Dramatic Variability of X-ray Absorption Lines in the Black Hole
  Candidate Cygnus X-1
Dramatic Variability of X-ray Absorption Lines in the Black Hole
Candidate Cygnus X-1
Chulhoon Chang1 and Wei Cui1
Department of Physics, Purdue University, West Lafayette, IN 47907
ABSTRACT
We report results from a 30 ks observation of Cygnus X-1 with the High En-
ergy Transmission Grating Spectrometer (HETGS) on board the Chandra X-ray
Observatory. Numerous absorption lines were detected in the HETGS spectrum.
The lines are associated with highly ionized Ne, Na, Mg, Al, Si, S, and Fe, some
of which have been seen in earlier HETGS observations. Surprisingly, however,
we discovered dramatic variability of the lines over the duration of the present
observation. For instance, the flux of the Ne X line at 12.14 Å was about 5×10−3
photons cm−2 s−1 in the early part of the observation but became subsequently
undetectable, with a 99% upper limit of 0.06 × 10−3 photons cm−2 s−1 on the
flux of the line. This implies that the line weakened by nearly two orders of
magnitude on a timescale of hours. The overall X-ray flux of the source did also
vary during the observation but only by 20–30%. For Cyg X-1, the absorption
lines are generally attributed to the absorption of X-rays by ionized stellar wind
in the binary system. Therefore, they may provide valuable diagnostics on the
physical condition of the wind. We discuss the implications of the results.
Subject headings: binaries: general — black hole physics — stars: winds,outflows
— stars: individual (Cygnus X-1) — X-rays: binaries
1. Introduction
Cygnus X-1 is the first dynamically-determined black hole system (Webster & Murdin
1972; Bolton 1972). It is in a binary system with a massive O9.7 Iab supergiant, and the
orbital period was determined optically to be 5.6 days. Cyg X-1 is thus intrinsically different
from the majority of known black hole candidates (BHCs) whose companion stars are much
1Email: chang40@physics.purdue.edu, cui@physics.purdue.edu
http://arxiv.org/abs/0704.0452v1
– 2 –
less massive (see review by McClintock & Remillard 2006). Curiously, those that have a
high-mass companion (including Cyg X-1, LMC X-1 and LMC X-3) are all persistent X-
ray sources, while those that have a low-mass companion are exclusively transient sources.
Perhaps, stellar wind from the companion star plays a significant role in this regard (Cui,
Chen, & Zhang 1997).
Unlike transient BHCs, in which mass accretion is mediated by the companion star
overfilling its Roche-lobe, Cyg X-1 is thought to be a wind-fed system. In this case, however,
the wind is thought to be highly focused toward the black hole, because the companion star
is nearly filling its Roche lobe (Gies & Bolton 1986). The observed orbital modulation of
the X-ray flux (Wen et al. 1999; Brocksopp et al. 1999a) has provided tentative evidence for
wind accretion, because it is probably caused by varying amount of absorption through the
wind (Wen et al. 1999). On the other hand, it is generally believed that an accretion disk
is also present, based on the presence of an ultra-soft component, as well as Fe Kα emission
line in the X-ray spectrum (e.g., Ebisawa et al. 1996; Cui et al. 1998).
Cyg X-1 is probably still the most studied BHC. It is a fixture in the target list for
all major X-ray missions. Much has been learned from modeling the X-ray continuum of
the source, as well as from studying its X-ray variability. A recent development is the
availability of high-resolution X-ray data, which may shed further light on the accretion
process and the environment within the binary system. Cyg X-1 has been observed on many
occasions with the High Energy Transmission Grating Spectrometer (HETGS) on board the
Chandra X-ray Observatory and the Reflection Grating Spectrometer on board the XMM-
Newton Observatory. The high-resolution spectra have revealed the presence of numerous
absorption lines that are associated with highly ionized material (Marshall et al. 2001; Schulz
et al. 2002; Feng et al. 2003; Miller et al. 2005).
In this work, we report the detection of absorption lines, some of which have been seen
previously but are much stronger here, and, more surprisingly, the discovery of dramatic
variability of the lines, based on data from our HETGS/Chandra observation of Cyg X-1
during its 2001 state transition. The fluxes of some of the lines varied by nearly two orders
of magnitude over the duration of the observation, while the overall flux of the source varied
only mildly.
2. Observations and Data Reduction
Cyg X-1 made a rare transition between the low-hard state and the high-soft state, as
seen by the All-Sky Monitor (ASM) on the Rossi X-ray Timing Explorer (RXTE), about
– 3 –
five years after a similar episode in 1996 (Cui et al. 1997a and 1997b). Figure 1 shows the
ASM light curve that covers the entire period. In this case, the flux of the source stayed
high for ∼400 days, about twice as long as in 1996. Otherwise, the two episodes are very
similar, including the flux levels of the two states, prominent X-ray flares in both states, and
rapid transitions.
Triggered by the ASM alert, the source was observed from 2001 October 28 16:13:52
to October 29 00:33:52 (UT) with the HETGS on Chandra (ObsID #3407). The HETGS
consists of two gratings: Medium Energy Grating (MEG) and High Energy Grating (HEG).
After passing the gratings, the photons are recorded and read out with the spectroscopic
array of the Advanced CCD Imaging Spectrometer (ACIS). To avoid photon pile-up in
the dispersed events, we chose to run the ACIS in continuous-clocking (CC) mode. We also
applied a spatial window around the aim point to accept every 10th event in the zeroth order,
to prevent telemetry saturation yet still have a handle on the position of the zeroth-order
image for accurate wavelength calibration.
The Chandra data were reduced and analyzed with the standard CIAO analysis package
(version 3.2). Following the CIAO 3.2 Science Threads1, we prepared and filtered the data,
produced the Level 2 event file from the Level 1 data products, constructed the light curves,
and made the spectra and the corresponding response matrices and auxilliary response files
(ARFs). We did have to work around a problem related to the use of maskfile for the CC-
mode data, when we were making the ARFs. The solution is now included in the CIAO
3.3 Science Threads. To verify wavelength calibrations, we compared the plus and minus
orders and found good agreement on the position of prominent lines. We did not subtract
background from either the light curves or the spectra, because it is not obvious how to select
background events from the CC-mode data. For a bright source like Cyg X-1, however, we
expect the effects to be entirely negligible.
In coordination with the Chandra observation, we also observed Cyg X-1 with the Pro-
portional Counter Array (PCA) and High Energy X-ray Timing Experiment (HEXTE) de-
tectors on board RXTE. The PCA and HEXTE covers roughly the energy ranges of 2–60
keV and 15-250 keV, respectively. In this work, we made use of the combined broad spec-
tral coverage of the two detectors to constrain the X-ray continuum more reliably. The
RXTE observation was carried out in six short segments, with exposure times of 1–3 ks.
The data were reduced with the standard HEASOFT package (version 5.2), along with the
associated calibration files and background models. We followed the usual procedures2 in
1See http://asc.harvard.edu/ciao3.2/threads/index.html
2See http://heasarc.gsfc.nasa.gov/docs/xte/recipes/cook book.html
http://asc.harvard.edu/ciao3.2/threads/index.html
http://heasarc.gsfc.nasa.gov/docs/xte/recipes/cook_book.html
– 4 –
preparing, filtering, and reducing the data, as well as in deriving light curves and spectra
from the standard-mode data. While the HEXTE background was directly measured from
off-source observations, the PCA background was estimated based on the background model
appropriate for bright sources.
3. Analysis and Results
3.1. Light Curves
Figure 2 shows the Chandra light curve of Cyg X-1 that is made from the MEG first-
order data (with plus and minus one orders coadded). For comparison, we have over-plotted
the average count rate of PCU #2 for each of the RXTE exposures. The agreement is quite
good. Besides the rapid flares and other short-term variability that are expected in Cyg X-1,
the source also varies significantly on a timescale of hours. The MEG count rate rises from
a baseline level of roughly 170 counts s−1 to a peak of roughly 220 counts s−1 within 3–4
hours and quickly goes back down to the baseline level. We somewhat arbitrarily divided
the light curve into two time periods (which are labeled as Period I and Period II in Fig. 2)
for subsequent analyses.
3.2. High-Resolution Spectroscopy
We analyzed and modeled the HETGS spectra in ISIS (version 1.2.9) (Houck & Denicola
2000)3. For this work, we focused only on the first-order spectra, because of the relatively
poor signal-to-noise ratio of higher-order spectra. For both the MEG and HEG, we first co-
added the plus and minus orders to produce an overall first-order spectrum, again following
the appropriate CIAO 3.2 science threads. The resolution of the raw data is about 0.01 Å
and 0.005 Å for the MEG and HEG, respectively, which represents a factor of 2 over-sampling
of the instrumental resolution. We applied no further binning of the data. Each spectrum
was broken into 3-Å segments for subsequent analyses. Each segment was fitted locally with
a model that consists of a multi-color disk component and a power law for the continuum
and negative or positive Gaussian functions for narrow absorption or emission features, with
the interstellar absorption taken into account. Such a continuum model is typical of BHCs.
However, the best-fit continuum differs among the segments or between the MEG and HEG,
presumably due to remaining calibration uncertainties. Since we are only interested in using
3See also http://space.mit.edu/CXC/ISIS/
http://space.mit.edu/CXC/ISIS/
– 5 –
the Chandra data to study lines, we think that the adopted procedure is justified. We use
the RXTE data to more reliably constrain the continuum.
One thing that one notices right away is the presence of many absorption lines in the
high-resolution spectra. For this work, we consider a feature real if it is present both in
the MEG and HEG data, with a significance of above 4σ. Figure 3 shows a portion of the
MEG first-order spectrum for Period I, to highlight the lines detected. No absorption lines
were found at wavelengths above 15 Å. We should point out that we also see the emission-
like features at 6.74 Å and 7.96 Å, which most likely instrumental artifacts associated with
calibration uncertainties around the Si K and Al K edges (Miller et al. 2005 and references
therein). We have identified the absorption lines with highly ionized species of Ne, Na, Mg,
Al, Si, S, and Fe, based mostly on the atomic data in Verner et al. (1996) and Behar et al.
(2002), but also in ATOMDB4 1.3.1 for additional transitions.
The line identification process involves three steps: (1) a transition is considered a
candidate if the theoretical wavelength is within 0.03 Å of the measured value; (2) in cases
where multiple candidates exist based on (1), the most probably one (based both on the
oscillator strength of the transition and the relative abundance of the element) is chosen;
and (3) consistency check is made to avoid the identification of a line when more probable
transitions of the same ion are not seen. The last step is critical. For instance, we initially
associated the lines at 10.051 Å and 11.029 Å with Fe XVIII 2s22p5–2s22p46d and Fe XVII
2s22p6–2s2p64p, respectively, because, assuming solar abundances, they are expected to be
stronger than Na XI 1s–2p and Na X 1s2–1s1p, which were also viable candidates from
Step 1. However, we did not detect other more probable transitions associated with Fe
XVIII and Fe XVII, which made the identifications highly unlikely. We had to conclude
that Fe is under-abundant by a factor of 2–3, at minimum, so that we could associate the
lines with Na instead. If so, the line at 7.480 Å would more likely be associated with Mg XI
1s2–1s4p, as opposed to Fe XXIII 2s2–2s5p, which we initially identified. The problem with
this is that we saw no hint of Mg XI 1s2–1s3p, which is more probable. Therefore, we had
to conclude similarly that Mg is also under-abundant by at least a similar amount (but not
by too much, as constrained by other Mg lines).
Table 1 show all of the absorption lines that we detected and identified. The flux and
equivalent width (EW) of each line shown were derived from the best-fit Gaussian for the
line, as well as the local continuum around the line for the latter. Note that in a few cases
our identifications are different from those in the literature. They include: Fe XXII at 8.718
Å and Fe XXI at 9.476 Å (cf. Marshall et al. 2001; Miller et al. 2005); Fe XX at 10.12 Å,
4See http://cxc.harvard.edu/atomdb
http://cxc.harvard.edu/atomdb
– 6 –
Na X at 11.0027 Å, and Fe XX at 12.82 Å (cf. Marshall et al. 2001); and Fe XXI 11.975 Å
(cf. Miller et al. 2005). Only about half of the absorption lines that we detected have been
seen previously in Cyg X-1 (Marshall et al. 2001; Schulz et al. 2002; Feng et al. 2003; Miller
et al. 2005). Many of these lines are much stronger in our case. On the other hand, we
detected nearly all of the reported absorption lines. The exceptions include: S XVI at 4.72
Å (Marshall et al. 2001; Feng et al. 2003), which is detectable only at the 3σ level in our
case (and is thus not included in the table), Fe XIX at 14.53 Å (Miller et al. 2005), which is
actually detectable at about the 4σ level here (but just misses our threshold), and Fe XIX
at 14.97 Å and Fe XVIII at 16.01 Å (Miller et al. 2005), whose presence is not apparent in
our data (< 3σ). Miller et al. (2005) also reported a line at 7.85 Å (Mg XI) but only at
the 3σ level. The line can also be seen in our data at the 4σ level (but also just misses our
threshold). Therefore, we already see some indication that the absorption lines in Cyg X-1
may be variable from a comparison of our results with those published.
Still, it is surprising that almost all of the absorption lines become undetectable in
Period II. Figure 4 shows the MEG first-order spectrum for this time period, which can be
directly compared to results in Fig. 3. The only exception is the line at 14.608 Å, which is
detected with a significance of 5σ in Period II. We derived 99% upper limits on the flux and
EW of each line seen in Period I but not in Period II. The results are also summarized in
Table 1, for direct comparison. As an example, we examine the Ne X line at 12.1339 Å in
the two periods. The integrated flux of the line is about 5× 10−3 photons cm2 s−1 in Period
I, while its 99% upper limit for Period II is only 0.06×10−3 photons cm2 s−1. Therefore, the
line weakened by nearly two orders of magnitude in flux over a timescale of merely several
hours. This is the first time that such dramatic variability of the lines has been observed in
any BHC. While this is the most extreme case, other lines also show large variability (see
Table 1).
To quantify the column density of each ion required to account for the corresponding
absorption line detected and its variability, we carried out curve-of-growth analysis, following
Kotani et al. (2000). The atomic data used in the analysis were again taken from Verner
et al. (1996), Behar et al. (2002), and also ATOMDB 1.3.1 in some cases. The analysis
assumes that the width of the lines is due entirely to thermal Doppler broadening. For
resolved lines, we derived the characteristic temperature from the measured widths. For
unresolved lines, on the other hand, we adopted a temperature that would lead to a line
width equal to the resolution of the MEG (0.023 Å at FWHM). In these cases, therefore, the
derived column density only represents a lower limit. The results are shown in Table 1. This
explains why, e.g., the density of Ne X derived from the 12.144-Å line (which is unresolved)
is significantly lower than that from the 10.245-Å line. Note, however, that the latter is
much lower that that from the 9.727-Å line. We think that the inconsistency arises from the
– 7 –
fact that the 9.727-Å line is likely a mixture of the Ne line and the Fe XIX line at 9.700 Å.
Similar inconsistency is also apparent in a few other cases (see Table 1), which may originate
similarly in line blending. It is also worth noting that most lines that we have analyzed fall
on the linear part of the curve of growth. All resolved lines at wavelengths λ > 11.7 Å are
saturated; so is the unresolved line at 14.203 Å. To show the degree of variability, we also
derived a 99% upper limit on the column density of each of the ions for Period II (assuming
the same characteristic temperatures).
3.3. X-ray Continuum
We now use the RXTE data to constrain the X-ray continuum of Cyg X-1. Both the
PCA and HEXTE data were used. For the PCA, we used only data from the first xenon layer
of each PCU, which is most accurately calibrated. Consequently, the PCA spectral coverage
is limited to roughly 2.5–30 keV. We relied on the HEXTE data to extend spectral coverage
to higher energies. The PCA consists of five detector units, known as PCUs. Not all PCUs
were always operating. For simplicity, we used only data from PCU #0 and PCU #2, which
stayed on throughout the observation, in the subsequent modeling. We chose to derive a
spectrum for each PCU separately, as well as for each of the two HEXTE clusters. Residual
calibration errors were accounted for by adding 1% systematic uncertainty to the data. We
also rebinned the HEXTE spectra to a signal-to-noise ratio of at least 3 in each bin. The
individual spectra were then jointly fitted with the same model that includes a multi-color
disk component, a broken power-law that rolls over exponentially beyond a characteristic
energy, and a Gaussian, taking into account the interstellar absorption. We also multiply
the model by a normalization factor that is fixed at unity for PCU #2 but is allowed to
vary for other detectors, in order to account for any uncalibrated difference in the overall
throughput among the detectors. The model fits the data well for all six segments in the
sense that the reduced χ2 is near unity (with 169 degrees of freedom).
The spectral shape of Cyg X-1 varied little from one observation to the other. The best-
fit photon indices are ∼2.1 and ∼1.7 below and above ∼10 keV. The roll-over energy stays at
20–21 keV and the e-folding energy roughly at 120–130 keV. Neither the disk component nor
the absorption column density is well constrained by the data, due to the lack of sensitivity
(and, to some extent, large calibration uncertainty) at low energies. The results can be
compared directly with those of Cui et al. (1997a), who applied the same empirical model
to the RXTE spectra of Cyg X-1 during the 1996 transition. It is quite apparent that the
spectra here are significantly harder, implying that the source was certainly not yet in the
true high-soft state (see Cui et al. 1997a and 1997b for discussions on the “settling period”).
– 8 –
From the long-term ASM monitoring data, we can see that Cyg X-1 was brighter during
our observation than during any of the earlier Chandra observations (Marshall et al. 2001;
Schulz et al. 2002; Miller et al. 2005), but not as bright as during a later observation (Feng
et al. 2003), when the high-soft state appears to have been reached.
3.4. Photoionization Modeling
To shed light on the physical properties of the absorber, we carried out a photo-ionization
calculation with XSTAR version 2.1 kn35. The underlying assumption is that the absorber
is photo-ionized by the X-ray radiation from the vicinity of the black hole. The input
parameters include the 0.5-10 keV luminosity (Lx = 3.11× 10
37 erg s−1, for a distance of 2.5
kpc), power-law photon index (2.1), both of which are based on results from modeling one of
the RXTE spectra with an assumed NH value of 5.5×10
21 cm−2 (Ebisawa et al. 1996; Cui et
al. 1998). We should note that it is in general risky to extrapolating the assumed power-law
spectrum to lower energies, because it could severely over-estimate the flux there. For the
lines of interest here, however, only ionizing photons with energies > 1 keV contribute and
the spectrum of those photons is described fairly well by a power law, because the effective
temperature of the disk component is expected to very low (e.g., Ebisawa et al. 1996; Cui
et al. 1998).
One of the outputs of the calculation is abundances of the ions of interest, as a function
of the ionization parameter, ξ = Lx/nr
2, where n is the density of the absorber and r is the
distance of the absorber to the source of ionizing photons. Using these abundance curves,
we could, in principle, constrain ξ to a range that is consistent with the ratio of the densities
of any two ions of an element. The challenge in practice is, as already mentioned, that many
of the lines are likely a blend of multiple transitions (of comparable probabilities), which
makes it difficult to reliably determine the densities. Nevertheless, we made an attempt
at deriving such constraints with the resolved, non-mixed lines. Figure 5 summarizes the
results. The intervals do not all overlap, which implies that no single value of ξ could account
for all the data. This is supported by the fact that we have detected all the lines that are
expected for ionization parameters in the range of roughly 102.5–104.5. If one assumes that
the “absorbers” are thin shells along the line of sight and that they all have the same density,
e.g., n = 1011 cm−3 (see, e.g., Wen et al. 1999), one would have 1011 cm . r . 1012 cm.
Compared with the estimated the distance between the compact object and the companion
star (∼ 1.4 × 1012 cm; LaSala et al. 1998), this would put the “absorbers” within the
5See http://heasarc.gsfc.nasa.gov/docs/software/xstar/xstar.html
http://heasarc.gsfc.nasa.gov/docs/software/xstar/xstar.html
– 9 –
binary system. One should, however, take the results with caution, because of, e.g., gross
over-simplification regarding the geometry of the “absorbers”.
4. Discussion
The observed dramatic variability of the absorption lines might be caused by a change
in the degree of ionization in the wind. Since the overall X-ray luminosity varied only mildly,
we speculate that it probably arises from a sudden change in the density of the wind. There
is evidence that such a change could occur during a state transition or during flares (Gies et
al. 2003). If the moderate decrease in the ionizing flux is accompanied by a more dramatic
reduction in the density of the wind from Period I to Period II, the ionization parameter
might increase sufficiently to cause a total ionization of the wind in Period II and thus the
disappearance (or significant weakening) of the lines. It is worth noting that in Period I
lighter elements seen are all H- or He-like but Fe is in an intermediate ionization state (as
indicated by the absence of H- or He-like ions), suggesting a high but not extreme degree of
ionization in the period. Conversely, a dramatic increase in the wind density could achieve
the same effect. Numerical simulations of similar wind-accretion systems (e.g., Blondin,
Stevens, & Kallman 1991) have revealed not only a significant jump in the column density
at late orbital phases (& 0.6) that is associated with tidal streams but also large variability of
the absorbing column. It is conceivable that Period II might coincide with a sudden increase
in the column density. Since we found no apparent absorption lines that correspond to a
lower degree of ionization in Period II, however, such lines must be outside the spectral range
covered with our data, in order for the scenario to be viable. A quantitative assessment of
these scenarios is beyond the scope of this work.
Many of the absorption lines detected by Miller et al. (2005) are much stronger during
Period I of our observation (see Table 1). Using only the lines that are detected with a
significance above 5σ, we looked for a systematic red- or blue- shift of the lines, following
up on the reported redshift of the lines by Marshall et al. (2001) based on data taken in
the low-hard state. The results are summarized in Figure 6 (in the left panel). In this
case, although the lines are still systematically redshifted on average, there is not an obvious
single-velocity solution. Interestingly, if we limit the results only to those lines that were
used by Marshall et al. (2001), as shown in the right panel of Figure 6, we would arrive at
an average velocity that is very close to what Marshall et al. reported, although the scatter
of data points is much larger in our case. On the other hand, our observation spans binary
orbital phases from 0.85 to 0.92, according to the most updated ephemeris (Brocksopp et al.
1999b), while Marshall et al.’s covers a phase range of 0.83–0.86. If the redshift of the lines
– 10 –
is related to the focused-wind scenario advocated by Miller et al. (2005), we ought to see a
larger (by about 30%) redshift. Given the large uncertainties, as well as the possibility that
the wind geometry might be different for different states, it is difficult to draw any definitive
conclusions.
Feng et al. (2003) reported the detection of a number of absorption lines of asymmetric
profile, when Cyg X-1 was in the high-soft state, which they interpreted as evidence for
inflows. The same lines are also present in our data during Period I and are, in fact, much
stronger (except for S XVI, as noted in § 3.2). Figure 7 shows an expanded view of the
Si and Mg lines, which are the strongest in the group. As is apparent from the figure, the
line profile can be fitted fairly well by a Gaussian function in all cases. Therefore, the lines
show no apparent asymmetry here. This also seems to be the case for the S and Fe lines,
although the statistics of the data are not as good. Taken together, our results and Feng et
al.’s imply that that the phenomenon is either unique to the high-soft state (in which Feng
et al. made the observation) or is intermittent in nature. We should also note that Feng et
al’s observation was carried out around the superior conjunction (i.e., zero orbital phase),
where absorption due to the wind is expected to be the strongest (e.g., Wen et al. 1999).
It is not clear, however, how such additonal absorption would lead to an asymmetry in the
profile of lines.
No emission lines are apparent in our data. Evidence for weak emission lines has been
presented (Schulz et al. 2002; Miller et al. 2005) but the significance is marginal in all cases.
On the other hand, several absorption edges are easily detected in our data (see Figs. 3 and
4), as first reported and studied in detail by Schulz et al. (2002). The edges can almost
certainly be attributed to the interstellar absorption.
We thank Harvey Tananbaum for approving this DDT observation, Norbert Schulz and
Herman Marshall for helpful discussion on the pros and cons of various observing configura-
tions, John Houck for help with the use of ISIS and Tim Kallman for help on using XSTAR,
and David Huenemoerder for looking into issues related to the Chandra data products. We
acknowledge the use of the curve-of-growth analysis script that Taro Kotani has made pub-
licly available. We also thank the anonymous referee for a number of useful comments that
led to significant improvement of the manuscript. Support for this work was provided in
part by NASA through the Chandra Award DD1-2011X issued by the Chandra X-ray Ob-
servatory Center, which is operated by the Smithsonian Astrophysical Observatory for and
on behalf of NASA under contract NAS8-03060, and through the LTSA grant NAG5-9998.
– 11 –
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Webster, B. L., & Murdin, P., 1972, Nature, 235, 37
Wen, L., Cui, W., Levine, A. M., & Bradt, H. V. 1999, ApJ, 525, 968
This preprint was prepared with the AAS LATEX macros v5.2.
– 13 –
Table 1: Detected Absorption Lines
Theoretical Measured Shift Flux(I) Flux(II) EW(I) EW(II) Nz(I) Nz(II)
Ion and Transition (Å) (Å) (km s−1) (10−3ph cm−2s−1) (mÅ) (mÅ) (10−16cm−2) (10−16cm−2)
S XV 1s2-1s2p 5.039b 5.041(3) 120±180 1.8(3) <0.9 3.8(7) <1.5 2.5(5) ≤1.0
Si XIV 1s-2p 6.1822a 6.189(1) 330±50 3.0(2) <0.6 6.9(4) <1.0 5.4(3) ≤0.7
Si XIII 1s2-1s2p 6.648b 6.657(2) 410±90 2.6(2) <1.7 6.1(5) <2.9 2.4(2) ≤1.1
Mg XII 1s-3p 7.1062a 7.119(3) 540±130 1.2(2) <0.9 2.6(5) <1.4 8(2) ≤4.0
Al XIII 1s-2p 7.1727a 7.191(
) 760
1.3(2) <1.7 3.0(5) <2.7 1.6(3)d ≤1.5
Fe XXIII 1s22s2p-1s22s6d 7.2646c 7.268(5) 140±210 1.4(3) <2.0 3.1(6) <3.3 29(6) ≤31
Fe XXIII 1s22s2-1s22s5p 7.4722a 7.480(
) 310
0.9(2) <1.4 1.9(5) <2.3 5(1) ≤6.5
Fe XXIV 1s22s-1s24p 7.9893a 8.004(5) 550±190 2.2(3) <1.3 4.7(7) <2.2 9(1) ≤4.0
Fe XXII 1s22s22p-1s22s25d 8.0904c 8.096(3) 210±110 1.0(2) <0.6 2.1(5) <1.0 8(2)d ≤3.6
Fe XXII 1s22s22p-1s22s25d 8.1684c 8.166(3) -90±110 1.0(2) <1.1 2.3(5) <1.9 9(2)d ≤7.2
Fe XXIII 1s22s2-1s22s4p 8.3029a 8.319(2) 580±70 3.0(3) <0.6 6.6(6) <1.0 6.4(6) ≤0.9
Mg XII 1s-2p 8.4210a 8.428(1) 250±40 5.0(3) <0.4 10.8(6) <0.7 4.7(3) ≤0.3
Fe XXI 1s22s22p2-1s22s22p5d 8.573a 8.577(5) 140±170 1.3(3) <0.7 2.8(7) <1.2 6(2) ≤2.6
Fe XXII 1s22s22p-1s22s2p1/24p3/2 8.718
c 8.735(
) 580
2.2(3) <0.5 4.7(
) <0.9 7(1) ≤1.3
Fe XXI 1s22s22p2-1s22s2p1/22p3/24p3/2 8.8254
c 8.823(
) -80
1.4(3) <0.9 3.0(
) <1.4 21(
) ≤9.4
Fe XXII 1s22s22p-1s22s24d 8.98a 8.978(
) -70
1.8(3) <0.9 3.9(6) <1.6 4.6(8)d ≤1.8
Fe XXII 1s22s22p-1s22s24d 9.07a 9.083(
) 430
) <0.6 4.0(
) <1.0 5.2
Mg XI 1s2-1s2p 9.170b 9.192(
) 720
6.2(4) − 13.5(9) − 2.9(2) −
Fe XXI 1s22s22p2-1s22s22p4d 9.356a 9.378(5) 700±160 1.9(4) <2.8 4.3
<4.8 9(2) ≤10
Fe XXI 1s22s22p2-1s22s22p4d 9.476a 9.478(1) 60±30 3.8(3) <0.8 8.6(
) <1.5 6.6(
)d ≤1.0
Fe XIX 1s22s22p4-1s22s22p3(2D)5d 9.68a 9.700(4) 620±120 4.4(
) <1.8 9(1) <3.1 30(3) ≤9.8
Ne X 1s-4p 9.7082a 9.727(
) 580
) <0.7 5.3(
) <1.2 24(4) ≤5.0
Fe XX 1s22s22p3-1s22s22p2(3P)4d 9.991a 10.000(1) 270±30 3.7(4) <0.7 8.3(8) <1.3 5.6(6)d ≤0.8
Na XI 1s-2p 10.0250a 10.051(2) 780±60 5.6(
) <2.0 13(1) <3.5 4.0(3) ≤1.0
Fe XX 1s22s22p3-1s22s22p2(3P)4d 10.12a 10.127(3) 210±90 1.7(4) <0.3 3.8(9) <0.6 2.3(6)d ≤0.4
Ne X 1s-3p 10.2389a 10.245(
) 180
) <0.5 6(1) <0.9 9(2) ≤1.2
Fe XXIV 1s22s-1s23p 10.619a 10.631(
) 340
9.1(6) <4.6 22(2) <8.9 10(1) ≤3.7
Fe XXIV 1s22s-1s23p 10.663a 10.674(3) 310±80 5.6(6) <3.2 14(2) <6.2 12(2) ≤5.0
Fe XIX 1s22s22p4-1s22s22p3(4S)4d 10.816c 10.818(5) 60±140 7.0(9) <3.0 18(2) <6.0 14(2) ≤4.3
Fe XXIII 1s22s2-1s22s3p 10.981a 10.990(1) 230±30 5.5(5) <2.7 14(1) <5.6 2.4(2)d ≤0.8
Na X 1s2-1s2p 11.0027a 11.029(2) 720±50 6.3(
) <4.4 17(2) <9.2 2.7(4) ≤1.3
Fe XVIII 1s22s22p5-1s22s22p4(1D)4d 11.326c 11.33(1) 100±260 8(1) <5.1 23(
) <11 24(
) ≤11
Fe XXII 1s22s22p-1s22s2p(3P0)3p 11.44a 11.431(1) -240±30 7.3(6) <2.5 21(2) <5.7 8(1)d ≤1.6
Fe XXII 1s22s22p-1s22s2p(3P0)3p 11.51a 11.500(3) -260±80 4.3(
) <0.6 12(2) <1.4 22(
) ≤2.3
Fe XXII 1s22s22p-1s22s23d 11.77a 11.781(
) 280
12.6(9) <3.8 39(3) <9.3 7.5
Fe XXI 1s22s22p2-1s22s2p23p 11.975c 11.982(2) 180±50 9.1(9) − 29(3) − 17(
Ne X 1s-2p 12.1339a 12.144(
) 250
4.7(7) <0.06 16(2) <0.2 3.9(
)d ≤0.04
Fe XXI 1s22s22p2-1s22s22p3d 12.259a 12.247(
) -290
) <5.4 14(3) <14 10(3)d ≤10
Fe XXI 1s22s22p2-1s22s22p3d 12.285a 12.304(2) 460±50 22(1) <6.8 75(5) <18 12(2) ≤1.6
Fe XXI 1s22s22p2-1s22s22p3d 12.422c 12.438(2) 390±50 3.9(8) <2.9 14(3) <8.1 3.0(
)d ≤1.6
Fe XX 1s22s22p3-1s22s2p33p 12.576c 12.583(
) 170
7(1) <3.6 26(4) <10 19(
) ≤5.2
Fe XX 1s22s22p3-1s22s22p2(3P)3d 12.82a 12.844(2) 560±50 26(2) <6.5 100(6) <20 34
Fe XX 1s22s22p3-1s22s22p1/22p3/23d 12.912
c 12.914(3) 50±70 12(2) <4.7 47(6) <14 33
Fe XX 1s22s22p3-1s22s22p1/22p3/23d 12.965
c 12.953(
) -280
12(1) − 48(6) − 56
Ne IX 1s2-1s2p 13.448b 13.448(
9(2) <8.5 38(9) <30 6(
) ≤4.1
Fe XIX 1s22s22p4-1s22s22p1/22p
3d 13.479c 13.482(3) 70±70 5(1) <3.3 20(
) <12 1.0(
)d ≤0.5
Fe XIX 1s22s22p4-1s22s22p3(2D)3d 13.518c 13.523(3) 110±70 16(2) − 70(
) − 24
Fe XVIII 1s22s22p5-1s22s22p4(1D)3d 14.203a 14.220(3) 360±40 6(1) <4.0 32(7) <18 4(
)d ≤1.5
Fe XIX 1s22s22p4-1s22s22p3(2P)3s 14.60a
14.608(5)e 100±100 19(3) 20(4) 81±13 73±14 200
−60Fe XVIII 1s22s22p5-1s22s22p4(3P)3d 14.610a
Notes. — Results for Periods I and II are both shown for comparison. The errors in parentheses indicate uncertainty in the last digit of the
measurement; 1σ errors are shown. Negative flux or EW upper limits (indicating emission) are not shown.
a Verner et al. (1996); b Behar et al. (2002); c ATOMDB 1.3.3; d Unresolved; e The two transitions are equally probable. The average wavelength
was used to derive the Doppler-shift of the line.
– 14 –
Fig. 1.— Daily-averaged ASM Light Curve of Cyg X-1 during the 2001 state transition. The
vertical line indicates the time of the Chandra observation.
– 15 –
Fig. 2.— X-ray Light Curves of Cyg X-1. The solid curve shows data from the MEG first
order, while the horizontal bars show the average count rates from PCU #2. The error bars
are negligible in both cases. The dashed line defines the two time periods for subsequent
analyses (see text).
– 16 –
Fig. 3.— MEG first-order spectrum of Cyg X-1 for Period I. No binning was applied. The
presence of absorption lines are apparent. The identifications of the lines are shown. Note
that the emission-like features at 6.74 Å and 7.96 Å are likely instrumental (see text). The
solid line shows the best-fit to the (local) continuum.
– 17 –
Fig. 4.— As in Fig. 3, but for for Period II. Note the absence of nearly all the absorption
lines seen in Fig. 3.
– 18 –
2.5 3 3.5 4
Ionization Parameter (log ξ)
Fe XX/Fe XXII
Fe XX/Fe XXIV
Fe XX/Fe XIX
Fe XX/Fe XXI
Fe XX/ Fe XXIII
Si XIV/Si XIII
Ma XII/Mg XI
Ne X/Ne IX
Fig. 5.— Allowed ranges of the ionization parameter, each of which is inferred from the ratio
of the average densities of two ions of the same element.
4 6 8 10 12 14
Wavelength (Å)
4 6 8 10 12 14
Wavelength (Å)
Fig. 6.— Inferred Doppler shift of the selected absorption lines, (left) all the lines with a
significance above 5σ and (right) only the lines that were used by Marshall et al. (2001).
The dotted line shows the average Doppler velocity in both cases.
– 19 –
Fig. 7.— Profiles of the selected absorption lines. In each case, the dot-dashed histogram
shows a fit to the profile with a Gaussian function.
	Introduction
	Observations and Data Reduction
	Analysis and Results
	Light Curves
	High-Resolution Spectroscopy
	X-ray Continuum
	Photoionization Modeling
	Discussion
