# Spectral analysis of the AMXP IGR J17591–2342 during its 2018 outburst

A. Manca,1 A. F. Gambino,2 A. Sanna,1,3 G. K. Jaisawal,5 T. Di Salvo,2,3,4 R.
Iaria,2 S. M. Mazzola,1 A. Marino,8,9 A. Anitra,2 E. Bozzo,7 A. Riggio,1,4 and
L. Burderi,1,3,4
1Dipartimento di Fisica, Università degli Studi di Cagliari, SP Monserrato-
Sestu, KM 0.7, Monserrato, 09042 Italy
2Università degli Studi di Palermo, Dipartimento di Fisica e Chimica - Emilio
Segrè, via Archirafi 36 - 90123 Palermo, Italy
3INFN, Sezione di Cagliari, Cittadella Universitaria, 09042 Monserrato, CA,
Italy
4INAF - Osservatorio Astronomico di Cagliari, via della Scienza 5, 09047
Selargius (CA), Italy
5National Space Institute, Technical University of Denmark, Elektrovej
327-328, 2800 Lyngby, Denmark
6Istituto Nazionale di Astrofisica, IASF Palermo, Via U. La Malfa 153, I-90146
Palermo, Italy
7ISDC, Department of Astronomy, University of Geneva, Chemin d’Écogia 16, 1290
Versoix, Switzerland
8Institute of Space Sciences (ICE, CSIC), Campus UAB, Carrer de Can Magrans
s/n, E-08193 Barcelona, Spain
9Institut d’Estudis Espacials de Catalunya (IEEC), E-08034 Barcelona, Spain
E-mail<EMAIL_ADDRESS>
(Accepted XXX. Received YYY; in original form ZZZ)

###### Abstract

The Accreting Millisecond X-ray Pulsar IGR J17591–2342 is a LMXB system that
went in outburst on August 2018 and it was monitored by the NICER observatory
and partially by other facilities. We aim to study how the spectral emission
of this source evolved during the outburst, by exploiting the whole X-ray data
repository of simultaneous observations. The continuum emission of the
combined broad-band spectra is on average well described by an absorbed
Comptonisation component scattering black-body-distributed photons peaking at
(0.8$\pm$0.5) keV, by a moderately optically thick corona ($\tau$=2.3$\pm$0.5)
with temperature of (34$\pm$9) keV. A black-body component with temperature
and radial size of (0.8$\pm$0.2) keV and (3.3$\pm$1.5) km respectively is
required by some of the spectra and suggests that part of the central
emission, possibly a fraction of the neutron star surface, is not efficiently
scattered by the corona. The continuum at low energies is characterised by
significant residuals suggesting the presence of an absorption edge of O viii
and of emission lines of Ne ix ions. Moreover, broad Fe i and Fe xxv K$\alpha$
emission lines are detected at different times of the outburst, suggesting the
presence of reflection in the system.

###### keywords:

X-rays: binaries – stars: neutron – stars:individual:IGR J17591–2342 – line:
formation – line: profiles

††pubyear: 2015††pagerange: Spectral analysis of the AMXP IGR J17591–2342
during its 2018 outburst–A

##  1 Introduction

Accreting Millisecond X-ray Pulsars (AMXPs) are Low Mass X-ray binary systems
(LMXBs) hosting neutron stars (NS) that show coherent pulsations with periods
lower than 10 ms. Neutron Stars in these systems are characterised by low
magnetic fields, generally between $10^{8}$–$10^{9}$ G (see, e.g., Di Salvo &
Sanna, 2020). Their relatively small spin periods are now established to be a
direct consequence of the mass transfer occurring via Roche lobe overflow from
a low mass (< 1 M⊙) companion star onto a slow rotating NS. This property
makes AMXPs the progenitors of the rotation-powered millisecond pulsars
emitting on a large fraction of the electomagnetic spectrum, i.e. from the
radio to the gamma-ray band (Alpar et al., 1982). AMXPs usually show sporadic
outbursts during which the X-ray luminosity can attain values between
$10^{36}$ and $10^{37}$ erg/s. So far, 24 sources are included in this
subclass of objects (see Bult et al., 2022; Ng et al., 2021). Even though the
spectral evolution of these sources was rarely monitored in detail during the
whole outburst, AMXPs are generally observed in hard spectral states with no
hard to soft transitions. For this reason AMXPs are usually referred as hard
X-ray transients (Di Salvo & Sanna, 2020).

The AMXP system IGR J17591–2342 was discovered in outburst by the
INTernational Gamma-Ray Astrophysics Laboratory (INTEGRAL) on August 2018
(Ducci et al., 2018). Nowak et al. (2019) inferred the celestial coordinates
of the source from a pointed Chandra/HETG observation, i.e. 17h 59m 02.83s,
-23∘ 43’ 10.2" (J2000) with an error box of 0.6" representing the 90%
confidence level of Chandra positional accuracy. Observations performed by the
Australia Telescope Compact Array (ATCA) detected the radio counterpart of IGR
J17591–2342 providing an improvement on the previous coordinates that resulted
to be 17h 59m 02.86s $\pm$ 0.04, -23∘ 43’ 08.3" $\pm$ 0.1" (J2000) (Russell et
al., 2018).

Using NICER and NuSTAR observations, the timing analysis performed by Sanna et
al. (2018) allowed to constrain coherent pulsations at a period of 1.9 ms and
an orbital period of 8.8 hrs from the Doppler modulation. Moreover, the X-ray
data show a peculiar spin-down behaviour that is compatible with a
magnetically threaded disc (Sanna et al., 2020). The results reported by Sanna
et al. (2018) led to discuss a scenario in which the companion star is a late
spectral type star having an age between 8–12 Gyrs and with a mass between
0.85–0.92 M⊙, assuming a mass of 1.4 M⊙ for the NS. This scenario is
consistent with a value of the inclination angle between 28∘–30∘.

Broad-band spectral analysis was performed on this source, but considering an
average spectrum obtained by pointed observations at different times of the
outburst. Sanna et al. (2018) analysed a combined averaged spectrum consisting
in Swift, INTEGRAL/ISGRI and NuSTAR FPMA and FPMB data. They modelled the
spectrum using an absorbed soft black-body plus a Comptonisation component.
Their work reports a value of the hydrogen column density $N_{H}$ of
(3.6$\pm$1.1)$\times 10^{22}$ cm-2, and a Comptonisation component that is
characterised by a photon index $\Gamma$ of about 1.8, a temperature $kT_{e}$
of the corona of 22${}^{+4}_{-3}$ keV, and a temperature of the seed photons
(assuming a black-body spectrum) of $kT_{seed}=0.79\pm 0.09$ keV. The soft
black-body direct emission is characterised by a temperature that is
compatible with that of the seed photons and in line with the emission from a
region of a few kilometers. These authors also detected a weak Gaussian line
centred at an energy compatible with the iron K$\alpha$ complex, and
interpreted as a possible signature of disc reflection.

Nowak et al. (2019) modelled a 1–9 keV Chandra-NICER spectrum of the source,
taken at about 58353 MJD, using a model consisting in an absorbed black-body
and Comptonisation component. They inferred a value of $N_{H}$ of
(2.9$\pm$0.5)$\times 10^{22}$ cm-2 and a black-body temperature of 0.06 keV,
by fixing the $kT{{}_{e}}$ value to that observed by Sanna et al. (2018). In
addition, based on a Si xiii absorption line detected in the Chandra spectrum,
they propose the presence of an outflow with a velocity of about 2800 km s-1
in the system. They also found evidences of possible Ca lines in the HETGS
spectra and hypothesized that the NS could be formed via accretion induced
collapse of a white dwarf in a rare, calcium-rich Type Ib supernova explosion.

Kuiper et al. (2020) studied the broad-band 0.3–150 keV averaged spectrum of
the source using data from NuSTAR FPMA and FPMB modules, from XMM-Newton RGS,
Epic-pn and Epic-MOS2 instruments and from INTEGRAL/ISGRI. They modelled the
spectrum with an absorbed Comptonisation component finding a value of $N_{H}$
of (2.09$\pm$0.05)$\times 10^{22}$ cm-2 and a Comptonisation component in
which the temperature of the seed photons is about 0.64 keV, the coronal
temperature $kT_{e}$ is (38.8$\pm$1.2) keV and in which the optical depth of
the corona is $\tau=1.59\pm 0.04$. They excluded the presence of a local
absorber into the system, due to the fact that their estimation of $N_{H}$
results in line with the one expected according to the total Galactic
absorption (i.e. 2.2$\times 10^{22}$ cm-2) resulting from the optical
reddening maps (Russell et al., 2018). They also found evidences of an
emission line in the Iron K-$\alpha$ region. However, they considered the
detection of this line spurious and related to blending of lines from
different Fe ionisation stages, or alternatively to uncertainties in the XMM
Epic-pn response (Kuiper et al., 2020). In the same work, these authors
estimated an upper limit on the distance to the source of $d=(7.6\pm 0.7)$ kpc
according to the analysis of a burst showing clues of Photospheric Radius
Expansion (PRE) in one INTEGRAL/JEM-X observation.

In this work we perform a spectral study of IGR J17591–2342 based on the
entire NICER data set and including a large sample of the available
observations in the X-ray archive, with the aim of studying its spectral
evolution in detail during the whole outburst. The paper is structured as
follows: in section 2 we describe the data selection and reduction, in section
3 we report the data analysis, in section 4 we discuss the obtained results,
and in section 5 we summarize these results advancing some conclusions.

##  2 Observations and data reduction

Figure 1: NICER light curve of IGR J17591–2342 during the outburst of 2018.
Superimposed we show the mid-observation times at which other observatories
pointed the source and the shaded areas representing the time windows of the
INTEGRAL observations. Figure 2: The 3–6 keV (upper panel) and 6–10 keV
(middle panel) NICER light curve of IGR J17591–2342 during the outburst of
2018. In the bottom panel we report the hardness ratio evaluated between these
two energy bands.

A complete follow up of the outburst of IGR J17591–2342 was performed by the
X-ray timing instrument (XTI, Gendreau et al. 2012; Gendreau et al. 2016) on
board the Neutron star Interior Composition Explorer (NICER) mission, hosted
onto the International Space Station (ISS). This instrument collected data of
the source between 2018-08-14T23:59:42 (ObsID 1200310101) and
2018-10-17T05:32:22 (ObsID 1200310139). We reduced the NICER data using the
nicerl2111https://heasarc.gsfc.nasa.gov/docs/nicer/analysis_threads/nicerl2/
script available under HEAsoft version 6.28. We processed the data using the
latest gain and calibration database files of version 20200722. We applied
standard filtering criteria based on elevation angle from the Earth limb,
pointing offset, offset from the bright Earth, and South Atlantic Anomaly
region in our study. Following the above criteria, we selected the Good Time
Intervals using the nimaketime task. We extracted the source spectrum using
the filtered clean event file in the xselect environment, while we obtained
the background spectra corresponding to each observation using the
nibackgen3C50222https://heasarc.gsfc.nasa.gov/docs/nicer/toolsnicer_bkg_est_tools.html
tool (Remillard et al., 2021). For the spectral analysis, we considered the
response matrix and ancillary response file of version 20200722. We added a
systematic error of 1% to the obtained source spectra, as suggested by the
NICER team (see, e.g., Jaisawal et al., 2019).

From Figure 1, it is possible to follow the complete light curve of the source
during the outburst as observed by NICER in the band 0.2–10 keV. The first
observation of the source shows an average count-rate of about 17 c/s that
increased up to a first peak of 29 c/s at 58362 MJD. After this first maximum,
the count-rate decreased until it reached a value of about 17 c/s at 58367
MJD, when it started to increase again up to the absolute maximum of about 42
c/s reached at 58380 MJD. Later on, the source started to decrease its
luminosity down to a count-rate of 14 c/s, reached on 58397 MJD. In addition,
as highlighted by the hardness ratio reported in Figure 2 and evaluated
between the energy band 3–6 keV and 6–10 keV, the spectral state of IGR
J17591–2342 maintains quite constant, suggesting that no significant spectral
changes occurred during the outburst, in agreement with what usually observed
for AMXPs (Di Salvo & Sanna, 2020).

The outburst of IGR J17591–2342 was also observed by INTEGRAL between
revolutions 1986 and 2009. In addition, further single pointed observation of
the Nuclear Spectroscopic Telescope Array (NuSTAR), Chandra, XMM-Newton and
ASTROSAT were performed during the first part of the outburst.

NuSTAR (Harrison et al., 2013) observed IGR J17591–2342 on 2018-08-13T22:36:09
(ObsID 90401331002) and on 2018-08-17T20:01:09 (ObsID 80301311002). The data
have been reprocessed with the nupipeline routine while the source and
background spectra were extrapolated with the nuproducts pipeline, using
circular extraction regions of 60" of radius centred at the coordinates of the
source and far away from the source, respectively. The spectra were modelled
in the 2–80 keV energy band.

Chandra observed IGR J17591–2342 on 2018-08-23T17:40:05 for 20 ks using the
High Energy Transmission Grating spectrometer (HETG, Canizares et al., 2005).
The events were reprocessed using the official Chandra analysis environment
CIAO v. 4.13 with calibration files updated to the version 4.9.4. The data
were reprocessed with the chandra_repro pipeline and the first order High-
Energy Grating (HEG) and Medium-Energy Grating (MEG) spectra were combined
using the combine_spectra tool. The obtained spectra were grouped to have at
least 25 counts per energy bin and analysed in the 1.2–7.8 keV energy band.

The XMM-Newton observatory observed the source with both the Reflection
Grating Spectrometer (RGS, Den Herder et al., 2001) and the European Imaging
cameras (EPIC) on 2018-09-03T18:44:55 during a time window in which the count-
rate of the source momentarily decreased, as it can be inferred from Figure 1.
The science data reduction was performed using the XMM-Newton Science Analysis
System (SAS) version 19.0.0. The (imaging) EPIC-pn instrument (Strüder et al.,
2001) operated in Timing Mode, while the MOS-1 and MOS-2 cameras (Turner et
al., 2001), acquired data in Small window and Timing uncompressed modes,
respectively. The Epic-pn spectra were extracted from the RAWX coordinates
between 31–44, while the background spectrum was obtained from RAWX
coordinates centred between 5 and 15. The MOS 1 data were collected in Small
window mode and, as already noticed by Kuiper et al. (2020), are extremely
corrupted by pile-up effects and for this reason they are not considered in
our work. On the contrary, MOS 2 data are not affected by pile-up, since the
instrument was operating in Timing uncompressed mode, which allows to collect
data up to 35 mCrab without severe pile-up distortions. The source and
background spectra for the MOS2 were extrapolated in the RAWX range 290–320,
while the background spectrum was extracted by a box having a width and a
height of 8432.64 and 4285.44 in physical coordinates, respectively. The
analysed energy range is 2.2–10 keV. The RGS operated in the standard
Spectroscopy HighEventRate mode with Single Event Selection. No soft proton
background flares were detected during the observation. Then, we combined the
first order data of RGS1 and RGS2 using the rgscombine tool included in the
SAS package, and grouping the obtained spectrum to have at least 25 counts per
energy bin. However, the data counts are considerably low for these data for
energies below 1.2 keV, and for this reason we preferred not to consider these
data for our analysis, since the energy range above 1 keV is well covered by
the rest of the considered space missions.

The ASTROSAT mission observed the source with the Large Area X-ray
Proportional Counter (LAXPC) instrument (Yadav et al., 2016; Antia et al.,
2017; Agrawal et al., 2017) on 2018-08-23T01:10:15 (ObsID 9000002320) and on
2018-08-27T00:00:00 (ObsID 9000002332) for a net exposure of 30 ks and 37.7
ks, respectively. A thermonuclear type I X-ray burst occurred during the ObsID
9000002320 and has been removed before extracting the persistent spectrum of
the source. For both the observations, the LAXPC 30 was not working. On the
other hand, the LAXPC 10 was active on a low voltage gain
setting333http://astrosat-ssc.iucaa.in/. Since the source is faint, we
extracted spectra from the top layer of each detector to avoid unnecessary
background. However, the LAXPC 10 spectra appear to be affected by the
background and gain change, and for this reason they were not included in the
analysis. In addition, the LAXPC 20 spectrum extracted for ObsID 9000002320
shows a significant mismatch with respect to the spectra of the closest
available observations from other missions, possibly introduced by a bad
calibration. For this reason, this observation was also excluded from the
analysis. A systematic error of 2% has been added to the LAXPC 20 spectrum of
ObsID 9000002332 (see, e.g., Misra et al., 2017), which was also grouped to
have at least 25 counts per energy bin. The analysis was conducted in the 4–17
keV energy range.

INTEGRAL observed the outburst from the source during the period spanning from
58340 MJD to 58405 MJD, corresponding to satellite revolutions from 1986 to
2016. We considered data from the two JEM-X units (Lund et al., 2003) and from
IBIS/ISGRI (Ubertini et al., 2003; Laurent et al., 2003). We analysed the
ISGRI, JEM-X 1 and JEM-X2 data with the Offline Scientific Analysis (OSA)
software version 11.0 distributed by the ISDC (Courvoisier et al., 2003).
Different source spectra were initially extracted for the two JEM-X and
IBIS/ISGRI from each revolution. A grouping of 16 bins has been used for the
JEM-X (IBIS/ISGRI) spectra extraction, following the standard practice for
similar sources. We excluded the science window (SCW) 50 in revolution 2001
due to the presence of a type I X-ray burst (see also Kuiper et al., 2020).
The JEM-X spectra were analysed in the 3–20 keV energy range, while the
IBIS/ISGRI spectra were analysed in the 25–200 keV range.

Figure 3: Evolution of the main spectral parameters describing the continuum
emission of IGR J17591–2342 as a function of time. The associated errors are
reported at a level of statistical confidence of 90%. In the upper panel the
mean count-rate collected during each NICER observation is reported. In each
panel, the star-shaped points represent fixed values for the parameters.
Figure 4: Evolution of the iron line spectral parameters and statistical
significance of the observed feature. In the plot of the line centroids
(second plot from the top) we also reported the rest-frame energy of the Fe i
K$\alpha 1$ , Fe xxv (resonance), and Fe xxvi Ly$\alpha 1$ transitions, as
reference. In the bottom panel, we highlighted in pink the detection area
considering as lower limit $\sigma$=2.5 for a weak detection. The dotted line
represents the 3$\sigma$ detection threshold. Figure 5: Best fit model and
associated residuals obtained for the broad-band spectrum associated to the
NICER ObsID 1200310106 (see Table 2 and Table 3). In black the NICER spectrum,
in red, green and blue the INTEGRAL JEM-X1, JEM-X2 and ISGRI spectra of
revolution 1992 respectively, and in cyan the ASTROSAT LAXPC20 spectrum of
ObsID 9000002332.

##  3 Spectral Analysis

The spectral analysis was entirely performed using the X-ray spectral fitting
package Xspec v. 12.11.1c (Arnaud, 1996). We chose to fit as many spectra as
possible in order to follow the evolution of the spectral parameters during
the outburst with the highest possible temporal resolution guaranteed by the
available data. For this reason, we paired each NICER spectrum with any quasi-
simultaneaous, i.e. taken within two days, data-set from the other
observatories, when possible. As widely discussed in Sanna et al. (2018) and
Kuiper et al. (2020), after some preliminary tests on the available data, we
noted that the model which best describes the continuum was an absorbed
Comptonisation model. For this reason, we adopted this model in the following
analysis. Some of the performed fits, however, revealed that in some cases a
soft excess remained in the residuals and this could be corrected with the
addition of a black-body emission from the source, as already reported in
Sanna et al. (2018). We used the chemical abundances of Wilms et al. (2000a)
and the cross sections reported in Verner et al. (1996). We modelled the
continuum spectral emission of IGR J17591–2342 adopting the Tuebingen-Boulder
ISM absorption model Tbabs, a black-body component bbodyrad, and the thermal
Comptonisation component nthcomp. The nthcomp model is described by the
asymptotic power-law photon index $\Gamma$, the electron temperature $kT_{e}$
of the hot electron corona, the seed photon temperature $kT_{seed}$, the
$inp\\_type$ parameter, which can assume values 0 or 1 for considering black-
body or disc-black-body distributions for the seed photons, respectively, and
the $redshift$ parameter. We assumed that the seed photons are distributed
accordingly to a black-body law by fixing the $inp\\_type$ parameter to 0.
Moreover, we adopted a value of redshift equal to zero for all the subsequent
analysis. The NICER spectra were fitted in the energy range 0.6–10 keV. On the
other hand, the NuSTAR FPMA and FPMB spectra of ObsID 80301311002 were fitted
in the range 2–60 keV, while those of ObsID 90401331002 in the range 2–80 keV.
The XMM-Newton spectra of ObsID 0795750101 were fitted in the range 3–10 keV
and 2.2–10 keV for the EPIC-pn and EPIC-MOS2 spectra, respectively. The
Chandra MEG and HEG spectra of ObsID 20173 were fitted in the range 1.2–6.8
keV and 1.4–7.8 keV, respectively, while the ASTROSAT/LAXPC 20 spectrum of
ObsID 9000002332 was fitted in the energy range 4–17 keV.

We inspected with xspec all INTEGRAL data of the source, but for our broad-
band analysis we finally only made use of the spectra extracted during
revolutions from 1992 and 2006 because they are more reasonably close in time
to the available NICER data. During revolutions from 1986 to 1989, the source
was caught by INTEGRAL during the earliest stages of the outburst evolution
and the INTEGRAL data were characterised by a number of counts too low to
perform any meaningful spectral analysis. A similar conclusion applies to the
data collected during revolutions 2008–2016, corresponding to the rapid decay
of the source flux toward the end of the outburst. Depending on the source
flux, we limited our spectral analysis for the JEM-X data roughly in the
interval 3–20 keV. During revolutions 1989, 1994, 1995, 2002, 2004 and 2006,
the JEM-X data did not have a large number of counts, in order to be used in
the spectral analysis and thus we do not mention these data any longer in the
following sections. IBIS/ISGRI data were used, following the OSA 11.0
recommendations, from an energy of 25 keV up to roughly 200 keV, above which
the source fell below the detection threshold of the instrument. The higher
energy bound of the IBIS/ISGRI spectrum varies across the different
revolutions considered depending on the flux of the source but remains in all
cases well above 100 keV.

The NICER spectra, in combination with those of other missions, generally
ensured a good energy coverage over a large energy range. However, some of the
NICER spectra could not be combined with other spectra, keeping hard to
constrain the high-energy cut-off $kT_{e}$ for the Comptonisation component.
In all these cases, this parameter was fixed to the value obtained by Kuiper
et al. (2020) (i.e. $kT_{e}=38.8\;keV$). In Figure 3 we show the evolution of
the spectral parameters obtained from the best fit model of each observation
as a function of time. In all the plots, the star-shaped points indicate
parameters that were frozen in the fit. This plot shows how the $kT_{e}$
parameter has been constrained for many observations, with an average value
obtained during the whole outburst of about 34 keV, which is in line with that
inferred by Kuiper et al. (2020). Some of the NICER spectra showed the
presence of residuals in absorption at about 0.87 keV, which are consistent
with the presence of an O viii absorption edge. To fit the low energy
residuals, as a first step, we tried to replace the Tbabs ISM absorption
component with the Tbfeo and the Tbvarabs components that take into account
variable abundances for the chemical species in the ISM. However, even if both
these models return perfectly compatible values for the $N_{H}$, they are
unable to fit the observed residuals. For this reason we manually described
the observed feature by recurring to the edge component in Xspec. The energy
of the edge was kept frozen for all the fits. Other localised residuals
characterise some of the NICER spectra at about 0.922 keV, in line with the
presence of Ne ix ions. This feature was modelled with a Gaussian component in
which the energy centroid was fixed to the rest-frame energy of the
aforementioned ion. Moreover, the spectral resolution of the NICER/XTI
instrument around 1 keV was not sufficient to constrain the width of the
modelled line. For this reason, we fixed the parameter $sigma$ to 0.085 keV
(i.e. the NICER/XTI spectral resolution at 1 keV). On the other hand, the
normalisation of the Gaussian component was left free to vary. Firstly, the
neutral column density NH in the Tbabs component was left free to vary during
the whole duration of the outburst, but we noticed a correlation with the
kTseed parameter of the Comptonisation component, with a consequent scattering
of the NH component. Therefore, we decided to keep this parameter frozen at
the value of (2.09$\pm$0.05)$\times 10^{22}$ cm-2, in accordance with Kuiper
et al. (2020) for the rest of the analysis.

The Comptonisation component is on average characterised by stable values of
$\Gamma$ and $kT_{e}$, equal to 1.9, with a standard deviation of about 0.2,
and 34 keV with a standard deviation of about 9 keV, respectively. On the
other hand, during the outburst the temperature of the seed photons, assumed
to be injected with a black-body-like spectrum, shows a variation in
accordance with the advance of the outburst. The mean value of this parameter
during the outburst was of 0.8 keV with an associated standard deviation of
0.5 keV.

The bbodyrad component appears to be required by the fit only in some of the
observations. We plotted the resulting black-body temperature $kT$ in the
fifth panel from the top of Figure 3. This parameter shows an almost constant
trend during the outburst with a mean value of 0.8 keV and a standard
deviation of about 0.2 keV. The associated errors at 90% of confidence level
strongly suffer from the statistics of the fitted spectra. However, for all
the reported black-body components the significance of its detection results
to be higher than 3$\sigma$ (the magenta points have a significance of more
than 5$\sigma$, whilst the dark violet ones have a significance between
3$\sigma$ and 4$\sigma$).

Some of the observed combined spectra also showed significant residuals in the
energy range at which the Fe xxv line is expected, i.e at about 6.5 keV. Once
the residuals have been detected, we tried to model them with a Gaussian
component. For each of these observations we evaluated the significance of
this local feature, by evaluating how much the normalization of this component
deviates from the continuum in units of $\sigma$. In the lower panel of Figure
4, we report the significance of the observed iron line. In the upper and
medium panels of the same figure, we show how the parameters of this feature
evolve in time during the outburst. In particular, we can observe that, with
the exception of the second point for which we have an upper limit on the
energy centroid, the rest of the lines have been well constrained, even though
it was not possible to obtain lower limits for the associated width. However,
these observations suggest that the line is marginally detected and that
results to be in agreement with a Fe i K$\alpha 1$ line the first phase of the
burst. After this phase, the line starts to be consistent with the presence of
more ionised species as Fe xxv. Since this line tends to be generally broad,
when its width could be constrained, we tried to fit this line also by using
the Relativistic model discline (Fabian et al., 1989) or the most recent
shaddisc (La Placa et al., 2020). These models, however, with the exception of
the centroid energy, returned totally unconstrained spectral parameters and
generally not physically reasonable due to the low number of counts on the
line profile.

As an example, we show the best fit model obtained for the broad-band spectrum
associated to the NICER ObsID 1200310106 in Figure 5.

##  4 Discussion

###  4.1 Spectral evolution

We report hereon the spectral analysis of the whole set of NICER observations
of the outburst of IGR J17591–2342 occurred in 2018, integrating these
observations with all the available quasi-simultaneous observations stored in
the X-ray data archive. The NICER light curve during the outburst is
characterised by a double peak. The spectral model adopted to fit the spectra
of the source suggests that during the first peak, occurring at 58362 MJD, the
source showed an unabsorbed bolometric flux of (4.6$\pm$0.2)$\times 10^{-10}$
erg cm-2 s-1 in the range 0.1–100 keV, which increases up to
7.2${}^{+0.7}_{-0.5}\times 10^{-10}$ erg cm-2 s-1 in the same energy band at
58380 MJD, when the second peak of the outburst occurred. The flux errors are
reported at 3$\sigma$ c.l., as derived from the cflux convolution model.

Comptonisation is the physical process for which we observe the major
contribution in terms of flux. Actually, this component, on average,
contributes to the total unabsorbed flux for 95% in the case of IGR
J17591–2342. According to the obtained results, the Comptonisation component
appears to be characterised by an electron corona with a temperature $kT_{e}$
of about 34 keV, which is in line with the value of 38.8 keV reported by
Kuiper et al. (2020). Unfortunately, due to the lack of coverage at the higher
energies in some points of the outburst, we were unable to constrain this
parameter for all the observations performed by NICER. However, we constrained
this parameter in 29% of the NICER observations, inferring that at the
beginning of the outburst the cloud was characterised by a slightly lower
temperature of about 24 keV.

The photons that are Comptonised in this corona appear to be mainly with a
temperature of about 0.8 keV, with the tendency to slightly increase during
the peaks of the outburst, always, however, at the limit of statistical
compatibility at the 90% confidence level. This could be due to episodes of
increased heating of the neutron star surface/boundary layer, plausibly caused
by occasional rises in the mass-accretion rate. The asymptotic power law index
$\Gamma$, on the contrary, remains almost stable between 1.7 and 2.1.

The black-body component has been tested for each spectrum and only 26% of the
best fit models require this component at a level of confidence that is
greater than 3$\sigma$. The equivalent radius of emission of the black-body
component has been obtained from the normalization as Rbb=$N_{bb}d^{2}_{10}$,
where Nbb is the normalization of the bbodyrad component and $d^{2}_{10}$ is
the distance to the source in units of 10 kpc. Assuming the distance to the
source obtained in this work (i.e., (7.2$\pm$0.8) kpc), we obtained the radii
reported in Figure 6.

Figure 6: Trend of the equivalent radius of emission of the black-body
component (upper panel) and of the Comptonisation seed photons (middle panel)
as a function of time. In the lower panel, the trend followed by the optical
depth of the Comptonising electron cloud is also reported. The arrows indicate
upper limits on the values of the parameters.

From the observed trend, it is possible to notice how the radius of the
emitting black-body is mainly at a value of about $\overline{R_{bb}}$ =
(3.3$\pm$1.5) km, where the associated error in this case coincides with the
standard deviation of the distribution of measurements. This component,
peaking mainly at $\overline{kT}$ = (0.8$\pm$0.2) keV, could arise from a
direct emission by a fraction of the NS surface. It is interesting to notice
that the size of this region varies in accordance with the variation of the
flux extrapolated in the band 0.1–100 keV during the outburst. In particular,
for lower values of the flux, it appears smaller with respect to the case in
which the flux is higher. One possible explanation could be represented by the
fact that when the mass accretion rate increases, the region responsible for
the emission becomes larger, reaching a maximum in proximity of the two peaks
of the outburst, where the values are consistent with a radius of about 5 km,
i.e half of the NS size, assuming a NS of 10 km of radius. Then, we tried to
test if the direct observation of the black-body emission from the NS might be
related to a low optical depth of the hot corona, using the relation of
Zdziarski et al. (1996):

$\Gamma=\left[\dfrac{9}{4}+\dfrac{1}{\tau\left(1+\dfrac{\tau}{3}\right)\left(\dfrac{kT_{e}}{m_{e}c^{2}}\right)}\right]^{1/2}-\dfrac{1}{2},$
(1)

with the spectral parameters obtained in the best fit model for each
observation. The evolution of this parameter as a function of the observing
time is reported in Figure 6 in the lower panel. As shown by this trend, the
value of the optical depth remains almost stable during the outburst at a
value of about $\tau\sim$2.3 and a standard deviation of 0.5, at least for the
parameters that were constrained. For this reason, it is possible that the
corona Comptonises the majority of the photons emitted by the NS surface
leaving, however, only a small fraction ($\sim$ 10%) of them not significantly
Comptonised. The latter could contribute to the observed direct black-body
component, peaking approximately at the same temperature of the seed photons,
as evidenced in the obtained results.

A direct proof of this scenario could be provided by estimating the equivalent
radius of the region emitting the seed photons for the Comptonisation, which
we expect more or less similar to the radius obtained for the direct black-
body emission region. An estimation of this physical parameter can be obtained
by using the relation of In ’t Zand et al. (1999), assuming a spherical
geometry of the corona,

$R_{0}=3\times
10^{4}\;d\left(\dfrac{f_{bol}}{1+y}\right)^{1/2}\left(kT_{seed}\right)^{-2},$
(2)

where $d$ is the distance to the source in kpc inferred in this work,
$f_{bol}$ is the unabsorbed bolometric flux extrapolated from the
Comptonisation component in erg cm-2 s-1, $kT_{seed}$ is the temperature of
the seed photons in keV, and $y=4kT_{e}max[\tau,\tau^{2}]/(m_{e}c^{2})$ is the
Compton parameter, in which $kT_{e}$ is the electron temperature in keV. The
radius $R_{0}$ obtained for each observation is plotted as a function of the
observing time in the middle panel of Figure 6. These parameters, at least in
the cases in which the statistics of the data allowed to constrain them, are
scattered around a mean value of $\overline{R_{0}}$ = (5.0$\pm$2.8) km, which
is in agreement with the obtained value for $\overline{R_{bb}}$.

###  4.2 Spectral lines

The spectral continuum of the source is characterised by the presence of
several local features. On the one hand, almost the totality of the analysed
spectra present an absorption edge at 0.871 keV, which suggests the presence
of O viii ions. Despite the fact that the energy of this feature needed to be
fixed due to the complexity of the residuals at the lowest energies in the
NICER spectra, the significance of this feature is always higher than
3$\sigma$ for all the results reported in Table 2. The low energy residuals of
some spectra are well fitted by taking into account a Ne ix emission line for
all the observations (see Table 3). In accordance with the statistics of the
available spectra and with the spectral resolution of NICER/XTI, which is not
sufficient to constrain the line width, we fixed the parameter $sigma$ of the
Gaussian component to a value of 85 eV, which is the spectral resolution of
the instrument at 1 keV. Trying to leave this parameter free to vary during
the fit, indeed results in a value of width that is of the order of the
spectral resolution of NICER at 1 keV, or slightly higher, even though not
constrained. This could probably suggest two possible scenarios: on one hand,
the lines could be produced approximately in the same region of an accretion
disc where they are relativistically broadened due to the proximity to the NS.
This scenario, however, is not supported by evidences of any further black-
body or multicolour-disc black-body component needed to fit the spectrum and
attributable to emission from the disc. On the other hand, this line could be
smeared only as a consequence of Compton broadening, or as an effect of the
not sufficient spectral resolution of the XTI instrument to resolve a complex
of lines at those energies. The test of this latter scenario deserves
observations with high spectral resolution instruments. Indeed, the low energy
spectrum of the source is completely dominated by the background flux below 1
keV in the case of the pointed observations performed by the Chandra/HETG and
the two MOS instruments on board the XMM-Newton mission.

About 19% of the observations show a significant evidence ($>3\sigma$) of Fe
emission lines. In Figure 4, we show the evolution followed by the parameters
describing the line profile, assumed to be Gaussian. The detection of this
line is significant only for the time window including the range spanned by
the first peak of the outburst and by part of the rise of flux towards the
main peak of the burst, i.e during the periods of higher statistics of the
NICER spectra. Our analysis shows how, in the reported cases, it is possible
to constrain the centroid energy of the lines, with the exception of the
observations occurred on 58357 MJD and on 58395 MJD, for which we find only an
upper limit of 6.3 and 6.7 keV, respectively. The remaining sample of line
profiles result to be in line with the rest-frame energy corresponding to the
Fe xxv K$\alpha$ line, usually observed in bright LMXB systems showing
evidences of reflection. The observed sigma for these lines results to be on
average of about 0.5 keV and its distribution shows a standard deviation of
about 0.5 keV. On the basis of this result, it seems that the line profile is
tendentially broad. Kuiper et al. (2020) detected a similar line profile in
the average spectrum of the source. The feature they observed is centred at
about 6.7 keV and has a width of 0.69 keV, in line with our results. However,
they consider this feature to be characterised by values of the centroid
energy and sigma that are too far off and too broad. They propose that this
feature could be produced by a blend of Fe lines deriving from different
ionisation stages, which can not be distinguished individually by the MOS due
to its spectral resolution. Moreover, they propose that this detection could
reflect uncertainties in the XMM-Newton EPIC-pn response for observations
taken in timing mode, also because there is no detection for Chandra/HETG
data, having a higher spectral resolution. For this reason they do not
consider the detection as real. On the contrary, we detected this line at
different times during the outburst, including the case occurring on 58353 MJD
for which we observe a detection at about 3$\sigma$, and that was obtained by
considering a combined spectrum that also includes a Chandra/HETG spectrum. In
accordance with this evidence, and to the fact that also other spectra during
the outburst show the same feature, we conclude that this feature is real.

The poor statistics offered by the available data on the iron line profile,
and the lack of detected lines for close observations, does not allow a deeper
investigation on the nature of this feature that could be produced by
reflection. A simple modelling with a Relativistic component as discline or
shaddisc returns unconstrained values for the physical parameters describing
the line profile. Further observations of the source in outburst, performed
with future space missions that will be provided with higher effective area
and spectral resolution over a wider energy range, will be fundamental to
investigate the reflection component in IGR J17591–2342, for a deeper
comprehension of the ionisation state of the matter and of the geometry of the
system.

###  4.3 A new measurement of the distance

The analysis of PRE type I X-ray bursts is largely used in literature for
inferring the distance to the sources, the method being often the only
possible way to obtain such an estimate. However, this method suffers from
systematic uncertainties (see, e.g., Marino et al., 2019, and references
therein). For instance, the flux reached by PRE bursts for the same source
generally scatters around a mean value with variations of about 15% (see,
e.g., Kuulkers et al., 2003; Galloway et al., 2003, 2008). Moreover, the value
used for the Eddington luminosity of the source may not be accurate without
details on, e.g., the composition of the NS atmosphere and and therefore of
its opacity. The latest value for the distance of the source is $d=(7.6\pm
0.7)$ kpc, obtained by Kuiper et al. (2020) from the analysis of a type-I
X-ray burst which was suggested to reach the Eddington limit. However, without
strong evidences of PRE, the obtained distance could be an upper limit rather
than an actual estimate. Furthermore, the aforementioned systematic
uncertainties that affect the method, demand to try an alternative way to find
the distance and eventually confirm the measurement by Kuiper et al. (2020).

From the value of the hydrogen column density, we can derive an estimate of
the distance of the system to compare with the existing one by Kuiper et al.
(2020), by invoking the 3D extinction map of the radiation in the Ks band for
our Galaxy of Chen, B. Q. et al. (2013). The map provides a profile of the
radiation extinction in the direction of the Galactic bulge, as a function of
the distance to the source. We used the profile at galactic coordinates
$l=0.00$, $b=1.00$, reported in Figure 7.

Figure 7: Expected profile of the extinction in the Ks band, as a function of
the distance to the source in the direction of IGR J17591–2342. The blue
vertical lines indicates the best fit value (navy blue line) of the distance
to the source inferred by Kuiper et al. (2020) and its relative error (lighter
blue lines). The green line represents the best fit parabolic function that
fits the profile in the range 5–10 kpc.

The value of NH is related to the visual extinction of the source radiation AV
through the relation of Güver & Özel (2009):

$N_{H}=(2.21\pm 0.09)\times 10^{21}A_{V}.$ (3)

The visual extinction is then related to the extinction of the radiation in
the Ks band (A${}_{K_{S}}$) through the relation of Nishiyama et al. (2008):

$A_{K_{S}}=(0.062\pm 0.005)\;A_{V}\;{\rm mag}.$ (4)

We fitted the profile of Figure 7 with a parabolic function in the range 5–10
kpc, i.e. in the region corresponding to values of the distance inferred by
Kuiper et al. (2020), that is d = (7.6$\pm$0.7) kpc. Assuming the value of NH
found by Kuiper et al. (2020), the value of expected extinction is
A${}_{K_{s}}$=0.59$\pm$0.01 mag, corresponding to a distance of $d=(7.2\pm
0.8)$ kpc.

The estimation of the distance through the hydrogen column density is strongly
affected by uncertainties on this parameter. The accuracy in the evaluation of
the column density depends on the quality of the data and on the energy range
in which the fit is conducted. A small energy range and a low resolution can
lead to correlation effects with other parameters of the model. The use of
solar abundances as initial values for the parameters can also lead to
differences around 5% in the estimation (see Wilms et al., 2000b). Moreover,
the accuracy extinction map itself depends on the quality of the data used,
whether they are recent data with higher quality, and the analysed region of
the sky, which can be affected by strong intrinsic variations in the
interstellar medium. Nonetheless, the two estimates of the distance, by means
of the PRE analysis of Kuiper et al. (2020) and through the spectral evolution
of this work, are compatible.

##  5 Conclusions

We analysed a large sample of the available observations of IGR J17591–2342 in
the X-ray archive with the aim of characterising the spectral emission of the
source and its evolution during the outburst. The source is well fitted by an
absorbed Comptonisation component that on average contributes to 95% of the
whole budget of the unabsorbed emitted flux. No significant spectral changes
were found for the source during the whole outburst. The estimation of the
distance is in line with the same values reported in literature, and equal to
$d=(7.2\pm 0.8)$ kpc.

The spectral continuum, especially in the time window at which the flux
reaches higher values during the outburst, needs to be fitted with a black-
body component peaking at a temperature that is correlated with the NICER
count-rate. This component is characterised by a temperature of about 0.8 keV
and appears to be emitted from a region with a radius of (3.3$\pm$1.5) km that
could be compatible with a fraction of the NS surface, or possibly with the
boundary layer. No significant variations are observed on the electron
temperature of the Comptonising cloud that mainly shows a temperature of about
34 keV. The corona appears to be characterised by an optical depth of about
2.3, which could in part explain the direct black-body component that could
arise from a fraction of photons emitted by a small region of the NS surface
and that are not significantly scattered in the corona. The spectral continuum
appears to be characterised by a Ne ix emission line. This line, however,
seems to have a broad profile that is compatible or some times higher than the
spectral resolution of NICER at 1 keV. This could be produced by relativistic
effects, if this feature is originated in the innermost parts of the accretion
disc (the evidence of which is not detectable from the analysed data), or
could be an effect of the spectral resolution of NICER, being unable to
resolve a complex of lines at those energies.

A broad iron emission line has been detected in the spectra of a number of
observations. The energy line appears to be correlated to the phase of the
outburst. In particular, for lower flux regimes, we observe a broad ($\sim 1$
keV) Fe i K$\alpha$ line, while in proximity of the peak of the outburst we
observed a Fe xxv K$\alpha$ line. Observations of future outbursts of IGR
J17591–2342 with instruments equipped with larger effective area over a wider
energy range, as for example the enhanced X-ray Timing and Polarimetry mission
(eXTP, Zhang et al., 2019; In’t Zand et al., 2019), could provide important
constraints on the possible reflection component in this system, and then also
on the ionization state of the matter and on the inclination angle of the
system that can not be inferred by the current statistics.

## Acknowledgements

The authors acknowledge financial contribution from the agreement ASI-INAF
n.2017-14-H.0 from INAF mainstream (PI: A. De Rosa), and from the HERMES
project financed by the Italian Space Agency (ASI) Agreement n. 2016/13 U.O
and from the ASI-INAF Accordo Attuativo HERMES Technologic Pathfinder n.
2018-10-H.1-2020. We also acknowledge support from the European Union Horizon
2020 Research and Innovation Framework Programme under grant agreement HERMES-
Scientific Pathfinder n. 821896. RI and TDS acknowledge the research grant
iPeska (PI: Andrea Possenti) funded under the INAF national call Prin-SKA/CTA
approved with the Presidential Decree 70/2016. RI acknowledges financial
contribution from the agreement ASI-INAF n.2017-14-H.0, from INAF mainstream
(PI: T. Belloni). A. Marino is supported by the H2020 ERC Consolidator Grant
"MAGNESIA" under grant agreement No. 817661 (PI: Rea) and National Spanish
grant PGC2018-095512-BI00. This work was also partially supported by the
program Unidad de Excelencia Maria de Maeztu CEX2020-001058-M, and by the
PHAROS COST Action (No. CA16214).

## Data availability

The data utilized in this article are publicly available in the Heasarc Data
Archive at https://heasarc.gsfc.nasa.gov/cgi-bin/W3Browse/w3browse.pl. The
ASTROSAT data are publicly available in the ISRO Science Data Archive at
https://webapps.issdc.gov.in/astro_archive/archive/Search.jsp.

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##  Appendix A Tables

Table 1: Observations considered in this work. |  |  |
---|---|---|---
ObsID | Satellite | Start Time (UT) | Stop Time (UT)
90401331002 | NuSTAR | 2018-08-13T22:36:09 | 2018-08-14T14:26:09
1200310101 | NICER | 2018-08-14T23:59:42 | 2018-08-15T14:08:00
80301311002 | NuSTAR | 2018-08-17T20:01:09 | 2018-08-18T13:21:09
Rev. 1989 | INTEGRAL | 2018-8-17 10:45:21 | 2018-8-19 19:44:37
1200310102 | NICER | 2018-08-18T02:07:53 | 2018-08-18T03:56:20
1200310103 | NICER | 2018-08-23T00:58:20 | 2018-08-23T23:01:36
9000002320 | AstroSAT | 2018-08-23T01:10:15 | 2018-08-24T00:41:18
20173 | Chandra | 2018-08-23T17:40:05 | 2018-08-23T23:31:42
1200310104 | NICER | 2018-08-24T01:51:10 | 2018-08-24T14:39:23
Rev. 1992 | INTEGRAL | 2018-8-25 13:08:59 | 2018-8-27 19:11:30
1200310105 | NICER | 2018-08-27T00:49:37 | 2018-08-27T22:39:23
9000002332 | AstroSAT | 2018-08-27T00:05:56 | 2018-08-27T21:59:56
1200310106 | NICER | 2018-08-27T23:55:20 | 2018-08-28T01:44:44
1200310107 | NICER | 2018-08-30T15:18:00 | 2018-08-30T17:35:00
Rev. 1994 | INTEGRAL | 2018-8-30 20:48:49 | 2018-9-2 3:31:23
1200310108 | NICER | 2018-08-31T22:13:38 | 2018-08-31T22:55:20
1200310109 | NICER | 2018-08-31T23:47:15 | 2018-09-01T17:05:18
1200310110 | NICER | 2018-09-02T06:40:10 | 2018-09-02T23:58:36
Rev. 1995 | INTEGRAL | 2018-9-2 12:39:01 | 2018-9-4 19:22:19
795750101 | XMM-Newton | 2018-09-03T18:47:34.08 | 2018-09-04T03:58:54.86
1200310111 | NICER | 2018-09-03T01:12:30 | 2018-09-03T07:41:56
1200310112 | NICER | 2018-09-05T17:59:40 | 2018-09-05T21:55:00
1200310113 | NICER | 2018-09-06T00:24:00 | 2018-09-06T06:49:00
1200310114 | NICER | 2018-09-07T08:42:40 | 2018-09-07T09:26:40
1200310115 | NICER | 2018-09-08T01:42:11 | 2018-09-08T11:42:51
1200310116 | NICER | 2018-09-09T03:58:11 | 2018-09-09T22:59:33
Rev. 1998 | INTEGRAL | 2018-9-10 12:07:30 | 2018-9-12 18:48:26
1200310118 | NICER | 2018-09-11T00:44:36 | 2018-09-11T17:58:50
1200310119 | NICER | 2018-09-11T23:54:45 | 2018-09-12T21:49:14
1200310120 | NICER | 2018-09-13T05:15:23 | 2018-09-13T22:53:29
1200310121 | NICER | 2018-09-14T07:57:04 | 2018-09-14T17:04:00
1200310122 | NICER | 2018-09-15T00:54:21 | 2018-09-15T19:31:40
Rev. 2001 | INTEGRAL | 2018-9-18 11:33:28 | 2018-9-20 18:17:35
1200310124 | NICER | 2018-09-19T06:26:49 | 2018-09-19T22:06:40
1200310125 | NICER | 2018-09-20T05:35:18 | 2018-09-20T15:09:00
Rev. 2002 | INTEGRAL | 2018-9-21 4:47:37 | 2018-9-23 10:06:50
1200310126 | NICER | 2018-09-21T00:07:18 | 2018-09-21T11:14:40
1200310127 | NICER | 2018-09-23T07:45:04 | 2018-09-23T11:11:45
Rev. 2003 | INTEGRAL | 2018-9-23 19:12:26 | 2018-9-26 1:57:06
1200310128 | NICER | 2018-09-24T02:17:04 | 2018-09-24T02:41:20
1200310129 | NICER | 2018-09-25T01:27:04 | 2018-09-25T15:49:20
1200310130 | NICER | 2018-09-26T03:55:00 | 2018-09-26T04:11:20
Rev. 2004 | INTEGRAL | 2018-9-26 11:02:00 | 2018-9-28 17:47:47
1200310131 | NICER | 2018-09-28T00:30:47 | 2018-09-28T05:38:00
Rev. 2005 | INTEGRAL | 2018-9-29 4:14:04 | 2018-10-1 9:38:30
Rev. 2006 | INTEGRAL | 2018-10-1 18:42:44 | 2018-10-4 0:28:03
1200310132 | NICER | 2018-10-04T01:52:19 | 2018-10-04T22:15:00
1200310133 | NICER | 2018-10-05T18:10:40 | 2018-10-05T21:19:40
1200310134 | NICER | 2018-10-10T20:33:44 | 2018-10-10T20:48:40
1200310135 | NICER | 2018-10-11T01:11:02 | 2018-10-11T01:28:20
1200310136 | NICER | 2018-10-12T06:32:14 | 2018-10-12T06:56:18
1200310137 | NICER | 2018-10-13T13:24:10 | 2018-10-13T15:15:40
1200310138 | NICER | 2018-10-16T03:14:20 | 2018-10-16T05:06:20
1200310139 | NICER | 2018-10-17T05:32:22 | 2018-10-17T18:13:20
Table 2: Best fit model continuum for each observation of IGR J17591–2342. The associated errors are reported at 90% confidence level. |  | Tbabs | Edge | Bbodyrad | nthComp | $\chi^{2}/dof$
---|---|---|---|---|---|---
ObsID | MJD |  NH | EdgeE | MaxTau | kT | Norm | Gamma | $kT_{e}$ |  kTseed | Norm |
|  |  ($\times 10^{22}$) | (keV) |  | (keV) |  |  | (keV) | (keV) |  |
1200310101 | 58345.0 | 2.09* | $0.871*$ | $2.0\pm 0.2$ | $1.10\pm 0.06$ | $0.8\pm 0.2$ | $1.76\pm 0.02$ | $22^{+4}_{-3}$ | $0.43\pm 0.03$ | $0.0153^{+0.0010}_{-0.0009}$ | $2444.3/2147$
1200310102 | 58348.09 | 2.09* | $0.871*$ | $2.5^{+0.6}_{-0.2}$ | – | – | $1.77\pm 0.02$ | $22^{+8}_{-4}$ | $1.03^{+0.06}_{-0.07}$ | $0.0019^{+0.0003}_{-0.0002}$ | $1541.3/1515$
1200310103 | 58353.04 | 2.09* | $0.871*$ | $2.0\pm 0.1$ | – | – | 1.9* | 38.8* | $0.58\pm 0.01$ | $0.0150\pm 0.0004$ | $1507.4/1641$
1200310104 | 58354.08 | 2.09* | $0.871*$ | $1.7\pm 0.2$ | – | – | $1.91\pm 0.05$ | 38.8* | $0.63\pm 0.02$ | $0.0139^{+0.0007}_{-0.0006}$ | $1346.2/1563$
1200310105 | 58357.03 | 2.09* | $0.871*$ | $2.2\pm 0.3$ | $1.02^{+0.08}_{-0.09}$ | $4.4^{+2.1}_{-1.6}$ | $1.65^{+0.07}_{-0.07}$ | 38.8* | $0.44^{+0.06}_{-0.07}$ | $0.016\pm 0.002$ | $760.8/718$
1200310106 | 58358.0 | 2.09* | $0.871*$ | $2.5^{+0.3}_{-0.7}$ | $1.15^{+0.09}_{-0.07}$ | $3\pm 1$ | $1.71^{+0.07}_{-0.06}$ | $40^{+57}_{-13}$ | $0.42^{+0.10}_{-0.13}$ | $0.019^{+0.009}_{-0.004}$ | $538.0/557$
1200310107 | 58360.64 | 2.09* | – | – | – | – | $1.82\pm 0.05$ | $31^{+55}_{-10}$ | $0.54\pm 0.02$ | $0.012\pm 0.002$ | $657.4/623$
1200310108 | 58361.93 | 2.09* | $0.871*$ | $2.1\pm 0.3$ | – | – | $1.84^{+0.07}_{-0.06}$ | $32^{+86}_{-11}$ | $0.55\pm 0.04$ | $0.022\pm 0.002$ | $587.1/540$
1200310109 | 58361.99 | 2.09* | $0.871*$ | $1.8\pm 0.2$ | – | – | $1.90\pm 0.04$ | $\geq 42$ | $0.60\pm 0.02$ | $0.0207^{+0.0009}_{-0.0008}$ | $798.7/769$
1200310110 | 58363.28 | 2.09* | $0.871*$ | $1.9\pm 0.2$ | $0.58\pm^{+0.04}_{-0.05}$ | $58^{+10}_{-8}$ | $2.3^{+0.3}_{-0.2}$ | 38.8* | $1.1\pm 0.2$ | $0.0061^{+0.0028}_{-0.0015}$ | $865.2/745$
1200310111 | 58364.05 | 2.09* | $0.871*$ | $1.6\pm 0.2$ | – | – | $1.79\pm 0.02$ | 38.8* | $0.57\pm 0.02$ | $0.0164\pm 0.0008$ | $1190.3/840$
1200310112 | 58366.75 | 2.09* | $0.871*$ | $1.4\pm 0.6$ | $0.61^{+0.09}_{-0.10}$ | $36^{+21}_{-8}$ | 1.9* | 38.8* | $1.3^{+0.5}_{-0.4}$ | $\leq 0.0023$ | $533.0/489$
1200310113 | 58367.02 | 2.09* | $0.871*$ | $2.4\pm 0.4$ | – | – | $1.78^{+0.06}_{-0.05}$ | 38.8* | $0.45\pm 0.05$ | $0.017\pm 0.002$ | $561.6/535$
1200310114 | 58368.36 | 2.09* | $0.871*$ | $1.4\pm 0.3$ | – | – | 1.9* | 38.8* | $0.60\pm 0.03$ | $0.0132^{+0.0010}_{-0.0009}$ | $492.6/496$
1200310115 | 58369.07 | 2.09* | $0.871*$ | $1.9\pm 0.2$ | – | – | $1.87^{+0.06}_{-0.05}$ | 38.8* | $0.57\pm 0.03$ | $0.0148^{+0.0010}_{-0.0009}$ | $683.9/643$
1200310116 | 58370.17 | 2.09* | $0.871*$ | $2.0\pm 0.2$ | – | – | $1.93^{+0.06}_{-0.05}$ | $\geq 41$ | $0.62\pm 0.03$ | $0.0168^{+0.0013}_{-0.0009}$ | $735.3/718$
1200310118 | 58372.03 | 2.09* | $0.871*$ | $1.3\pm 0.3$ | $0.73^{+0.08}_{-0.09}$ | $45^{+8}_{-40}$ | $1.9\pm 0.2$ | 38.8* | $1.8^{+0.8}_{-0.5}$ | $0.005^{+0.014}_{-0.002}$ | $530.0/499$
1200310119 | 58373.0 | 2.09* | $0.871*$ | $2.0\pm 0.3$ | $0.65^{+0.06}_{-0.10}$ | $45^{+8}_{-40}$ | $1.8^{+0.5}_{-0.4}$ | $\geq 27$ | $1.2^{+0.4}_{-0.7}$ | $0.005^{+0.014}_{-0.002}$ | $719.1/648$
1200310120 | 58374.22 | 2.09* | $0.871*$ | $2.2\pm 0.3$ | – | – | $1.84\pm 0.06$ | 38.8* | $0.62\pm 0.03$ | $0.021\pm 0.001$ | $641.6/642$
1200310121 | 58375.33 | 2.09* | – | – | – | – | $2.19\pm 0.02$ | 38.8* | $0.87\pm 0.04$ | $0.0130\pm 0.005$ | $561.0/426$
1200310122 | 58376.04 | 2.09* | $0.871*$ | $2.4\pm 0.3$ | – | – | $1.81^{+0.08}_{-0.07}$ | 38.8* | $0.60\pm 0.04$ | $0.025\pm 0.002$ | $601.8/548$
1200310124 | 58380.27 | 2.09* | $0.871*$ | $2.2\pm 0.3$ | – | – | $1.90^{+0.09}_{-0.07}$ | $58^{+289}_{-23}$ | $0.64\pm 0.04$ | $0.028\pm 0.002$ | $511.7/554$
1200310125 | 58381.23 | 2.09* | $0.871*$ | $3.1\pm 0.5$ | – | – | $1.77^{+0.09}_{-0.07}$ | $33^{+20}_{-8}$ | $0.53\pm 0.6$ | $0.033^{+0.005}_{-0.004}$ | $401.3/420$
1200310126 | 58382.01 | 2.09* | $0.871*$ | $2.2\pm 0.2$ | – | – | $1.89^{+0.07}_{-0.06}$ | $40^{+54}_{-13}$ | $0.62\pm 0.03$ | $0.029\pm 0.002$ | $605.5/621$
1200310127 | 58384.32 | 2.09* | $0.871*$ | $2.1\pm 0.3$ | – | – | $1.91^{+0.09}_{-0.08}$ | $33^{+33}_{-10}$ | $0.63\pm 0.04$ | $0.025\pm 0.002$ | $555.6/561$
1200310128 | 58385.1 | 2.09* | $0.871*$ | $2.0\pm 0.3$ | – | – | $1.90^{+0.08}_{-0.07}$ | 38.8* | $0.61\pm 0.04$ | $0.023\pm 0.002$ | $525.1/523$
1200310129 | 58386.06 | 2.09* | $0.871*$ | $2.1\pm 0.3$ | – | – | $1.90\pm 0.06$ | 38.8* | $0.61^{+0.03}_{-0.04}$ | $0.0215^{+0.002}_{-0.001}$ | $688.0/630$
1200310130 | 58387.16 | 2.09* | $0.871*$ | $2.0\pm 0.4$ | $0.68^{+0.07}_{-0.08}$ | $49^{+14}_{-10}$ | $2.0\pm 0.1$ | 38.8* | $1.4\pm 0.4$ | $0.004^{+0.003}_{-0.001}$ | $468.7/465$
1200310131 | 58389.02 | 2.09* | $0.871*$ | $2.0\pm 0.3$ | – | – | $1.91\pm 0.06$ | 38.8* | $0.57\pm 0.04$ | $0.019^{+0.002}_{-0.001}$ | $619.8/613$
1200310132 | 58395.08 | 2.09* | $0.871*$ | $2.1\pm 0.4$ | – | – | $1.87^{+0.21}_{-0.07}$ | $\geq 45$ | $0.45\pm 0.05$ | $0.0168^{+0.003}_{-0.002}$ | $453.4/453$
1200310133 | 58396.76 | 2.09* | – | – | – | – | $2.0^{+0.3}_{-0.2}$ | 38.8* | $0.62\pm 0.07$ | $0.0075^{+0.0007}_{-0.0006}$ | $134.4/120$
Table 3: Best fit parameters for the emission lines detected in the IGR J17591–2342 spectrum for each observation. The line energy for the Ne ix ion was fixed to the rest-frame energy, while the sigma was fixed to 0.085 keV for each ion transition (i.e. the spectral resolution of NICER/XTI at 1 keV). All the associated errors are reported at 68% confidence level. |  | Ne ix | Fe xxv
---|---|---|---
ObsID | MJD | Norm | LineE | Sigma | Norm
|  |  | (keV) | (keV) |
1200310101 | 58345.0 | $0.015\pm 0.002$ | – | – | –
1200310102 | 58348.09 | $0.008^{+0.003}_{-0.002}$ | – | – | –
1200310103 | 58353.04 | $0.021\pm 0.002$ | $6.6^{+0.3}_{-0.4}$ | $0.5\pm 0.3$ | $0.00015^{0.00006}_{-0.00005}$
1200310104 | 58354.08 | $0.019^{+0.003}_{-0.002}$ | – | – | –
1200310105 | 58357.03 | $0.0108^{+0.0016}_{-0.0014}$ | $\leq 6.3$ | $1.5\pm 0.2$ | $0.0026^{+0.0007}_{-0.0008}$
1200310106 | 58358.0 | $0.024^{+0.007}_{-0.009}$ | $6.59\pm 0.08$ | $0.1*$ | $0.00018^{+0.00002}_{-0.00005}$
1200310107 | 58360.64 | $0.034^{+0.007}_{-0.006}$ | – | – | –
1200310108 | 58361.93 | $0.028^{+0.007}_{-0.006}$ | – | – | –
1200310109 | 58361.99 | $0.016\pm 0.002$ | – | – | –
1200310110 | 58363.28 | $0.016\pm 0.002$ | – | – | –
1200310111 | 58364.05 | $0.016^{+0.003}_{-0.002}$ | $6.7\pm 0.2$ | $0.8^{+0.3}_{-0.2}$ | $0.00019^{+0.00005}_{-0.00004}$
1200310112 | 58366.75 | $0.0068^{+0.0019}_{-0.0015}$ | – | – | –
1200310113 | 58367.02 | $0.017^{+0.005}_{-0.004}$ | – | – | –
1200310114 | 58368.36 | $0.016^{+0.004}_{-0.003}$ | $6.53^{+0.07}_{-0.06}$ | $0.02^{+0.10}_{-0.02}$ | $0.00009\pm 0.00004$
1200310115 | 58369.07 | $0.017\pm 0.003$ | – | – | –
1200310116 | 58370.17 | $0.019\pm 0.003$ | – | – | –
1200310118 | 58372.03 | $0.016^{+0.005}_{-0.004}$ | – | – | –
1200310119 | 58373.0 | $0.019^{+0.004}_{-0.003}$ | – | – | –
1200310120 | 58374.22 | $0.030^{+0.006}_{-0.005}$ | – | – | –
1200310121 | 58375.33 | $0.048^{+0.020}_{-0.014}$ | – | – | –
1200310122 | 58376.04 | $0.041^{+0.011}_{-0.009}$ | – | – | –
1200310124 | 58380.27 | $0.031^{+0.009}_{-0.007}$ | – | – | –
1200310125 | 58381.23 | $0.08^{+0.04}_{-0.03}$ | – | – | –
1200310126 | 58382.01 | $0.033^{+0.007}_{-0.006}$ | – | – | –
1200310127 | 58384.32 | $0.019^{+0.005}_{-0.004}$ | – | – | –
1200310128 | 58385.1 | $0.024^{+0.007}_{-0.005}$ | – | – | –
1200310129 | 58386.06 | $0.019^{+0.004}_{-0.003}$ | – | – | –
1200310130 | 58387.16 | $0.017^{+0.006}_{-0.005}$ | – | – | –
1200310131 | 58389.02 | $0.018^{+0.004}_{-0.003}$ | – | – | –
1200310132 | 58395.08 | $0.013^{+0.004}_{-0.003}$ | $6.63^{+0.06}_{-0.05}$ | $0.04331^{+0.09}_{-0.04}$ | $0.00013^{+0.00005}_{-0.00004}$
1200310133 | 58396.76 | – | – | – | –