12 1,2,3
GK Per 0154550101 2002-03-09 2,4
HT Cam 0144840101 2003-03-24 2,5,6
PQ Gem 0109510301 2002-10-08 2,7
NY Lup 0105460301 2000-09-07 2,8
UU Col 0201290201 2004-08-21 9
V1223 Sgr 0145050101 2003-04-13 1,2
V405 Aur 0111180401 2001-10-05 2,10
V2400 Oph 0105460601 2001-08-30 2
WX Pyx 0149160201 2003-05-20 11
References. — (1) Cropper et al. (2002), (2)
Evans & Hellier (2005b), (3) Evans et al. (2004), (4)
Vrielmann et al. (2005) (5) de Martino et al. (2005),
(6) Evans & Hellier (2005a), (7) Evans, Hellier & Ram-
say (2006), (8) Haberl et al. (2002), (9) de Martino et
al. (2006) (10) Evans & Hellier (2004), (11) Schlegel
(2005).
tions of twelve IPs from the public archive. We ana-
lyzed the data from the EPIC-MOS and pn instruments
(Turner et al. 2001; Strüder et al. 2001) which provide
high-throughput, medium-resolution spectroscopy across
the 0.2–12 keV energy range. The higher resolution RGS
instruments (den Herder et al. 2001) have only 20 per
cent of the effective area of the MOS cameras and the
data are not used here.
A summary of the observations used is given in Table 1.
We re-ran the pipeline processing for these observations
using xmm-sas v7.0.0. The observations of GK Per, NY
Lup and V2400 Oph suffered from pile-up, and thus only
the wings of the PSF were included in the source extrac-
tion. The MOS-1 observation of EX Hya was so badly
piled up that we excluded it from our analysis.
RC04 used only the EPIC-pn data as it was better cal-
ibrated than the EPIC-MOS data at soft energies. Us-
ing the better calibrations of xmm-sas v.7 we extracted
http://arxiv.org/abs/0704.1388v1
2 Evans & Hellier
spectra from all three EPIC instruments. Response ma-
trices were created for each spectrum, using the xmm-sas
rmfgen and arfgen tasks. We then modelled the spectra
using xspec v11. For each star, all model parameters
were tied between the EPIC instruments, except for the
normalisation which we allowed to vary in order to com-
bat the effects of cross-calibration uncertainties.
Although IP spectra can vary considerably over the
spin cycle, for the majority of the systems in this paper,
we do not have enough geometric information to identify
phase regions when the hard/soft components are best
presented to us (as RC04 did), so we extracted spectra
covering the entire observation. Note that the results of
our spectroscopy are thus weighted averages from across
the spin cycle; this was taken into account when inter-
preting our results .
To reproduce the hard component we used the strat-
ified accretion column model of Cropper et al. (1999).
This models the spectrum in terms of the white dwarf
mass (MWD) and specific accretion rate (i.e. accretion
rate per unit area, ṁ), from which it calculates the
temperature and density profile of the column. This is
then divided into 100 bins, evenly distributed in veloc-
ity space, each bin emitting as an optically thin plasma
(a mekal). To the stratified column model, we added
narrow Gaussians for the 6.4-keV iron fluorescence line
and the 0.547 keV Oxygen vii photoionisation line where
necessary. We then applied to this emission a simple pho-
toelectric absorber. For most systems this did not give
an acceptable fit, so we added either one or two partial-
covering absorbers as necessary.
Next, we added a blackbody component to the mod-
els. Since absorption at the densities of the partial-
covering components (typically ∼ 1023 cm−2) will com-
pletely smother any soft X-ray emission and thus be re-
dundant with model normalisation, the blackbody com-
ponent was absorbed only by the simple absorption,
which was of order 1019–1021 cm−2.
For some systems the addition of a blackbody did not
improve the fit. For these systems we manually raised the
blackbody normalisation until it significantly reduced the
fit quality, thus finding an upper limit. Since this will
be temperature dependent, we did this for blackbody
temperatures of 40, 60 and 80 eV.
We quote, in Table 2, the f-test statistic to judge
the the significance of adding the blackbody component.
However, this test will produce false positives in the pres-
ence of calibration systematics. We have thus estimated
the systematics by fitting a model optimised for the MOS
data to the pn data (allowing only the normalisation to
change) and recording the change in χ2(=∆χ2system). We
claim the presence of a blackbody only if it improves the
χ2 by more than ∆χ2system. This method is more con-
servative than using the f-test alone. We include this
estimate of the systematics in all the errors quoted in
this paper.
Details of the fits are given in Table 3. The ṁ was un-
constrained for every system, so is not given. We do not
quote errors on the partial-covering absorbers as they do
not affect the softness ratio. The ratio is sensitive, how-
ever, to the metal abundance in the column, as there is a
forest of iron L lines in the 0.5–1.2 keV range, affecting
the model fit at the soft end.
For all twelve systems we then calculated the flux from
the hard and soft components. Following RC04 we de-
fined the softness ratio as Fs/4Fh, where Fs and Fh are
the fluxes of the soft and hard components respectively.
The factor of four arises because the hard component is
optically thin and thus radiates isotropically, whereas the
hard component is optically thick. Where the blackbody-
emitting region is seen foreshortened, the observed ratio
will be an underestimate.
The softness ratios are shown in Figs. 1 and 2. We show
the observed ratio, the ratio of unabsorbed fluxes over
the 0.2–12 keV range, and the ratio of unabsorbed fluxes
calculated over all energies. These bolometric fluxes and
softness ratios are given in Table 4. For the systems
with no detectable soft component we show the upper
limit calculated for a 60-eV blackbody, and present the
fluxes and ratios for a range of blackbody temperatures
in Table ??.
3. RESULTS
We show the spectra for the systems with a blackbody
component in Fig. 3, and for those without in Fig. 4. For
the latter we have also shown the upper limit determined
for a 60-eV blackbody component.
For FO Aqr, AO Psc, V1223 Sgr and HT Cam
we found no evidence for a soft component, in agree-
ment with previous observations, (see Norton et al.
(1992); Hellier et al. (1996); Beardmore et al (2000);
Evans & Hellier (2005a) respectively).
3.1. V405 Aur
The XMM observation of V405 Aur contains system-
atic discrepancies between the two EPIC-MOS instru-
ments below 0.4 keV. However, when processed under
sas 7.0 these are at a much lower level than when
Evans & Hellier (2004) analysed the data, and we have
made no allowance for these discrepancies in the fit. Note
also that as there is no pn data for V405 Aur, we have
no estimate of the effects of systematics discussed in Sec-
tion 2, so our errors are likely to be underestimates.
The best-fitting blackbody temperature was kT =
64.78+0.81
−1.11 eV. This is significantly higher than the 40±4
eV reported by Evans & Hellier (2004) analysing the
same observation, however they used two mekals to
fit the hard component whereas we used the stratified
column model. Since the calibration has also changed
since Evans & Hellier (2004), we analysed our better-
calibrated data using their model, and found a fit in
agreement with theirs. This demonstrates that the re-
sults are somewhat model dependent; the stratified col-
umn model is likely to be the more physically realis-
tic. Fitting the hard component with a single, high
temperature plasma, Haberl et al. (1994) found a black-
body temperature of 49–64 eV (from ROSAT data) and
de Martino et al. (2004) found 73 ± 14 eV (using Bep-
poSAX ).
Our fitted hydrogen column of 3.46+0.41
−0.31 × 10
20 cm−2
agrees with that of de Martino et al. (2004) [(4±2)×1020
cm−2] but not with those of Haberl et al. (1994) or
Evans & Hellier (2004) who reported (5.7 ± 0.3) × 1020
and (10.6+0.9
−1.2) × 10
20 cm−2 respectively. However, the
fitted column will depend on the emission model used,
so some discrepancy is expected.
Soft X-rays from IPs 3
TABLE 2
Fit statistics for each star with and without a blackbody.
Star χ2 (dof) χ2 (dof) (with bb) f-test ∆χ2system χ
(No bb) (with bb)
AO Psc 3372.74 (2888) 3772.71 (2884) 0.93 700 0.03
FO Aqr 1462 (1864) 1444 (1860) 1.6×10−4 168 18
HT Cam 1325 (1273) 1310 (1269) 4.1×10−3 79 16
V1223 Sgr 3840.8 (3128) 3840.6 (3124) 0.99 143 0.2
EX Hya 14045 (4876) 10158 (4873) < 10−99 1284 3887
GK Per 17040 (4079) 5024 (4075) < 10−99 84 12016
NY Lup 902 (699) 669 (695) 8×10−14 40 233
PQ Gem 20435 (2439) 2940 (2435) < 10−99 144 17495
UU Col 1676 (817) 910 (813) < 10−99 76 766
V2400 Oph 1019 (1003) 953 (999) 9×10−14 50 66.15
V405 Aur 17626 (997) 1146 (994) < 10−99 * 16480
WX Pyx 594 (477) 495 (473) 8×10−18 52 99
Note. — The f-test gives the probability that no blackbody is present, making no
allowance for systematics. The ∆χ2system is the change in χ
2 in fitting the same model
to the MOS and pn cameras, thus giving an estimate of the systematic errors. The last
column is the improvement in χ2 when a blackbody is added. We consider this significant if
it exceeds ∆χ2system.
∗ There was no pn data for V405 Aur, so ∆χ2system was not estimated.
TABLE 3
Spectral components used in the fitted models.
Star wabs nH blackbody kT Part Abs (1) Part Abs (2) MWD Abundance
(1020cm−2) (eV) (nH , Cv Frc) (nH , Cv Frc) (M⊙) (solar)
V405 Aur 3.46 (+0.41, −0.31) 64.78 (+0.81, −1.11 ) 17, 0.49 3.0, 0.63 0.40 (+0.05, −0.06) 0.069 (+0.024, −0.021)
GK Per 23.3 (+2.0, −1.9) 62 (±2) 23, 0.74 4.7, 0.45 0.92 (+0.39, −0.13) 0.21 (+0.14, −0.07)
NY Lup 7.8 (±3.9) 104 (+21, −23) 14, 0.49 0.38, 0.71 0.96 (+0.40, −0.55) 0.68 (+0.51, −0.59)
V2400 Oph 7.0 (+2.9, −4.9) 117 (+33, −44) 11, 0.52 0.61, 0.53 0.69 (+0.06, −0.24) 0.33 (+0.12, −0.10)
PQ Gem 0 (+0.30) 47.6 (+2.9, −1.4) 42, 0.60 3.4, 0.56 0.70 (+0.16, −0.14) < 0.08
EX Hya 9.76 (+2.2, −0.86) 31.0 (+1.3, -2.4) 75, 0.35 4.0, 0.29 0.449 (+0.005, −0.013) 0.514 (+0.01, −0.0029)
UU Col 0 (+0.59) 73 (+20, −9) 10, 0.34 1.23 (+0.17, −0.29) 0.66 (+1.0, -0.62)
WX Pyx 8.4 (+3.8, −2.9) 82 (+11, −15) 1.4 (+0, −0.09) < 2.87
FO Aqr 0 (+2.1) 21, 0.80 6.4, 0.98 1.19 (+0.11, −0.31) 0.31 (+0.20, −0.23)
AO Psc 3.89 (+0.69, −1.44) 14, 0.62 1.8, 0.75 0.594 (+0.13, −0.040) 0.362 (+0.20, −0.064)
HT Cam 3.86 (+0.81, -0.88) 0.687 (+0.094, −0.061) 0.52 (+0.24, −0.11)
V1223 Sgr 1.03 (+0.36, −0.52) 13, 0.46 1.3, 0.63 1.046 (+0.049, −0.012) 0.398 (+0.090, -0.049)
Note. — The column density of the partial absorption is given in units of 1022 cm−2. Errors are quoted to the same power of ten as the corresponding
parameter.
TABLE 4
The unabsorbed, bolometric fluxes of the soft and hard components for those systems which
show blackbody emission.
Object Fh,bol Fs,bol Ratio
(erg s−1 cm−2) (erg s−1 cm−2)
V405 Aur 5.1× 10−11 (+3.6, −1.1) 4.3× 10−11 (+2.4, −1.2) 0.211 (+0.018, -0.038)
GK Per 1.20× 10−9 (+0.25, −0.06) 2.29× 10−10 (+0.95, −0.62) 4.8× 10−2 (+1.8, −1.3)
NY Lup 4.15× 10−11 (+12.7, −0.18) 4.3× 10−12 (+9.9, −1.4) 2.6× 10−2 (+1.6, −1.1)
V2400 Oph 9.2× 10−11 (+4.1, −1.9) 3.3× 10−12 (+2.1, −1.5) 8.9× 10−3 (+4.6, −2.8)
PQ Gem 1.07× 10−10 (+0.40, −0.23) 1.33× 10−11 (+0.13, −0.17) 3.11× 10−2 (+0.74, −0.83)
EX Hya 3.95× 10−10 (+0.26, −0.20) 1.59× 10−10 (+3.47, −0.62) 1.00 (+0.48, −0.20)
WX Pyx 7.51× 10−12 (+0.35, −0.93) 6.0× 10−13 (+5.9, −2.9) 2.00× 10−2 (+2.38, −0.93)
UU Col 6.81× 10−12 (+2.29, −0.80) 3.04× 10−13 (±0.98) 1.12× 10−2 (±0.45)
Note. — The ratio is defined as in Fig. 1. Errors are given to the same power of ten as the values.
4 Evans & Hellier
Fig. 1.— Softness ratios of the IPs observed with XMM, defined as Fs/4Fh, where Fs and Fh are the fluxes of the soft blackbody and
hard plasma components respectively, calculated over the 0.2–12 keV energy range covered by XMM. Upper panel : Ratios calculated from
the spectral fits. Lower panel : The ratios calculated from the spectral fits, after the absorption components were removed.
Fig. 2.— As for Fig. 1, but with the effects of absorption removed and Fs and Fh extended over all energies. The ratios for the polars
given in RC04 are also shown (hollow squares); RC04 did not quote errors. The uppermost six systems are the polars in which RC04 found
no blackbody component. We have found an upper limit for these systems as we did for the IPs. The dashed line corresponds to a softness
ratio of 0.5; systems with a higher softness ratio exhibit a ‘soft excess’.
Soft X-rays from IPs 5
TABLE 5
The unabsorbed, bolometric fluxes from the systems with no detectable soft
X-ray component, and the upper limit of the softness ratio, for a range of
temperatures.
Object Fh,bol Ratio40eV Ratio60eV Ratio80eV
(erg s−1 cm−2)
FO Aqr 2.71× 10−10(+0.65, −0.18) < 4.6× 10−4 < 1.4× 10−4 < 8.0× 10−5
AO Psc 1.51× 10−10(+0.09, −0.11) < 4.3× 10−3 < 1.1× 10−3 < 6.2× 10−4
HT Cam 8.48× 10−12(+0.53, −0.38) < 2.5× 10−2 < 5.3× 10−3 < 2.6× 10−3
V1223 Sgr 2.96× 10−10( ±0.13) < 8.5× 10−4 < 3.5× 10−4 < 2.3× 10−4
3.2. GK Per
A soft blackbody component was necessary to model
the XMM spectrum of GK Per as previously found by
Vrielmann et al. (2005). They reported a blackbody
temperature of 59.6 ± 0.2 eV absorbed by a column of
(3.2±0.2)×1021 cm−2. Our temperature of 62±2 eV and
column of (2.3± 0.2) cm−2 are very similar, though not
formally in agreement. Note that Vrielmann et al. (2005)
parameterised the hard emission using a bremsstrahlung
component and a mekal, supporting our assertion above
that these results are model dependent.
3.3. NY Lup
Haberl et al. (2002) analysed this XMM observation
of NY Lup (=RXJ154814) and found a soft component
with a blackbody temperature of 84–97 eV and a column
density of (11.7–15.5)×1020 cm−2. Our values of kTbb =
104+21
−23 eV and nH = (7.8± 3.9)× 10
21 cm−2 agree.
3.4. V2400 Oph
V2400 Oph was identified as a soft IP by
de Martino et al. (2004), who analysed a BeppoSAX ob-
servation and reported a blackbody temperature of 103±
10 eV and absorption column (46+12
−13)×10
20 cm−2. We
find a blackbody temperature of 117+33
−44 eV, in agreement
with this result, but a slightly lower absorption column of
(7.0+2.9
−4.9)× 10
20 cm−2. This is probably because de Mar-
tino used a single mekal and a single partial-covering
absorber to model the hard emission, whereas we used
the stratified column model and two partial covering ab-
sorbers.
3.5. PQ Gem
PQ Gem was the first IP found to have a soft-X-ray
component (Mason et al. 1992). This component was
also present in the XMM data, with a best-fitting black-
body temperature of 47.6+2.9
−1.4 eV, in agreement with the
46+12
−23 eV of Duck et al. (1994) from ROSAT data, and
56+12
−14 eV of de Martino et al. (2004) from BeppoSAX
data. The fitted column density goes to zero, which is
likely to be an artefact of fitting a complex absorption
with too simple a model. We quote an upper limit of
3× 1019 cm−2 based on the phase-resolved modelling of
Evans et al. (2006).
3.6. EX Hya
The best-fitting model for EX Hya used a blackbody
component, which has not been previously reported in
this system. However even with this component, the
procedure outlined in Section 2 resulted in a poor fit
=2.1). A possible reason for this is our choice of ab-
sorption model. We have used a cold absorber in our
models since the data do not warrant the extra param-
eters in ionised absorption models, even though one ex-
pects any absorbing material (e.g. the accretion curtains)
to be ionised. We therefore tried various ionised absorp-
tion models, but gained only a minor improvement to
the fit. We thus reverted to the cold absorber model for
consistency with the rest of this paper. We also tried us-
ing phase-resolved spectroscopy, in case the poor fit was
the result of averaging phase-variant parameters, how-
ever this still did not yield an acceptable fit. We have
nonethless included our results for EX Hya, for complete-
ness, but due to the poor fit, we do not much place much
weight on the EX Hya data when considering our results.
As the distance to EX Hya is known (64.5±1.2 pc:
Beuermann et al. (2003)), we can determine the size of
the accretion footprint from the soft X-ray flux. Table 4
gives this as (1.59+3.47
−0.62) × 10
−10 ergs cm−2 s−1 with a
temperature of 31.0+1.3
−2.4 eV, from which we compute an
emitting area (8.4+29.8
−4.2 ) × 10
13 cm2. Suleimanov et al
(2005) gave the mass of the white dwarf in EX Hya as
0.5±0.05M⊙, thus the observed blackbody emitting area
in EX Hya covers (7.3+29.3
−4.0 ) × 10
−4 of the white-dwarf
surface.
3.7. UU Col
UU Col was identified as a soft IP by Burwitz et al.
(1996). de Martino et al. (2006) have recently confirmed
this with a detailed analysis of the XMM observation.
They reported a blackbody temperature of 49.7+5.6
−2.9 eV,
which is lower than our value of 73+20
−9 eV, however in
their model the blackbody is absorbed by the partial cov-
ering absorber, and no simple absorber is present.
3.8. WX Pyx
The XMM observation of WX Pyx, the only X-ray
observation of this star to date, has a relatively low sta-
tistical quality. It was previously analysed by Schlegel
(2005) who did not report looking for a blackbody com-
ponent. However, we find that adding a blackbody does
significantly improve the fit.
3.9. Comparison with the polars
In Fig. 2 we have plotted the softness ratios of both the
IPs and the polars (from RC04). For the polars which
RC04 reported not to have a blackbody, we obtained the
spectra as extracted an calibrated by RC04 (Ramsay,
private communication), and fitted them in the same way
as the IPs (Section 2) to obtain an upper limit.
6 Evans & Hellier
Fig. 3.— The EPIC-pn spectra of the eight IPs for which the best-fitting models contain a blackbody component. The solid line shows
the hard component; the broken line the blackbody. For V405 Aur we have shown the MOS-1 spectrum, since the pn camera did not collect
any data.
The chief difference in the two distributions is that
while several polars show a softness ratio > 0.5, no IP
can be confirmed to do this, and it can be excluded for
all but EX Hya – for which our results are uncertain
(Section 3.6). The ‘soft excess’ in polars is believed to
arise due to ‘blobby accretion’ (e.g. Kuijpers & Pringle
(1982)). In this model, dense blobs of matter penetrate
into the white dwarf photosphere and the energy is ther-
malised to a blackbody.
Whether such accretion occurs in IPs has not been
widely discussed in the literature. Hellier & Beardmore
(2002) suggested that viscous interactions in an accre-
tion disc would destroy blobs, although Vrielmann et al.
(2005) interpreted flares in the lightcurve of GK Per as
resulting from the accretion of blobs. Our findings sug-
gest that blobby accretion is not significant in IPs.
4. DISCUSSION
The ‘polar’ class of magnetic cataclysmic variable has
long been known to be characterised by a soft black-
body component (e.g. King & Watson (1987)). This
Soft X-rays from IPs 7
Fig. 4.— The EPIC-pn spectra of the four IPs for which the best-fitting model does not contain a blackbody component. The solid line
shows the best-fitting model. The broken line shows the upper limit to a blackbody component, given a temperature of 60 eV. For FO Aqr
we show the MOS-1 data, as the signal-to-noise ratio of the pn data is worse.
is considered to arise from the white-dwarf surface,
heated either by reprocessing of hard X-rays from
the accretion column, or by thermalisation of blobs
of accretion (e.g. Kuijpers & Pringle (1982)). In con-
trast, IPs were thought to lack this component (e.g.
King & Lasota (1990)). However, observations with
ROSAT found a blackbody component in some IPs, lead-
ing Haberl & Motch (1995) to suggest that there were
two spectrally distinct classes of IP. This raised the ques-
tion of why.
To address this we have conducted a systematic survey
of the spectral characteristics of the IPs observed with
XMM-Newton, which has much greater spectral coverage
and throughput than ROSAT.
We find that, of twelve IPs analysed, eight show a soft
blackbody component while four do not. This suggests
that a blackbody is a normal component of IPs, and
hence of accretion onto magnetic white dwarfs, and that
the spectra differ only in degree.
We thus ask what causes the differing visibility
of the soft component. There does not appear to
be any correlation with the white-dwarf mass [see
Cropper et al. (1999); Ezuka & Ishida (1999); Ramsay
(2000); Suleimanov et al (2005) for mass estimates], nor
any obvious correlation with the orbital period.
In polars, systems with higher magnetic field strengths
appear to have higher softness ratios (e.g. Ramsay et al.
(1994)). Of the IPs in our sample showing polarisation,
and thus known to have a relatively strong field (5–20
MG), all (PQ Gem, V405 Aur and V2400 Oph) show
a blackbody component, while the four stars showing
no blackbody emission (FO Aqr, AO Psc, HT Cam &
V1223 Sgr) do not show polarisation. We give a possible
explanation for this after discussing the role of absorp-
tion.
We first consider the simple absorber, which is proba-
bly of interstellar origin. The detection of a blackbody
component in most systems shows that interstellar ab-
sorption is not sufficient to extinguish the soft emission.
Further, the systems with no detected blackbody compo-
nent do not have higher interstellar columns than those
with a blackbody (Table 3), so this absorption cannot
explain the differing visibility of the soft component.
We thus turn to the partial-covering absorption, which
in IPs is predominantly caused by the accretion cur-
tains crossing the line of sight. Here we find that
the systems where the lightcurves are dominated by
deep absorption dips owing to the accretion curtains
(FO Aqr, V1223 Sgr and AO Psc; see Beardmore et al.
(1998); Beardmore et al (2000); Hellier et al. (1991) re-
spectively) do tend to be those which lack a blackbody
component. In contrast, systems showing a blackbody
component, such as V405 Aur, NY Lup, EX Hya and
V2400 Oph, tend to be systems where the lightcurves
suggest that the accretion curtains do not hide the accre-
tion footprints (see Evans & Hellier (2004); Haberl et al.
(2002); Allan et al. (1998); Hellier & Beardmore (2002)
respectively).
We thus suggest that the major reason why some IPs
don’t show a blackbody component is simply that the
heated region near the accretion footprint is hidden by
the accretion curtains, while in other IPs it is not, the
difference being the result of the system inclination and
the magnetic colatitude (see Fig. 5). Coupled with this is
8 Evans & Hellier
to foreshortening
No BB seen
BB seen
BB seen
Little BB owing
Fig. 5.— The factors that affect blackbody emission in an IP.
a) When the upper magnetic pole is on the visible face, blackbody
emission will only be seen if the inclination is such that the heated
accretion region is visible above the accretion curtains.
b) When the lower pole is on the visible face, it will likely be too
foreshortened for us to detect blackbody emission.
c) In UU Col the magnetic axis is highly inclined, so the foreshort-
ening seen in b) is reduced and blackbody emission is seen.
the effect of foreshortening, such that an optically thick
heated region will not produce much blackbody emission
if it is only seen while on the white-dwarf limb, rather
than in the middle of the face.
A proper investigation of this idea would need knowl-
edge of the size and location of the accretion footprints
and of the surrounding heated polecaps, so that we could
estimate the difference absorbing columns of different
spectral components, and how these vary with spin-cycle
phase. However, this information is not known for the
majority of IPs. The softness ratio might conceivably
also vary with parameters such as accretion rate and
white-dwarf mass, which are again only poorly known.
However, as a test of our ideas, we can outline how
they might apply to the remaining systems in our sam-
ple which we did not consider when forming the model,
namely HT Cam, GK Per, PQ Gem and UU Col.
In PQ Gem the accretion curtains do cause an ab-
sorption dip when they obscure the accretion footprints.
However, the geometry of this star is relatively well de-
termined (Potter et al. 1997; Mason 1997; Evans et al.
2006) and it appears that the heated polecap is graz-
ingly visible above the accretion curtain for part of the
cycle; thus it shows both an absorption dip and a soft
blackbody, and is on the boundary between the two cases
illustrated in the upper panel of Fig. 5.
UU Col also shows an absorption dip when the ac-
cretion curtains obscure the upper pole, and also shows
blackbody emission. de Martino et al. (2006) proposed
that the blackbody emission comes from the lower pole,
viewed when that pole is closest to us (lowest panel of
Fig. 5). We thus suggest that UU Col has an abnor-
mally high inclination of the magnetic dipole, such that
the lower pole is not foreshortened as much as in other
IPs where no blackbody component is seen. V405 Aur
is another system that appears to have a highly inclined
dipole, such that blackbody emission from the lower pole
is significant, leading in that system to a double-peaked
soft-X-ray lightcurve (Evans & Hellier 2004).
In contrast to all the other IPs, the XMM data of
GK Per reported here were collected during an outburst.
Fig. 6.— Schematic diagram of GK Per in outburst. Accretion
occurs from all azimuths, resulting in a circular blackbody-emitting
region (dark ring). As can be seen, even when the accretion cur-
tains lie across our line-of-sight, part of this region is unobscured.
Hellier et al. (2004) have argued that during outburst the
accretion occurs from all azimuths, forming a complete
accretion ring at the poles. As illustrated in Fig. 6, this
means that some portion of the heated polecap is likely
to be visible ‘behind’ the magnetic pole, where accretion
does not normally occur. Thus in GK Per in outburst we
see a system with both strongly absorbed X-ray emission
(from in front of the magnetic pole) and a blackbody
component.
Lastly, we consider HT Cam. This shows very lit-
tle sign of absorption, and its lightcurve can explained
without any absorption effects (Evans et al. 2006). Yet
it shows no blackbody emission, in apparent contradic-
tion to our model. However, as previously suggested by
de Martino et al. (2006) and Evans et al. (2006), it ap-
pears that HT Cam has an exceptionally low accretion
rate (partly accounting for the lack of absorption). If so,
it could be that the blackbody component is simply too
cool to be detected in the XMM bandpass. We note that
the blackbody temperature in EX Hya, the other star in
our sample below the period gap, is lower than in the
others (Table 3), and that in HT Cam might be lower
still.
5. SUMMARY
We have analysed data from XMM observations of 12
intermediate polars and find that a soft blackbody com-
ponent is a common feature of their X-ray spectra. We
suggest that in the systems showing no blackbody emis-
sion the heated accretion polecaps are largely hidden by
the accretion curtains, or are only visible when on the
white dwarf limb and highly foreshortened. Thus IPs
with lightcurves dominated by absorption dips caused by
the passage of accretion curtains across the line of sight
tend to show no blackbody emission. Further, these are
also the systems least likely to show polarisation, since
the cyclotron-emitting column will also be obscured by
the accretion curtains, or would be beamed away from
us if the accretion region were on the white-dwarf limb.
After comparing the blackbody emission seen in IPs with
that seen in polars, we conclude that the blobby emis-
sion responsible for soft X-ray excesses in polars does not
occur in IPs.
ACKNOWLEDGEMENTS
We thank Gavin Ramsay for providing us with the
spectra of the polars with no detectable soft component.
Facilities: XMM ()
Soft X-rays from IPs 9
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  We make a systematic analysis of the XMM-Newton X-ray spectra of intermediate
polars (IPs) and find that, contrary to the traditional picture, most show a
soft blackbody component. We compare the results with those from AM Her stars
and deduce that the blackbody emission arises from reprocessing of hard X-rays,
rather than from the blobby accretion sometimes seen in AM Hers. Whether an IP
shows a blackbody component appears to depend primarily on geometric factors: a
blackbody is not seen in those that have accretion footprints that are always
obscured by accretion curtains or are only visible when foreshortened on the
white-dwarf limb. Thus we argue against previous suggestions that the blackbody
emission characterises a separate sub-group of IPs which are more akin to AM
Hers, and develop a unified picture of the blackbody emission in these stars.

Draft version October 29, 2018
Preprint typeset using LATEX style emulateapj v. 10/09/06
WHY DO SOME INTERMEDIATE POLARS SHOW SOFT X-RAY EMISSION? A SURVEY OF XMM-Newton
SPECTRA.
P.A. Evans
and Coel Hellier
Astrophysics Group, Keele University, Staffordshire, ST5 5BG, UK
Draft version October 29, 2018
ABSTRACT
We make a systematic analysis of the XMM-Newton X-ray spectra of intermediate polars (IPs) and
find that, contrary to the traditional picture, most show a soft blackbody component. We compare the
results with those from AM Her stars and deduce that the blackbody emission arises from reprocessing
of hard X-rays, rather than from the blobby accretion sometimes seen in AM Hers. Whether an IP
shows a blackbody component appears to depend primarily on geometric factors: a blackbody is not
seen in those that have accretion footprints that are always obscured by accretion curtains or are only
visible when foreshortened on the white-dwarf limb. Thus we argue against previous suggestions that
the blackbody emission characterises a separate sub-group of IPs which are more akin to AM Hers,
and develop a unified picture of the blackbody emission in these stars.
Subject headings: accretion, accretion discs – novae, cataclysmic variables – X-rays: binaries.
1. INTRODUCTION
Intermediate polars (IPs) – interacting binaries with a
magnetic white-dwarf primary – have traditionally been
noted for their hard X-ray emission. This arises as the
magnetic field of the white dwarf disrupts the accre-
tion disc and channels material towards the magnetic
polecaps. This material forms stand-off shocks, below
which it cools via free-free interactions, producing hard
X-rays. However, a growing number of systems have
been shown to emit a distinct blackbody component
in softer X-rays (e.g. Mason et al. (1992); Haberl et al.
(1994); de Martino et al. (2004)), reminiscent of the soft
component prominent in the X-ray spectra of many AM
Her stars (also known as polars). These systems are
similar to IPs but the white-dwarf has a magnetic field
strong enough to prevent an accretion disc from forming
at all. In these systems, the soft blackbody component
is though to arise from a heated polecap surrounding the
accretion column (See Warner (1985); Hellier (2001) for
a review of these objects).
Currently it is unclear why the blackbody component
is seen in some IPs and not others. Haberl & Motch
(1995) suggested that there are two distinct classes of IP,
with the ‘soft’ systems being evolutionary progenitors of
polars. They argued that the ‘hard IPs’ may have larger
and cooler polecaps, pushing the soft emission into the
EUV and explaining the difference in the spectra.
We present here a study of XMM-Newton X-ray data
of 12 IPs, aimed at discovering why some IPs show a
blackbody component while others don’t. Our method
is similar to that of Ramsay & Cropper (2004) (hereafter
RC04) who analysed the XMM-Newton data of twenty-
one polars, which enables us to compare the IPs with the
polars.
2. OBSERVATIONS AND DATA ANALYSIS
The XMM-Newton observatory (Jansen et al. 2001)
was launched in 1999, and we have obtained observa-
1 Current address: Department of Physics and Astronomy, Uni-
versity of Leicester, Leicester, LE1 7RH, UK
TABLE 1
The XMM observations of IPs analysed in this
paper.
Object ObsID Date References
AO Psc 0009650101 2001-06-09 1,2
EX Hya 0111020101 2000-07-01 1,2
0111020201 2000-07-01 1,2
FO Aqr 0009650201 2001-05-12 1,2,3
GK Per 0154550101 2002-03-09 2,4
HT Cam 0144840101 2003-03-24 2,5,6
PQ Gem 0109510301 2002-10-08 2,7
NY Lup 0105460301 2000-09-07 2,8
UU Col 0201290201 2004-08-21 9
V1223 Sgr 0145050101 2003-04-13 1,2
V405 Aur 0111180401 2001-10-05 2,10
V2400 Oph 0105460601 2001-08-30 2
WX Pyx 0149160201 2003-05-20 11
References. — (1) Cropper et al. (2002), (2)
Evans & Hellier (2005b), (3) Evans et al. (2004), (4)
Vrielmann et al. (2005) (5) de Martino et al. (2005),
(6) Evans & Hellier (2005a), (7) Evans, Hellier & Ram-
say (2006), (8) Haberl et al. (2002), (9) de Martino et
al. (2006) (10) Evans & Hellier (2004), (11) Schlegel
(2005).
tions of twelve IPs from the public archive. We ana-
lyzed the data from the EPIC-MOS and pn instruments
(Turner et al. 2001; Strüder et al. 2001) which provide
high-throughput, medium-resolution spectroscopy across
the 0.2–12 keV energy range. The higher resolution RGS
instruments (den Herder et al. 2001) have only 20 per
cent of the effective area of the MOS cameras and the
data are not used here.
A summary of the observations used is given in Table 1.
We re-ran the pipeline processing for these observations
using xmm-sas v7.0.0. The observations of GK Per, NY
Lup and V2400 Oph suffered from pile-up, and thus only
the wings of the PSF were included in the source extrac-
tion. The MOS-1 observation of EX Hya was so badly
piled up that we excluded it from our analysis.
RC04 used only the EPIC-pn data as it was better cal-
ibrated than the EPIC-MOS data at soft energies. Us-
ing the better calibrations of xmm-sas v.7 we extracted
http://arxiv.org/abs/0704.1388v1
2 Evans & Hellier
spectra from all three EPIC instruments. Response ma-
trices were created for each spectrum, using the xmm-sas
rmfgen and arfgen tasks. We then modelled the spectra
using xspec v11. For each star, all model parameters
were tied between the EPIC instruments, except for the
normalisation which we allowed to vary in order to com-
bat the effects of cross-calibration uncertainties.
Although IP spectra can vary considerably over the
spin cycle, for the majority of the systems in this paper,
we do not have enough geometric information to identify
phase regions when the hard/soft components are best
presented to us (as RC04 did), so we extracted spectra
covering the entire observation. Note that the results of
our spectroscopy are thus weighted averages from across
the spin cycle; this was taken into account when inter-
preting our results .
To reproduce the hard component we used the strat-
ified accretion column model of Cropper et al. (1999).
This models the spectrum in terms of the white dwarf
mass (MWD) and specific accretion rate (i.e. accretion
rate per unit area, ṁ), from which it calculates the
temperature and density profile of the column. This is
then divided into 100 bins, evenly distributed in veloc-
ity space, each bin emitting as an optically thin plasma
(a mekal). To the stratified column model, we added
narrow Gaussians for the 6.4-keV iron fluorescence line
and the 0.547 keV Oxygen vii photoionisation line where
necessary. We then applied to this emission a simple pho-
toelectric absorber. For most systems this did not give
an acceptable fit, so we added either one or two partial-
covering absorbers as necessary.
Next, we added a blackbody component to the mod-
els. Since absorption at the densities of the partial-
covering components (typically ∼ 1023 cm−2) will com-
pletely smother any soft X-ray emission and thus be re-
dundant with model normalisation, the blackbody com-
ponent was absorbed only by the simple absorption,
which was of order 1019–1021 cm−2.
For some systems the addition of a blackbody did not
improve the fit. For these systems we manually raised the
blackbody normalisation until it significantly reduced the
fit quality, thus finding an upper limit. Since this will
be temperature dependent, we did this for blackbody
temperatures of 40, 60 and 80 eV.
We quote, in Table 2, the f-test statistic to judge
the the significance of adding the blackbody component.
However, this test will produce false positives in the pres-
ence of calibration systematics. We have thus estimated
the systematics by fitting a model optimised for the MOS
data to the pn data (allowing only the normalisation to
change) and recording the change in χ2(=∆χ2system). We
claim the presence of a blackbody only if it improves the
χ2 by more than ∆χ2system. This method is more con-
servative than using the f-test alone. We include this
estimate of the systematics in all the errors quoted in
this paper.
Details of the fits are given in Table 3. The ṁ was un-
constrained for every system, so is not given. We do not
quote errors on the partial-covering absorbers as they do
not affect the softness ratio. The ratio is sensitive, how-
ever, to the metal abundance in the column, as there is a
forest of iron L lines in the 0.5–1.2 keV range, affecting
the model fit at the soft end.
For all twelve systems we then calculated the flux from
the hard and soft components. Following RC04 we de-
fined the softness ratio as Fs/4Fh, where Fs and Fh are
the fluxes of the soft and hard components respectively.
The factor of four arises because the hard component is
optically thin and thus radiates isotropically, whereas the
hard component is optically thick. Where the blackbody-
emitting region is seen foreshortened, the observed ratio
will be an underestimate.
The softness ratios are shown in Figs. 1 and 2. We show
the observed ratio, the ratio of unabsorbed fluxes over
the 0.2–12 keV range, and the ratio of unabsorbed fluxes
calculated over all energies. These bolometric fluxes and
softness ratios are given in Table 4. For the systems
with no detectable soft component we show the upper
limit calculated for a 60-eV blackbody, and present the
fluxes and ratios for a range of blackbody temperatures
in Table ??.
3. RESULTS
We show the spectra for the systems with a blackbody
component in Fig. 3, and for those without in Fig. 4. For
the latter we have also shown the upper limit determined
for a 60-eV blackbody component.
For FO Aqr, AO Psc, V1223 Sgr and HT Cam
we found no evidence for a soft component, in agree-
ment with previous observations, (see Norton et al.
(1992); Hellier et al. (1996); Beardmore et al (2000);
Evans & Hellier (2005a) respectively).
3.1. V405 Aur
The XMM observation of V405 Aur contains system-
atic discrepancies between the two EPIC-MOS instru-
ments below 0.4 keV. However, when processed under
sas 7.0 these are at a much lower level than when
Evans & Hellier (2004) analysed the data, and we have
made no allowance for these discrepancies in the fit. Note
also that as there is no pn data for V405 Aur, we have
no estimate of the effects of systematics discussed in Sec-
tion 2, so our errors are likely to be underestimates.
The best-fitting blackbody temperature was kT =
64.78+0.81
−1.11 eV. This is significantly higher than the 40±4
eV reported by Evans & Hellier (2004) analysing the
same observation, however they used two mekals to
fit the hard component whereas we used the stratified
column model. Since the calibration has also changed
since Evans & Hellier (2004), we analysed our better-
calibrated data using their model, and found a fit in
agreement with theirs. This demonstrates that the re-
sults are somewhat model dependent; the stratified col-
umn model is likely to be the more physically realis-
tic. Fitting the hard component with a single, high
temperature plasma, Haberl et al. (1994) found a black-
body temperature of 49–64 eV (from ROSAT data) and
de Martino et al. (2004) found 73 ± 14 eV (using Bep-
poSAX ).
Our fitted hydrogen column of 3.46+0.41
−0.31 × 10
20 cm−2
agrees with that of de Martino et al. (2004) [(4±2)×1020
cm−2] but not with those of Haberl et al. (1994) or
Evans & Hellier (2004) who reported (5.7 ± 0.3) × 1020
and (10.6+0.9
−1.2) × 10
20 cm−2 respectively. However, the
fitted column will depend on the emission model used,
so some discrepancy is expected.
Soft X-rays from IPs 3
TABLE 2
Fit statistics for each star with and without a blackbody.
Star χ2 (dof) χ2 (dof) (with bb) f-test ∆χ2system χ
(No bb) (with bb)
AO Psc 3372.74 (2888) 3772.71 (2884) 0.93 700 0.03
FO Aqr 1462 (1864) 1444 (1860) 1.6×10−4 168 18
HT Cam 1325 (1273) 1310 (1269) 4.1×10−3 79 16
V1223 Sgr 3840.8 (3128) 3840.6 (3124) 0.99 143 0.2
EX Hya 14045 (4876) 10158 (4873) < 10−99 1284 3887
GK Per 17040 (4079) 5024 (4075) < 10−99 84 12016
NY Lup 902 (699) 669 (695) 8×10−14 40 233
PQ Gem 20435 (2439) 2940 (2435) < 10−99 144 17495
UU Col 1676 (817) 910 (813) < 10−99 76 766
V2400 Oph 1019 (1003) 953 (999) 9×10−14 50 66.15
V405 Aur 17626 (997) 1146 (994) < 10−99 * 16480
WX Pyx 594 (477) 495 (473) 8×10−18 52 99
Note. — The f-test gives the probability that no blackbody is present, making no
allowance for systematics. The ∆χ2system is the change in χ
2 in fitting the same model
to the MOS and pn cameras, thus giving an estimate of the systematic errors. The last
column is the improvement in χ2 when a blackbody is added. We consider this significant if
it exceeds ∆χ2system.
∗ There was no pn data for V405 Aur, so ∆χ2system was not estimated.
TABLE 3
Spectral components used in the fitted models.
Star wabs nH blackbody kT Part Abs (1) Part Abs (2) MWD Abundance
(1020cm−2) (eV) (nH , Cv Frc) (nH , Cv Frc) (M⊙) (solar)
V405 Aur 3.46 (+0.41, −0.31) 64.78 (+0.81, −1.11 ) 17, 0.49 3.0, 0.63 0.40 (+0.05, −0.06) 0.069 (+0.024, −0.021)
GK Per 23.3 (+2.0, −1.9) 62 (±2) 23, 0.74 4.7, 0.45 0.92 (+0.39, −0.13) 0.21 (+0.14, −0.07)
NY Lup 7.8 (±3.9) 104 (+21, −23) 14, 0.49 0.38, 0.71 0.96 (+0.40, −0.55) 0.68 (+0.51, −0.59)
V2400 Oph 7.0 (+2.9, −4.9) 117 (+33, −44) 11, 0.52 0.61, 0.53 0.69 (+0.06, −0.24) 0.33 (+0.12, −0.10)
PQ Gem 0 (+0.30) 47.6 (+2.9, −1.4) 42, 0.60 3.4, 0.56 0.70 (+0.16, −0.14) < 0.08
EX Hya 9.76 (+2.2, −0.86) 31.0 (+1.3, -2.4) 75, 0.35 4.0, 0.29 0.449 (+0.005, −0.013) 0.514 (+0.01, −0.0029)
UU Col 0 (+0.59) 73 (+20, −9) 10, 0.34 1.23 (+0.17, −0.29) 0.66 (+1.0, -0.62)
WX Pyx 8.4 (+3.8, −2.9) 82 (+11, −15) 1.4 (+0, −0.09) < 2.87
FO Aqr 0 (+2.1) 21, 0.80 6.4, 0.98 1.19 (+0.11, −0.31) 0.31 (+0.20, −0.23)
AO Psc 3.89 (+0.69, −1.44) 14, 0.62 1.8, 0.75 0.594 (+0.13, −0.040) 0.362 (+0.20, −0.064)
HT Cam 3.86 (+0.81, -0.88) 0.687 (+0.094, −0.061) 0.52 (+0.24, −0.11)
V1223 Sgr 1.03 (+0.36, −0.52) 13, 0.46 1.3, 0.63 1.046 (+0.049, −0.012) 0.398 (+0.090, -0.049)
Note. — The column density of the partial absorption is given in units of 1022 cm−2. Errors are quoted to the same power of ten as the corresponding
parameter.
TABLE 4
The unabsorbed, bolometric fluxes of the soft and hard components for those systems which
show blackbody emission.
Object Fh,bol Fs,bol Ratio
(erg s−1 cm−2) (erg s−1 cm−2)
V405 Aur 5.1× 10−11 (+3.6, −1.1) 4.3× 10−11 (+2.4, −1.2) 0.211 (+0.018, -0.038)
GK Per 1.20× 10−9 (+0.25, −0.06) 2.29× 10−10 (+0.95, −0.62) 4.8× 10−2 (+1.8, −1.3)
NY Lup 4.15× 10−11 (+12.7, −0.18) 4.3× 10−12 (+9.9, −1.4) 2.6× 10−2 (+1.6, −1.1)
V2400 Oph 9.2× 10−11 (+4.1, −1.9) 3.3× 10−12 (+2.1, −1.5) 8.9× 10−3 (+4.6, −2.8)
PQ Gem 1.07× 10−10 (+0.40, −0.23) 1.33× 10−11 (+0.13, −0.17) 3.11× 10−2 (+0.74, −0.83)
EX Hya 3.95× 10−10 (+0.26, −0.20) 1.59× 10−10 (+3.47, −0.62) 1.00 (+0.48, −0.20)
WX Pyx 7.51× 10−12 (+0.35, −0.93) 6.0× 10−13 (+5.9, −2.9) 2.00× 10−2 (+2.38, −0.93)
UU Col 6.81× 10−12 (+2.29, −0.80) 3.04× 10−13 (±0.98) 1.12× 10−2 (±0.45)
Note. — The ratio is defined as in Fig. 1. Errors are given to the same power of ten as the values.
4 Evans & Hellier
Fig. 1.— Softness ratios of the IPs observed with XMM, defined as Fs/4Fh, where Fs and Fh are the fluxes of the soft blackbody and
hard plasma components respectively, calculated over the 0.2–12 keV energy range covered by XMM. Upper panel : Ratios calculated from
the spectral fits. Lower panel : The ratios calculated from the spectral fits, after the absorption components were removed.
Fig. 2.— As for Fig. 1, but with the effects of absorption removed and Fs and Fh extended over all energies. The ratios for the polars
given in RC04 are also shown (hollow squares); RC04 did not quote errors. The uppermost six systems are the polars in which RC04 found
no blackbody component. We have found an upper limit for these systems as we did for the IPs. The dashed line corresponds to a softness
ratio of 0.5; systems with a higher softness ratio exhibit a ‘soft excess’.
Soft X-rays from IPs 5
TABLE 5
The unabsorbed, bolometric fluxes from the systems with no detectable soft
X-ray component, and the upper limit of the softness ratio, for a range of
temperatures.
Object Fh,bol Ratio40eV Ratio60eV Ratio80eV
(erg s−1 cm−2)
FO Aqr 2.71× 10−10(+0.65, −0.18) < 4.6× 10−4 < 1.4× 10−4 < 8.0× 10−5
AO Psc 1.51× 10−10(+0.09, −0.11) < 4.3× 10−3 < 1.1× 10−3 < 6.2× 10−4
HT Cam 8.48× 10−12(+0.53, −0.38) < 2.5× 10−2 < 5.3× 10−3 < 2.6× 10−3
V1223 Sgr 2.96× 10−10( ±0.13) < 8.5× 10−4 < 3.5× 10−4 < 2.3× 10−4
3.2. GK Per
A soft blackbody component was necessary to model
the XMM spectrum of GK Per as previously found by
Vrielmann et al. (2005). They reported a blackbody
temperature of 59.6 ± 0.2 eV absorbed by a column of
(3.2±0.2)×1021 cm−2. Our temperature of 62±2 eV and
column of (2.3± 0.2) cm−2 are very similar, though not
formally in agreement. Note that Vrielmann et al. (2005)
parameterised the hard emission using a bremsstrahlung
component and a mekal, supporting our assertion above
that these results are model dependent.
3.3. NY Lup
Haberl et al. (2002) analysed this XMM observation
of NY Lup (=RXJ154814) and found a soft component
with a blackbody temperature of 84–97 eV and a column
density of (11.7–15.5)×1020 cm−2. Our values of kTbb =
104+21
−23 eV and nH = (7.8± 3.9)× 10
21 cm−2 agree.
3.4. V2400 Oph
V2400 Oph was identified as a soft IP by
de Martino et al. (2004), who analysed a BeppoSAX ob-
servation and reported a blackbody temperature of 103±
10 eV and absorption column (46+12
−13)×10
20 cm−2. We
find a blackbody temperature of 117+33
−44 eV, in agreement
with this result, but a slightly lower absorption column of
(7.0+2.9
−4.9)× 10
20 cm−2. This is probably because de Mar-
tino used a single mekal and a single partial-covering
absorber to model the hard emission, whereas we used
the stratified column model and two partial covering ab-
sorbers.
3.5. PQ Gem
PQ Gem was the first IP found to have a soft-X-ray
component (Mason et al. 1992). This component was
also present in the XMM data, with a best-fitting black-
body temperature of 47.6+2.9
−1.4 eV, in agreement with the
46+12
−23 eV of Duck et al. (1994) from ROSAT data, and
56+12
−14 eV of de Martino et al. (2004) from BeppoSAX
data. The fitted column density goes to zero, which is
likely to be an artefact of fitting a complex absorption
with too simple a model. We quote an upper limit of
3× 1019 cm−2 based on the phase-resolved modelling of
Evans et al. (2006).
3.6. EX Hya
The best-fitting model for EX Hya used a blackbody
component, which has not been previously reported in
this system. However even with this component, the
procedure outlined in Section 2 resulted in a poor fit
=2.1). A possible reason for this is our choice of ab-
sorption model. We have used a cold absorber in our
models since the data do not warrant the extra param-
eters in ionised absorption models, even though one ex-
pects any absorbing material (e.g. the accretion curtains)
to be ionised. We therefore tried various ionised absorp-
tion models, but gained only a minor improvement to
the fit. We thus reverted to the cold absorber model for
consistency with the rest of this paper. We also tried us-
ing phase-resolved spectroscopy, in case the poor fit was
the result of averaging phase-variant parameters, how-
ever this still did not yield an acceptable fit. We have
nonethless included our results for EX Hya, for complete-
ness, but due to the poor fit, we do not much place much
weight on the EX Hya data when considering our results.
As the distance to EX Hya is known (64.5±1.2 pc:
Beuermann et al. (2003)), we can determine the size of
the accretion footprint from the soft X-ray flux. Table 4
gives this as (1.59+3.47
−0.62) × 10
−10 ergs cm−2 s−1 with a
temperature of 31.0+1.3
−2.4 eV, from which we compute an
emitting area (8.4+29.8
−4.2 ) × 10
13 cm2. Suleimanov et al
(2005) gave the mass of the white dwarf in EX Hya as
0.5±0.05M⊙, thus the observed blackbody emitting area
in EX Hya covers (7.3+29.3
−4.0 ) × 10
−4 of the white-dwarf
surface.
3.7. UU Col
UU Col was identified as a soft IP by Burwitz et al.
(1996). de Martino et al. (2006) have recently confirmed
this with a detailed analysis of the XMM observation.
They reported a blackbody temperature of 49.7+5.6
−2.9 eV,
which is lower than our value of 73+20
−9 eV, however in
their model the blackbody is absorbed by the partial cov-
ering absorber, and no simple absorber is present.
3.8. WX Pyx
The XMM observation of WX Pyx, the only X-ray
observation of this star to date, has a relatively low sta-
tistical quality. It was previously analysed by Schlegel
(2005) who did not report looking for a blackbody com-
ponent. However, we find that adding a blackbody does
significantly improve the fit.
3.9. Comparison with the polars
In Fig. 2 we have plotted the softness ratios of both the
IPs and the polars (from RC04). For the polars which
RC04 reported not to have a blackbody, we obtained the
spectra as extracted an calibrated by RC04 (Ramsay,
private communication), and fitted them in the same way
as the IPs (Section 2) to obtain an upper limit.
6 Evans & Hellier
Fig. 3.— The EPIC-pn spectra of the eight IPs for which the best-fitting models contain a blackbody component. The solid line shows
the hard component; the broken line the blackbody. For V405 Aur we have shown the MOS-1 spectrum, since the pn camera did not collect
any data.
The chief difference in the two distributions is that
while several polars show a softness ratio > 0.5, no IP
can be confirmed to do this, and it can be excluded for
all but EX Hya – for which our results are uncertain
(Section 3.6). The ‘soft excess’ in polars is believed to
arise due to ‘blobby accretion’ (e.g. Kuijpers & Pringle
(1982)). In this model, dense blobs of matter penetrate
into the white dwarf photosphere and the energy is ther-
malised to a blackbody.
Whether such accretion occurs in IPs has not been
widely discussed in the literature. Hellier & Beardmore
(2002) suggested that viscous interactions in an accre-
tion disc would destroy blobs, although Vrielmann et al.
(2005) interpreted flares in the lightcurve of GK Per as
resulting from the accretion of blobs. Our findings sug-
gest that blobby accretion is not significant in IPs.
4. DISCUSSION
The ‘polar’ class of magnetic cataclysmic variable has
long been known to be characterised by a soft black-
body component (e.g. King & Watson (1987)). This
Soft X-rays from IPs 7
Fig. 4.— The EPIC-pn spectra of the four IPs for which the best-fitting model does not contain a blackbody component. The solid line
shows the best-fitting model. The broken line shows the upper limit to a blackbody component, given a temperature of 60 eV. For FO Aqr
we show the MOS-1 data, as the signal-to-noise ratio of the pn data is worse.
is considered to arise from the white-dwarf surface,
heated either by reprocessing of hard X-rays from
the accretion column, or by thermalisation of blobs
of accretion (e.g. Kuijpers & Pringle (1982)). In con-
trast, IPs were thought to lack this component (e.g.
King & Lasota (1990)). However, observations with
ROSAT found a blackbody component in some IPs, lead-
ing Haberl & Motch (1995) to suggest that there were
two spectrally distinct classes of IP. This raised the ques-
tion of why.
To address this we have conducted a systematic survey
of the spectral characteristics of the IPs observed with
XMM-Newton, which has much greater spectral coverage
and throughput than ROSAT.
We find that, of twelve IPs analysed, eight show a soft
blackbody component while four do not. This suggests
that a blackbody is a normal component of IPs, and
hence of accretion onto magnetic white dwarfs, and that
the spectra differ only in degree.
We thus ask what causes the differing visibility
of the soft component. There does not appear to
be any correlation with the white-dwarf mass [see
Cropper et al. (1999); Ezuka & Ishida (1999); Ramsay
(2000); Suleimanov et al (2005) for mass estimates], nor
any obvious correlation with the orbital period.
In polars, systems with higher magnetic field strengths
appear to have higher softness ratios (e.g. Ramsay et al.
(1994)). Of the IPs in our sample showing polarisation,
and thus known to have a relatively strong field (5–20
MG), all (PQ Gem, V405 Aur and V2400 Oph) show
a blackbody component, while the four stars showing
no blackbody emission (FO Aqr, AO Psc, HT Cam &
V1223 Sgr) do not show polarisation. We give a possible
explanation for this after discussing the role of absorp-
tion.
We first consider the simple absorber, which is proba-
bly of interstellar origin. The detection of a blackbody
component in most systems shows that interstellar ab-
sorption is not sufficient to extinguish the soft emission.
Further, the systems with no detected blackbody compo-
nent do not have higher interstellar columns than those
with a blackbody (Table 3), so this absorption cannot
explain the differing visibility of the soft component.
We thus turn to the partial-covering absorption, which
in IPs is predominantly caused by the accretion cur-
tains crossing the line of sight. Here we find that
the systems where the lightcurves are dominated by
deep absorption dips owing to the accretion curtains
(FO Aqr, V1223 Sgr and AO Psc; see Beardmore et al.
(1998); Beardmore et al (2000); Hellier et al. (1991) re-
spectively) do tend to be those which lack a blackbody
component. In contrast, systems showing a blackbody
component, such as V405 Aur, NY Lup, EX Hya and
V2400 Oph, tend to be systems where the lightcurves
suggest that the accretion curtains do not hide the accre-
tion footprints (see Evans & Hellier (2004); Haberl et al.
(2002); Allan et al. (1998); Hellier & Beardmore (2002)
respectively).
We thus suggest that the major reason why some IPs
don’t show a blackbody component is simply that the
heated region near the accretion footprint is hidden by
the accretion curtains, while in other IPs it is not, the
difference being the result of the system inclination and
the magnetic colatitude (see Fig. 5). Coupled with this is
8 Evans & Hellier
to foreshortening
No BB seen
BB seen
BB seen
Little BB owing
Fig. 5.— The factors that affect blackbody emission in an IP.
a) When the upper magnetic pole is on the visible face, blackbody
emission will only be seen if the inclination is such that the heated
accretion region is visible above the accretion curtains.
b) When the lower pole is on the visible face, it will likely be too
foreshortened for us to detect blackbody emission.
c) In UU Col the magnetic axis is highly inclined, so the foreshort-
ening seen in b) is reduced and blackbody emission is seen.
the effect of foreshortening, such that an optically thick
heated region will not produce much blackbody emission
if it is only seen while on the white-dwarf limb, rather
than in the middle of the face.
A proper investigation of this idea would need knowl-
edge of the size and location of the accretion footprints
and of the surrounding heated polecaps, so that we could
estimate the difference absorbing columns of different
spectral components, and how these vary with spin-cycle
phase. However, this information is not known for the
majority of IPs. The softness ratio might conceivably
also vary with parameters such as accretion rate and
white-dwarf mass, which are again only poorly known.
However, as a test of our ideas, we can outline how
they might apply to the remaining systems in our sam-
ple which we did not consider when forming the model,
namely HT Cam, GK Per, PQ Gem and UU Col.
In PQ Gem the accretion curtains do cause an ab-
sorption dip when they obscure the accretion footprints.
However, the geometry of this star is relatively well de-
termined (Potter et al. 1997; Mason 1997; Evans et al.
2006) and it appears that the heated polecap is graz-
ingly visible above the accretion curtain for part of the
cycle; thus it shows both an absorption dip and a soft
blackbody, and is on the boundary between the two cases
illustrated in the upper panel of Fig. 5.
UU Col also shows an absorption dip when the ac-
cretion curtains obscure the upper pole, and also shows
blackbody emission. de Martino et al. (2006) proposed
that the blackbody emission comes from the lower pole,
viewed when that pole is closest to us (lowest panel of
Fig. 5). We thus suggest that UU Col has an abnor-
mally high inclination of the magnetic dipole, such that
the lower pole is not foreshortened as much as in other
IPs where no blackbody component is seen. V405 Aur
is another system that appears to have a highly inclined
dipole, such that blackbody emission from the lower pole
is significant, leading in that system to a double-peaked
soft-X-ray lightcurve (Evans & Hellier 2004).
In contrast to all the other IPs, the XMM data of
GK Per reported here were collected during an outburst.
Fig. 6.— Schematic diagram of GK Per in outburst. Accretion
occurs from all azimuths, resulting in a circular blackbody-emitting
region (dark ring). As can be seen, even when the accretion cur-
tains lie across our line-of-sight, part of this region is unobscured.
Hellier et al. (2004) have argued that during outburst the
accretion occurs from all azimuths, forming a complete
accretion ring at the poles. As illustrated in Fig. 6, this
means that some portion of the heated polecap is likely
to be visible ‘behind’ the magnetic pole, where accretion
does not normally occur. Thus in GK Per in outburst we
see a system with both strongly absorbed X-ray emission
(from in front of the magnetic pole) and a blackbody
component.
Lastly, we consider HT Cam. This shows very lit-
tle sign of absorption, and its lightcurve can explained
without any absorption effects (Evans et al. 2006). Yet
it shows no blackbody emission, in apparent contradic-
tion to our model. However, as previously suggested by
de Martino et al. (2006) and Evans et al. (2006), it ap-
pears that HT Cam has an exceptionally low accretion
rate (partly accounting for the lack of absorption). If so,
it could be that the blackbody component is simply too
cool to be detected in the XMM bandpass. We note that
the blackbody temperature in EX Hya, the other star in
our sample below the period gap, is lower than in the
others (Table 3), and that in HT Cam might be lower
still.
5. SUMMARY
We have analysed data from XMM observations of 12
intermediate polars and find that a soft blackbody com-
ponent is a common feature of their X-ray spectra. We
suggest that in the systems showing no blackbody emis-
sion the heated accretion polecaps are largely hidden by
the accretion curtains, or are only visible when on the
white dwarf limb and highly foreshortened. Thus IPs
with lightcurves dominated by absorption dips caused by
the passage of accretion curtains across the line of sight
tend to show no blackbody emission. Further, these are
also the systems least likely to show polarisation, since
the cyclotron-emitting column will also be obscured by
the accretion curtains, or would be beamed away from
us if the accretion region were on the white-dwarf limb.
After comparing the blackbody emission seen in IPs with
that seen in polars, we conclude that the blobby emis-
sion responsible for soft X-ray excesses in polars does not
occur in IPs.
ACKNOWLEDGEMENTS
We thank Gavin Ramsay for providing us with the
spectra of the polars with no detectable soft component.
Facilities: XMM ()
Soft X-rays from IPs 9
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