11institutetext: 1INAF – Osservatorio Astronomico di Brera, Via E. Bianchi 46,
23807 Merate (LC), Italy 11email<EMAIL_ADDRESS>
2Max-Planck-Institut für extraterrestrische Physik, Gießenbachstraße 1, 85748,
Garching, Germany
3School of Astronomy and Space Science, Nanjing University, Nanjing 210046,
China
4Columbia Astrophysics Laboratory, Columbia University, Columbia, NY, 10027,
USA
5INAF – Istituto di Astrofisica Spaziale e Fisica Cosmica, via A. Corti 12,
20133 Milano, Italy
6Department of Physics and Astronomy, University of California, Los Angeles,
CA, 90095-1547, USA
7Institute of Space Sciences (ICE, CSIC), Campus UAB, Carrer de Can Magrans
s/n, E-08193 Barcelona, Spain
8Institut d’Estudis Espacials de Catalunya (IEEC), Carrer Gran Capità 2–4,
08034 Barcelona, Spain

# Periodicity from X-ray sources within the inner Galactic disk

Samaresh Mondal 11 Gabriele Ponti 1122 Tong Bao 33 Frank Haberl 22 Sergio
Campana 11 Charles J. Hailey 44 Shifra Mandel 44 Sandro Mereghetti Kaya Mori
5544 Mark R. Morris 66 Nanda Rea 7788 and Lara Sidoli 55

(Received XXX; accepted YYY)

###### Abstract

Aims. For many years it had been claimed that the Galactic ridge X-ray
emission at the Galactic Center (GC) is truly diffuse in nature. However, with
the advancement of modern X-ray satellites, it has been found that most of the
diffuse emission actually comprises thousands of previously unresolved X-ray
point sources. Furthermore, many studies suggest that a vast majority of these
X-ray point sources are magnetic cataclysmic variables (CVs) and active
binaries. One unambiguous way to identify these magnetic CVs and other sources
is by detecting their X-ray periodicity. Therefore, we systematically searched
for periodic X-ray sources in the inner Galactic disk, including the GC
region.

Methods. We used data from our ongoing XMM-Newton Heritage Survey of the inner
Galactic disk ($350\degr\lesssim l\lesssim+7\degr$ and $-1\degr\lesssim
b\lesssim+1\degr$) plus archival XMM-Newton observations of the GC. We
computed the Lomb-Scargle periodogram for the soft (0.2–2 keV), hard (2–10
keV), and total (0.2–10 keV) band light curves to search for periodicities.
Furthermore, we modeled the power spectrum using a power-law model to simulate
1000 artificial light curves and estimate the detection significance of the
periodicity. We fitted the energy spectra of the sources using a simple power-
law model plus three Gaussians, at 6.4, 6.7, and 6.9 keV, for the iron $K$
emission complex.

Results. We detected periodicity in 26 sources. For 14 of them, this is the
first discovery of periodicity. For the other 12 sources, we found periods
similar to those already known, indicating no significant period evolution.
The intermediate polar (IP) type sources display relatively hard spectra
compared to polars. We also searched for the Gaia counterparts of the periodic
sources to estimate their distances using the Gaia parallax. We found a likely
Gaia counterpart for seven sources.

Conclusions. Based on the periodicity, hardness ratio, and the equivalent
width of Fe $K$ line emission, we have classified the sources into four
categories: IPs, polars, neutron star X-ray binaries, and unknown. Of the 14
sources for which we detect the periodicity for the first time, four are
likely IPs, five are likely polars, two are neutron star X-ray binaries, and
three are of an unknown nature.

###### Key Words.:

X-rays:binaries – Galaxy:center – Galaxy:disk – white dwarfs – pulsars –
novae, cataclysmic variables

##  1 Introduction

In order to understand the star formation history of our Galaxy, it is
important to know the number of stars that ended their main-sequence life long
ago. Compact remnants of dead stars, such as black holes, neutron stars (NSs),
and white dwarfs (WDs), are commonly found in binary systems and are visible
in X-rays. Accreting WD binaries are the most common type of remnant in our
Galaxy as WDs are the end product of intermediate- and low-mass stars. Many of
these low-mass stars are born in binary systems with small separations that go
through one or more mass transfer phase, leading to the formation of
cataclysmic variables (CVs). More than a thousand CVs have been found in the
solar neighborhood (Downes et al., 2001; Ritter & Kolb, 2003). CVs are
categorized as magnetic or nonmagnetic (see Cropper, 1990; Patterson, 1994;
Mukai, 2017, for a review). Magnetic CVs are primarily categorized into two
subtypes: polar and intermediate polar (IP). Polars have a very strong
magnetic field ($>10$ MG), which synchronizes the spin and orbital motion
(i.e., $P_{\rm spin}=P_{\rm orb}$). The high magnetic field in polars is
confirmed by the observation of strong optical polarization and the
measurement of cyclotron humps (Warner, 2003). In polars, the accretion
directly follows the magnetic field lines from the L1 point, and no accretion
disk is formed. IPs have a relatively weak surface magnetic field of $1-10$
MG; therefore, they have less synchronization, and an accretion disk is
created. In these systems, the material leaving the L1 point forms an
accretion disk until the magnetic pressure becomes equal to the ram pressure
of the accreting gas. The X-ray emission from CVs originates from close to the
magnetic pole. The accreting material follows the magnetic field lines, and as
it approaches the WD surface, the radial in-fall velocity reaches supersonic
speeds of 3000-10000 km s-1. A shock front appears above the WD surface, and
the in-falling gas releases its energy in the shock, resulting in hard X-ray
photons (Aizu, 1973; Saxton et al., 2005).

Early observations of the Galactic Center (GC) revealed a diffuse X-ray
emission (Worrall et al., 1982; Warwick et al., 1985; Yamauchi et al., 1996)
called Galactic ridge emission. For many years a central point of debate has
been whether the Galactic ridge emission is truly diffuse or composed of
emission from many unresolved X-ray point sources. The advent of modern X-ray
satellites such as XMM-Newton and Chandra opened up the possibility of
detecting very faint X-ray sources in crowded regions such as the inner GC.
This is not possible in the optical waveband due to the high extinction toward
the GC. A deep Chandra observation of the inner Galactic bulge has
demonstrated that more than 80% of the Galactic ridge emission is produced by
CVs and coronally active stars (Wang et al., 2002; Revnivtsev et al., 2009;
Muno et al., 2003a, 2009; Zhu et al., 2018). Although this strongly indicates
that a large fraction of the X-ray sources observed toward the GC are magnetic
CVs, the physical nature of CVs in the GC remains unclear. Moreover, it was
suggested that, based on the hard power-law-type spectral shape and the
emission of Fe $K$ complex lines, the majority are IPs (Muno et al., 2004)

The Galactic ridge X-ray emission displays a copious amount of lines from
ionized iron at 6.7 and 6.9 keV. Some studies that compared the stellar mass
distribution with the Fe XXV (6.7 keV) line intensity map suggest the presence
of truly diffuse hard X-ray emission (Uchiyama et al., 2011b; Nishiyama et
al., 2013; Yasui et al., 2015). However, a recent study by our group found
that this diffuse hard emission in the GC can be explained if one assumes that
the GC stellar population has iron abundances $\sim 1.9$ times higher than
those in the Galactic bar/bulge (Anastasopoulou et al., 2023). Furthermore,
the 20–40 keV emission from the GC observed by NuSTAR is best explained by
two-temperature plasma models with $kT_{1}\sim 1$ keV and $kT_{2}\sim 7.5$
keV. The $\sim 1$ keV temperature component is attributed to emission from
supernovae heating the interstellar medium, coronally active stars, and
nonmagnetic WDs (Revnivtsev et al., 2009). The $\sim 7.5$ keV temperature
component is thought to be produced by emission from resolved and unresolved
accreting IPs (Perez et al., 2015). An additional component with a higher
plasma temperature, $kT\sim 35$ keV (Hailey et al., 2016), was recently
measured. In addition, Muno et al. (2003b) reported the discovery of eight
periodic sources in a $17^{\prime}\times 17^{\prime}$ field of the GC. Their
periods range from 300 s to 4.5 hr. All these sources exhibit hard power-law-
type spectral shapes (with photon index $\Gamma\sim 0$) and 6.7 keV iron-line
emission. These properties are consistent with magnetically accreting magnetic
CVs.

We are in the process of performing an X-ray scan of the inner Galactic disk
using XMM-Newton (Jansen et al. 2001; PI: G. Ponti). The main aim of this
survey is to constrain the flow of hot baryons that feed large-scale energetic
events such as the Galactic chimneys (Ponti et al., 2019, 2021), the _Fermi_
bubbles (Su et al., 2010), and the eROSITA bubbles (Predehl et al., 2020). In
early 2021, while performing this survey, we detected an X-ray source with
periodic modulation at 432 s (Mondal et al., 2022). The source was previously
observed by Suzaku in 2008 (Uchiyama et al., 2011a) and classified as an IP.
Furthermore, while examining XMM-Newton archival observations, we discovered
periodicity in two other sources within 1$\aas@@fstack{\circ}$5 of the GC
(Mondal et al., 2023). These two sources are also classified as IPs based on
the detected spin period and detection of an iron emission complex in the
spectra. Therefore, we took a systematic approach to hunt for such periodic
X-ray sources that might help us classify them. In this paper we report the
discoveries obtained from a periodicity search using XMM-Newton observations
of the Galactic disk and the GC.

##  2 Observations and data reduction

We have almost completed one side of the Galactic disk extending from $l\geq
350\degr$ to $l\leq+1.5\degr$ (see Fig. 1). The survey has an exposure of 20
ks per tile and is expected to cover the Galactic disk region in the range
$350\degr<l<+7\degr$ and $-1\degr<b<+1\degr$. During this campaign, we
detected thousands of X-ray point sources of various types. A forthcoming
paper will present a sample study of the X-ray point sources. Here we are
focusing on X-ray sources that show periodic modulations. While doing this
analysis, we considered including the GC region for positional comparison of
the sources located in the disk and GC (Ponti et al., 2013, 2015). For the GC,
we used the XMM-Newton archival observations. In total, we analyzed 444 XMM-
Newton observations, including our Galactic disk scanning observations plus
the archival observations of the GC. The observation data files were processed
using the XMM-Newton Science Analysis System (SAS,
v19.0.0)111https://www.cosmos.esa.int/web/xmm-newton/sas. We used the task
evselect to construct a high energy background light curve (energy between 10
and 12 keV for EPIC-pn and above 10 keV for EPIC-MOS1 and MOS2) by selecting
only PATTERN==0. The background light curve was used to filter high background
flaring activity and create good time intervals. We used the SAS task
emldetect for point source detection and source list creation. For each
individual detector, EPIC-pn, MOS1, or MOS2 (Strüder et al., 2001; Turner et
al., 2001), the source detection was performed in five energy bands: 0.2–0.5
keV, 0.5–1 keV, 1–2 keV, 2–4 keV, and 4–12 keV for a given single observation.
The source detection algorithm separately provides net counts and maximum
likelihood values for the five energy bands and three detectors: EPIC-pn,
MOS1, and MOS2. The source detection tool also provides the keyword EXT value
that indicates whether the emission is from a point-like or extended source.
We chose EXT $=0$ to select the point sources only. The total number of point
sources detected in our survey is $\sim 50000$. Then, we applied a filter in
which only sources with a total number of counts (EPIC-pn+MOS1+MOS2) higher
than 200 were chosen. This resulted in 2500 point sources for which we
extracted the light curves using the SAS task evselect after applying the
Solar System barycenter correction to the event files using the SAS task
barycen. We only selected events with PATTERN$\leq$4 and PATTERN$\leq$12 for
EPIC-pn and the MOS1 and MOS2 detectors, respectively. We chose a circular
region of 20″ radius for the source products extraction. The background
products were extracted from an annular region centered on the source position
with inner and outer radii of 25″ and 30″, respectively. The spectra were
binned to have a minimum of 20 counts in each energy bin. Many fields of the
GC have been observed more than once. If a source has been observed multiple
times, we searched for pulsations in each observation individually.

##  3 Results

Figure 1: Mosaic of the exposure maps created using the ongoing XMM-Newton
observations of the Galactic disk plus archival observations of the GC. The
small red, blue, and green circles show the positions of confirmed or likely
NSs, IPs, and polars, respectively. The black circles indicate the
unclassified sources.

###  3.1 Period search

The XMM-Newton observations suffer from gaps due to the filtering of high
background flaring activities. As the XMM-Newton observations suffer from
gaps, we used the Lomb-Scargle periodogram (Lomb, 1976; Scargle, 1982), which
is well known for detecting periodic signals in unevenly sampled time series
data. We computed the false alarm probability to estimate the statistical
significance of the periodogram peaks. The false alarm probability obtained is
based on the analytical approximation proposed by Baluev (2008), which employs
extreme value statistics to compute an upper bound of the false alarm
probability (or a lower limit of the significance of the detected
periodicity). For our timing analysis, we used the PYTHON-based astropy
(Astropy Collaboration et al., 2013, 2018, 2022) package’s time-series
module222https://docs.astropy.org/en/stable/timeseries/index.html.

We extracted the light curves in three different bands (0.2–2, 2–10, and
0.2–10 keV) for the three EPIC detectors (pn, MOS1, and MOS2). The light
curves were extracted with a time bin of 74 ms for the EPIC-pn and 3 s for the
MOS1 and MOS2 detectors in full frame mode of observation. The Lomb-Scargle
periodogram was computed for all nine light curves of each source, and a
periodicity search was conducted. The EPIC-pn detector has a frame time of 74
ms, which allowed us to probe a maximum frequency of $\sim 6$ Hz, whereas in
the case of the MOS1 and MOS2 detectors, we were able to probe a maximum
frequency of $\sim 0.16$ Hz. We imposed the criterion that the periodicity
detected at a frequency below 0.16 Hz should be present in all three
detectors. To search for periodicity at a higher frequency within the range
0.16–6 Hz, we used only the data from EPIC-pn. We have detected periodicity at
a significance above $3\sigma$ in 23 sources. Possible periodicities with
significance between 2 and 3$\sigma$ were found in another three sources333We
also detected a few sources with detection significance of the pulsation just
above the 1$\sigma$ confidence level. We did not list these sources in Table
2; one such example is the transient Galactic bulge IP XMMU J175035.2–293557
(Hofmann et al., 2018).. Figures 8 and 9 show the Lomb-Scargle periodograms of
the 26 sources, with the horizontal lines indicating the detection
significance levels.

Figure 1 shows the mosaic of the exposure maps created from our ongoing XMM-
Newton observations of the Galactic disk plus the XMM-Newton archival
observations of the GC. The small circles indicate the positions of the
periodic sources. We have completed the survey of one side of the Galactic
disk extending to $\sim 350^{\circ}$; however, most pulsators are concentrated
near the GC.

Table 2 shows the details of the X-ray properties of the periodic sources. The
period column shows the pulsation period obtained in our analysis and compares
it with the previously reported period. The pulse fraction is computed in 2–10
keV bands. The name of the sources is taken from the 4XMM catalog except for
the sources XMMU J173029.8–330920, XMMU J175441.9–265919, and XMMU
J180140.3–23422, which were not listed in the 4XMM (Webb et al., 2020) archive
as these sources were first detected in our campaign. The X-ray position and
$1\sigma$ positional error of the sources are taken from the 4XMM catalog. We
also list the source type based on previous studies, and for sources that were
not classified before, we give a tentative classification based on the X-ray
period, hardness ratio (HR) values, and spectral properties. Figure 2 shows
the distribution of the log of pulse period for various source types. Figure 3
shows the distribution of pulse fraction for the different categories. The
pulse fraction was computed as $\rm
PF=\frac{F_{max}-F_{min}}{F_{max}+F_{min}}$, where $\rm F_{max}$ and $\rm
F_{min}$ are the maximum and minimum counts in the folded light curves,
respectively. For sources with more than one XMM-Newton observation, we report
the periodicity from the multiple observations in Tables 3 and 4.

Figure 2: Distribution of log(period) for various source types. Figure 3:
Distribution of the 2-10 keV pulse fraction in percent for different source
types.

###  3.2 Caveats for the false alarm probability

Accretion-powered systems such as CVs and X-ray binaries are known to exhibit
aperiodic variability over a wide range of timescales. This irregularity,
often referred to as ”red noise,” constitutes a significant aspect of
aperiodic variability and has the potential to introduce spurious periodic
signals, especially at lower frequencies (Warner, 1989). Consequently, it is
essential to assess the likelihood of false detections among these periodic
signals found with the Lomb-Scargle periodogram method by using a large
simulated dataset (Bao & Li, 2022).

Specifically, we employed a power-law model to characterize the source power
spectrum, which has the form of

$P(\nu)=N\nu^{-1}+C_{\rm p}.$ (1)

In this equation, $N$ represents the normalization factor, and $C_{\rm p}$
accounts for the Poisson noise, which is influenced by the mean photon flux of
the source.

To begin, we estimated the power spectral density (PSD) using the standard
periodogram approach with an $\rm[rms/mean]^{2}$ PSD normalization (Vaughan et
al., 2003). However, as mentioned in Section 3.1, some of the light curves
suffer from gaps due to background flares. For these cases, we filled the gap
with artificial light curves of Poisson noise, assuming the mean flux is
consistent with that of the source. Although such processing results in little
differences in the described PSDs, for most of the periodic sources here these
gaps are fortunately negligible in terms of time (i.e., they take less than
0.5% of the total exposure time). Only one case exhibits a significant data
gap, which takes $\sim 1.4\%$ of the single observation, with
ObsID=0783160101. This source (4XMM J174816.9-280750) consistently exhibits
the same periodic signal across multiple observations (see Table 4). Thus, the
possible uncertainty of its confidence estimation by the process of filling
gaps will not impact the verification of its periodicity.

We fitted the power spectrum of each source with Eq. 1, using the maximum
likelihood function discussed in Vaughan (2010) and the Markov chain Monte
Carlo approach, employing the Python package _emcee_
444https://emcee.readthedocs.io/en/stable/ (Foreman-Mackey et al., 2013) to
derive the best-fit parameters and their associated uncertainties.

It turns out that only three of the periodic sources could be adequately
described by the power-law model with constrained normalization, implying a
potential influence of red noise. For the remaining sources, Poisson noise
actually dominates the source power spectrum. Thus, for the source with
potential red noise, simulated light curves for this best-ﬁfit power-law model
were constructed using the method of Timmer & Koenig (1995), which were
resampled and binned to have the same duration, mean count rate, and variance
as the observed light curve. As for sources where Poisson noise prevailed, we
followed a similar procedure to simulate their light curves, assuming pure
Poisson noise. A group of 1000 simulated time series was produced for each
source. To evaluate the false alarm level, we computed the maximum Lomb-
Scargle power for each simulated light curve. Specifically, we considered the
top 0.3% of the maximum Lomb-Scargle power from the 1000 Lomb-Scargle
periodograms, corresponding to the 3$\sigma$ confidence level estimation, and
the top 5% as the threshold corresponding to 2$\sigma$ (approximately 95%).
These simulated thresholds were then overlaid on the Lomb-Scargle periodogram
(Figs. 8 and 9), and the confidence levels, calculated using Baluev’s analysis
method (Baluev, 2008), were compared. It turns out that 23 sources exceed the
simulated-based threshold of 3$\sigma$, and 17 of them exceed the 3$\sigma$
threshold of Baluev’s method. The deviation between these two is mainly due to
that the Baluev method, by design, provides an upper limit to the false alarm
probability with little overestimation (Baluev, 2008).

###  3.3 Period and pulse fraction distribution

The top panel of Fig. 2 shows the period distribution of sources in our
sample. The distribution has two peaks at around $\sim 800$ s and $\sim 4800$
s. The first peak is associated with the population of IPs, and the second
peak corresponds to the population of polars.

The spin period of NS and likely NS systems in our sample ranges from 1.36 s
to 411.3 s. The spin period of Galactic NS high-mass X-ray binaries (HMXRBs)
ranges from a few to thousands of seconds, and the distribution has peaks
around $\sim 250$ s (Neumann et al., 2023). The red histogram in the top panel
of Fig. 2 shows the distribution of the period for NS X-ray binaries in our
sample.

The blue and cyan histogram in the middle panel of Fig. 2 shows the period
distribution for the known IPs plus the tentative identification of IPs in our
sample. The distribution has a peak of around 607 s. Typically, the spin
period of IPs ranges between 30 s and 3000 s, with a peak near 1000 s
(Scaringi et al., 2010). The middle panel of Fig. 3 (blue and cyan histogram)
shows the distribution of pulse fraction for IPs. The pulse fraction ranges
from 10% to 80% with a peak near 45%. One prominent feature in IPs is that the
pulse fraction or the modulation depth typically increases with decreasing
X-ray energy. This has been thought to be the effect of photoelectric
absorption (Norton & Watson, 1989). The distribution of pulse fraction of IPs
covers a wide range of scales and can vary from a few percent up to $\sim$100%
with an average around 24% (Norton & Watson, 1989; Haberl & Motch, 1995).

In polars, the spin and orbital periods are synchronized, ranging from 3000 s
to 30000 s (Scaringi et al., 2010). In our sample, the periods of polars vary
from 4093 s to 6784 s, with the peak at $\sim 4800$ s. The light curves of
polars show constant modulation of depth with X-ray energies. The depths are
generally higher compared to IPs (Norton & Watson, 1989). In the middle panel
of Fig. 3, it is evident that the pulse fraction for polars starts at higher
values, around 30% compared to IPs, and that more polar type sources are found
between 50% and 60%.

###  3.4 Spectral modeling

We performed time-averaged spectral modeling using the X-ray spectral fitting
software xspec555https://heasarc.gsfc.nasa.gov/xanadu/xspec/ (Arnaud, 1996).
We employed $\chi^{2}$ statics in our model fitting. The spectra were fitted
using a simple model composed of a power law and three Gaussian lines
(tbabs(power-law+g1+g2+g3)). The Galactic absorption component is represented
by tbabs (Wilms et al., 2000). For the continuum, we used a simple power-law
model, and g1, g2, and g3 represent the three Gaussian lines at 6.4, 6.7, and
6.9 keV, respectively, for iron emission complex. While doing the fit, we
freeze the line energies at the expected values, and the width of the lines is
fixed at zero eV. We jointly fit the spectra of EPIC-pn, MOS1, and MOS2
detectors. While fitting the spectra, we included a constant factor for cross-
calibration uncertainty, which is fixed to unity for EPIC-pn and allowed for
variation for MOS1 and MOS2. The spectral fitting results are summarized in
Table 5, and Figs. 10 and 11 show the fitted spectra of the sources.

Figure 4 shows the distribution of absorption column density $N_{\rm H}$
obtained from the X-ray spectral fitting. Overall, the $N_{\rm H}$
distribution has a peak near $10^{22}$ cm-2, and more than 50% of the sources
have $N_{\rm H}$ between $10^{21}-3.16\times 10^{22}$ cm-2. There are three
sources with high $N_{\rm H}$¿$10^{23}$ cm-2. The source 4XMM J175327.8–295716
has $N_{\rm H}=7\times 10^{20}$ cm-2, the lowest in our sample, which might
indicate that this source is the closest to us among our sample or has a soft
component that mimics the low $N_{\rm H}$.

Figure 5 shows the distribution of photon index $\Gamma$. The distribution has
a peak at $\Gamma\sim 0.6$. More than 50% of the sources have a flat spectral
shape with $\Gamma$¡1. A significant number of sources in our sample have a
softer spectrum with $\Gamma$¿1.2. We noticed that the majority of sources
with high $\Gamma$ values do not show any iron emission complex lines; only
two of the seven sources with $\Gamma\geq 1.3$, show strong emission lines in
the 6–7 keV band.

Figure 4: Distribution of the absorption column density, $N_{\rm H}$, for
different source types. Figure 5: Distribution of the photon index, $\Gamma$,
for different source types.

###  3.5 Gaia counterparts

A correct estimate of the distance to the source is required to derive their
luminosity. For this, we searched for counterparts in the Gaia DR3 catalog.
For each X-ray source, we computed the Gaia source density by counting the
number of Gaia sources within a circle of 1′ radius at the source position.
The Gaia density of sources is low and varies from 0.01–0.1 arcsec-2. We
compute the probability of the sources having a spurious association by
multiplying the Gaia source density with the area associated with the XMM-
Newton positional error. Table 1 lists the sources for which we found a Gaia
counterpart within the 3$\sigma$ positional uncertainty of XMM-Newton. We
found a likely Gaia counterpart for seven XMM-Newton sources. If a counterpart
is found, then we use the Gaia source ID to find the distance to the source
from Bailer-Jones et al. (2021). The distance to the sources for which a Gaia
counterpart was found varies from $\sim 1.5$ to $\sim 5$ kpc and the X-ray
luminosity is in the range $5\times 10^{32}-6\times 10^{33}$ erg s-1.

##  4 Discussion

###  4.1 Typical properties of different classes of sources

We analyzed 444 XMM-Newton observations of the GC and the Galactic disk. We
extracted X-ray light curves from nearly 2500 sources and systematically
searched for X-ray pulsation. We detected periodicity in 26 sources, 14 of
which are reported here for the first time. Many of the GC sources have a
luminosity of a few times $10^{32}$ erg s-1 (Muno et al., 2003a), which is
comparable to the luminosity typically observed in bright magnetic CVs
(Verbunt et al., 1997; Ezuka & Ishida, 1999). NS HMXRBs have much higher
luminosity and are detected during their outburst period, reaching
luminosities up to $10^{38}$ erg s-1. For the majority of the sources we did
not find a Gaia counterpart due to the high absorption column density toward
the GC and disk. Hence, the X-ray luminosity cannot be derived for a large
number of sources and we cannot use this information to classify them. The NS
HMXRBs are far less common in our Galaxy than the magnetic CVs. The nature of
the Galactic sources is a long-standing question. Identifying the magnetic CVs
or NS HMXRBs from X-ray periodicity alone can be difficult, as both types of
sources usually display periods in a similar range. The short-period
modulation in the X-ray light curve is thought to have originated from the
spin period of the magnetically accreting WD or NS. In our sample, the
smallest detected period is 1.36 s, and the maximum period detected is around
6784 s. The sample has a median period of 672 s. The pulse fraction of the
modulation ranges from 10% to 80%. The detected periods are consistent with
those of magnetic CVs and NSs in HMXRBs. A sample study of magnetic CVs
indicates the median spin period is 6000 s (Ritter & Kolb, 2003). There are
also a few magnetic CVs with very short spin periods; for example, CTCV
J2056–3014 has a spin period of 29.6 s (Lopes de Oliveira et al., 2020), and
V1460 Her has a spin period of 38.9 s (Ashley et al., 2020). The spin periods
of polars are mostly beyond 1 hr, while almost all IPs have WD spin periods
lower than 1 hr. In contrast, 85 Galactic HMXRB pulsars (both with Be and OB
supergiant companions) have a median (mean) spin period of 187 s ($\sim$970
s), with only four sources showing a period longer than 1 hr (Neumann et al.,
2023).

It is evident that the different classes of sources (NS HMXRBs and magnetic
CVs) exhibit a wide range of spin periods. Therefore, from the periodicity
alone, it is difficult to understand the nature of the unclassified sources.
Below we summarize a scheme to characterize the different classes of periodic
X-ray sources utilizing their X-ray spectral, timing properties, and
luminosity.

####  4.1.1 NS HMXRBs

The NS HMXRBs have properties that are very similar to IPs and they typically
have very hard spectra. Figure 6 shows the period versus HR plot for
classified and unclassified sources in our sample. The HR is calculated using
the net counts in the 2–5 keV and 5–10 keV bands. We did not choose an energy
band below 2 keV simply because it would be affected by Galactic absorption.
The known NS HMXRBs appear very hard, similar to IPs; however, it is clear
from Fig. 7 that they emit very little 6.7 keV iron line as compared to IPs.
In almost all NS HMXRBs, the dominant component of the Fe K emission complex
is the neutral 6.4 keV line emission and little to no ionized 6.7 and 6.9 keV
line emission. This is because HMXRBs are mainly wind-fed systems, so the
fluorescent iron line emission from the wind of the companion star is the main
spectral feature in their spectra, while the ionized iron emission lines
usually come from an accretion disk. The known NS HMXRBs in our sample – 4XMM
J172511.3–361657 and 4XMM J174906.8–273233 – show no 6.7 keV emission, with
upper limits on their equivalent widths (EWs) of 8 eV and 15 eV, respectively.
We define the following criteria for the characterization of the NS HMXRB: (i)
$P_{\rm spin}\lesssim$1000 s, (ii) HR¿-0.2, (iii) $\rm EW_{6.7}$¡50 eV, and
(iv) a typical X-ray luminosity of $10^{33}-10^{37}$ erg s-1.

####  4.1.2 IPs

One of the prominent features of IPs is the presence of strong ionized 6.7 keV
line emission. In our sample, all the confirmed IPs have a clear detection of
a 6.7 keV line, with the lowest EW of the sample being $78^{+34}_{-19}$ for
4XMM J174517.0–321358. Xu et al. (2016) studied a sample of bright 17 IPs
using Suzaku data. They found that the minimum and mean EW of the 6.7 keV line
of the sample is $58^{+10}_{-13}$ and $107\pm 17$ eV, respectively. The below
criteria can be used to characterize IPs. They typically have (i) a spin
period $P_{\rm spin}$¡2500 s, (ii) an HR¿-0.2, (iii) a strong 6.7 keV line
emission with $\rm EW_{6.7}$¿50 eV, and (iv) and an X-ray luminosity in the
range $10^{31}-10^{35}$ erg s-1 (Suleimanov et al., 2022).

####  4.1.3 Polars

The X-ray emission from polars is much softer than that of IPs. The spectra of
many polars are dominated by very soft blackbody-like emission from the WD
surface (Osborne et al., 1986; Ramsay et al., 1993; Clayton & Osborne, 1994).
However, toward the GC this component is difficult to observe due to the high
absorption. In general polars also show a strong 6.7 keV line with an EW
anywhere from $50$ eV to $\sim 450$ eV (Ezuka & Ishida, 1999; Xu et al.,
2016). As a whole, the detection of a 6.7 keV line in polar can be difficult
for faint sources as they are much softer than IPs. Polars can be tentatively
classified by having softer spectra and periods above 2500 s; however, a
secure classification would require the detection of the 6.7 keV line with
good quality X-ray spectra and strong circular polarization in the optical
band. The polars can be characterized by the following characteristics: (i)
$P_{\rm spin}=P_{\rm orb}$¿2500 s, (ii) HR¡-0.2, (iii) a strong 6.7 keV line
emission with $\rm EW_{6.7}$¿50 eV, and (iv) an X-ray luminosity below
$10^{33}$ erg s-1 (Suleimanov et al., 2022).

Figure 6: HR vs. period diagram. The HR is calculated using the net counts of
two bands: 2–5 keV and 5–10 keV. Figure 7: EW of the 6.7 keV line vs. period
diagram.

###  4.2 Known NS HMXRBs

The source 4XMM J172511.3–361657 was discovered on 9 February 2004 by INTEGRAL
and named as IGR J17252–3616 (Walter et al., 2004). XMM-Newton observed the
source on 21 March 2004. A period search was performed by Zurita Heras et al.
(2006), and a pulsation of $414.8\pm 0.5$ s was discovered. An orbital period
of $9.737\pm 0.004$ days was also reported by using the Rossi X-ray Timing
Explorer (RXTE) proportional counter array data (Markwardt & Swank, 2003;
Corbet et al., 2005). The source has a flat spectrum with
$\Gamma=0.82^{+0.04}_{-0.04}$, which can also be fitted by a flat power law
with an energy cutoff or a Comptonized model with $kT\sim 5.5$ keV (Zurita
Heras et al., 2006). The spectrum shows a 6.4 keV iron line with an EW of
$70^{+6}_{-7}$ eV. Previous studies indicate the source is a wind-fed
accreting pulsar with a supergiant companion star. The source has been
observed multiple times by XMM-Newton, and we searched for pulsation in all
the observations. The pulsations found in the different observations are
consistent with each other within the 1$\sigma$ error values. The source is
highly variable, and the flux of the source can vary from $2.19\times
10^{-13}$ to $7.42\times 10^{-11}$ erg s-1 cm-2. We noticed that whenever the
source flux drops below $\sim 5\times 10^{-13}$ erg s-1 cm-2 the pulsation was
undetectable.

The source 4XMM J174906.8–273233 was discovered in 1996 by ASCA. The source is
also known as AXJ1749.1–2733 (Sakano et al., 2002). In a 1995 ASCA
observation, the source was not detected, and in 2003 INTEGRAL caught a short
outburst, which indicates its transient nature (Grebenev et al., 2004). XMM-
Newton first observed 4XMM J174906.8–273233 on 31 March 2007, and Karasev et
al. (2008) analyzed EPIC-pn data and detected a spin period of 132 s. The
source was classified as a transient X-ray pulsar in a high-mass binary
system. The source has been observed twice by XMM-Newton in 2007 and 2008;
however, the pulsation was only detected in the 2007 observation (ObsID:
0510010401). The non-detection of pulsations in 2008 could be due to the
combination of two factors: (1) the source flux was almost an order of
magnitude fainter than in the 2007 observation, and (2) the 2008 observation
had a shorter exposure than the 2007 observation, which led to $\sim 22$ times
fewer net counts in the 2008 observation than in the 2007 observation. The
source spectrum is heavily absorbed and can be fitted by a steep power-law
model with $\Gamma=1.3^{+0.1}_{-0.1}$; adding an iron line at 6.4 keV improves
the fit minutely.

###  4.3 Known IPs

The source 4XMM J174517.0–321358 (Gong, 2022; Vermette et al., 2023) was
discovered by Chandra and serendipitously observed by XMM-Newton in 2010. An
iron-line emission complex and a pulsation of 614 s were detected using XMM-
Newton data. The source is classified as an IP with a WD of $0.8M_{\odot}$
(Vermette et al., 2023). The source has been observed twice by XMM-Newton, and
in both observations, we detected a pulsation of 613 s. The X-ray spectrum
looks like that of a typical IP with a flat spectral shape and iron emission
complex.

The source 4XMM J174033.8–301501 was discovered by Suzaku in 2008 (Uchiyama et
al., 2011a). Later, the source was observed by XMM-Newton on 18 March 2021
during a Galactic disk survey (Mondal et al., 2022). The source spectrum is
well described by emission from collisionally ionized diffuse gas with a
plasma temperature of $\sim 15.7$ keV plus an iron line emission complex. A
period of 432.4 s was detected in both Suzaku and XMM-Newton data. The source
has been observed twice by XMM-Newton in 2018 and 2021. In both XMM-Newton
observations, the detected pulsations are consistent. The source has a flat
spectrum with $\Gamma=0.5^{+0.1}_{-0.1}$ and an Fe emission complex in the 6–7
keV band.

4XMM J174954.6–294336 was first discovered by Chandra (Jonker et al., 2014).
The source is classified as an IP based on the spin period of 1002 s and hard
power-law spectral shape with complex iron line emission (Johnson et al.,
2017; Mondal et al., 2023). This is only the second known IP that shows
eclipses in X-rays. The source has been observed twice by XMM-Newton, and the
pulsation is not visible in ObsID 0801681401. Mondal et al. (2023) discuss the
possibility that the pulsation is suppressed due to a complex absorption
behavior and the eclipse seen in the X-ray light curve.

4XMM J174917.7–283329 is classified as IP (Mondal et al., 2023). A period of
1212 s was detected in a 2017 XMM-Newton observation. The continuum is best
fitted by a partially absorbed apec model with a plasma temperature of $13$
keV. The source has been observed three times by XMM-Newton, but the pulsation
was detected only once, when the source flux was one order of magnitude higher
than in the other two observations.

The source 4XMM J174816.9–280750 was observed by BeppoSAX during the GC survey
in 1997–1998 (Sidoli et al., 2006). The source has a spectrum with
$\Gamma=1.3^{+0.6}_{-0.6}$ and strong emission lines at 6–7 keV, plus a
coherent pulsation of period 593 s was found in Suzaku and XMM-Newton data.
These facts favor the source as an IP (Nobukawa et al., 2009). The source has
been observed ten times by XMM-Newton, displaying significant variation in the
pulsation period between different observations. A detailed, in-depth study of
the source is required to determine whether the pulsation period variation is
due to accretion or some other effects, such as the propeller phenomenon.

4XMM J174016.0–290337 was observed by XMM-Newton on 29 September 2005 (Farrell
et al., 2010). The source displays Fe $K_{\alpha}$ emission and a periodic
modulation with a period of 626 s. The source has been observed three times by
XMM-Newton, and in all cases, a pulsation period of 622 s is detected.

The source 4XMM J174009.0–28472 was first discovered by ASCA (Sakano et al.,
2000) and a period of 729 s was found. The source was classified as NS pulsar
based on the flat power-law-type spectrum shape (Sakano et al., 2000).
However, later near-infrared/optical studies suggested it is an IP (Kaur et
al., 2010). The source has been observed four times by XMM-Newton, and we
detected a similar pulsation period value in all observations. The source has
a very flat spectrum $\Gamma=0.1^{+0.1}_{-0.1}$ with strong emission lines.

4XMM J174622.7–285218 is classified as an IP (Nucita et al., 2022). The source
was first observed in a Chandra observation of the GC, and a periodic signal
of 1745 s was found (Muno et al., 2009). The spectrum is characterized by
$\Gamma=0.7^{+0.2}_{-0.2}$, and the 6.9 keV line is the strongest with an EW
of $242^{+81}_{-74}$.

###  4.4 Known polars

The source 4XMM J174728.9–321441 was first observed by Chandra during the
Galactic bulge survey (Jonker et al., 2014). The source is classified as a
polar based on its long period of 4860 s detected in X-rays and in He ii
$\lambda$5412 line emission (Wevers et al., 2017). This source has the
steepest spectrum in our sample with $\Gamma=1.8^{+0.4}_{-0.4}$, and no iron
emission complex was detected in the XMM-Newton spectrum. The non-detection of
iron lines could be due to low signal-to-noise in the data.

###  4.5 Unclassified sources

Below, we try to classify the unknown sources using the scheme defined in
Sect. 4.1. This is a tentative classification; further follow-up of the
individual sources is required to constrain their true nature. For many
sources, we do not have any Gaia counterpart; therefore, the estimation of the
distance to the source using parallax was not possible and hence the
luminosity is not calculated. In such a case, we only used the first three
criteria for classification.

####  4.5.1 Likely NS HMXRBs

The only two sources matching the NS HMXRB criteria are XMMU J175441.9–265919
and 4XMM J175525.0–260402. Both have relatively high HR values: $0.02\pm 0.06$
and $0.05\pm 0.02$, respectively. The upper limits on the EWs of the 6.7 keV
line for J175441.9 and J175525.0 are $\sim 28$ eV and $\sim 34$ eV,
respectively. The spin periods of these two systems are 1.36 s and 392.5 s.
The source J175525.0 was detected three times by XMM-Newton; however, the
pulsation was detected only in the longest observation. A luminosity
estimation was not possible as we did not find any counterparts in Gaia
catalogs.

####  4.5.2 Likely IPs

The sources we categorize as IP are 4XMM J173058.9–350812, 4XMM
J175301.3–291324, 4XMM J175740.5–285105, and 4XMM J175511.6–260315. The
periods found from these systems are below 2500 s and the HRs are above -0.2.
The 6.7 keV line EWs for the sources J173058.9, J175301.3, J175740.5, and
J175511.6 are $236^{+105}_{-83}$, $105^{+189}$, $112^{+61}_{-41}$, and
$194^{+162}_{-188}$ eV, respectively. The source J175301.3 was observed three
times by XMM-Newton, and a $\sim 672$ s period was consistently found in two
of those observations. The source J175511.6 was detected three times by XMM-
Newton; however, the period of $\sim 1135$ s was detected only in the longest
observation. We detected a likely Gaia counterpart for the sources J173058.9
and J175511.6 and the distances estimated from their parallaxes are
$3.2^{+2.2}_{-1.3}$ and $4.9^{+3.2}_{-2.4}$ kpc, respectively. The luminosity
of these two sources is in the range $(0.8-1.9)\times 10^{33}$ erg s-1, which
is typical for accreting magnetic CVs (Suleimanov et al., 2022).

####  4.5.3 Likely polars

The sources that are likely to be polars are 4XMM J173837.0–304818, 4XMM
J175327.8–295716, 4XMM J175244.4–285851, 4XMM J175328.4–244627, and XMMU
J180140.3–234221. These sources have very low HR values and long periods (see
Fig. 6), which suggests that these sources are most likely to be polars. The
long periods are most likely associated with the synchronized spin-orbital
period of the WDs. All these sources have relatively soft spectra of photon
index $\Gamma=0.8-2.1$. For most of the polar-type sources, we have an upper
limit on the EW of the 6.7 keV emission line. This is primarily because these
sources have very low net counts that give $\leq 50$ bins in the 0.2–10 keV
spectrum. The source J175328.4 is the brightest in the polar sample and has
the best signal-to-noise spectrum compared to the other four sources. In this
case, the EW of the 6.7 keV emission line is $28^{+26}_{-28}$ eV, which is
much smaller than the typical values found in IPs. The source J175327.8 was
observed six times by XMM-Newton; however, the periodicity was significantly
detected only in the two observations that have an exposure above 25 ks. We
found a likely Gaia counterpart for the source J175328.4 and its estimated
distance is $1.8^{+0.4}_{-0.2}$ kpc, giving a luminosity of $2\times 10^{33}$
erg s-1.

####  4.5.4 Unknowns

We classify the sources XMMU J173029.8–330920, 4XMM J174809.8–300616, and 4XMM
J175452.0–295758 as unknowns. These sources have high HR and periods similar
to IPs. However, the 6.7 keV line was not detected clearly and we could only
set an upper limit on its EW. The EW of the 6.7 keV line for the sources
J173029.8, J174809.8, and J175452.0 are ¡160, ¡172, and ¡206 eV, respectively.
The source J174809.8 has been observed twice by XMM-Newton. However, it is
relatively faint, with a flux of a few times $10^{-13}$ erg cm-2 s-1 and
therefore the period was detected only in the longer observation. The source
J175452.0 was detected three times by XMM-Newton; however, the pulsation was
detected only in the longest observation.

####  4.5.5 NS or WD?

The compact object in 4XMM J174445.6–271344 is not clearly identified. Also
known as HD 161103, it was observed by XMM-Newton on 26 January 2004. Lopes de
Oliveira et al. (2006) did a detailed multiwavelength spectroscopic study of
this source and suggested that the system hosts an NS; however, a WD scenario
was not excluded. From optical spectroscopy, the companion star of this system
is recognized as a Be star. We detected a periodicity of 3196 s from the X-ray
light curve. The X-ray spectra show strong 6.4, 6.7, and 6.9 keV emission
lines with EWs of $80^{+39}_{-40}$, $371^{+88}_{-76}$, and $109^{+43}_{-47}$
eV, respectively. Such strong 6.7 and 6.9 keV emission lines are not typically
seen in accreting NS HMXRBs. Also, the source has a much softer spectrum ($\rm
HR=-0.42\pm 0.03$ in Fig. 6) than the two confirmed NS HMXRBs (4XMM
J172511.3–361657 and 4XMM J172511.3–361657) in our sample.

##  5 Conclusion

We systematically searched for periodic X-ray sources in the inner Galactic
disk, which extends from $l\sim 350\degr$ to $l\sim+7\degr$ and includes the
GC, using XMM-Newton Heritage observations and archival data. We find 26
sources that show periodicity in their X-ray light curves, of which 12 have
previously reported periods. For these 12 sources, we have obtained periods
consistent with those previously reported. We have detected the periodicity in
the other 14 sources for the first time. We classified the sources based on
the values of the HR, period, and iron emission complex in the 6–7 keV band.
Of these 14 sources, we classify two as NS X-ray binaries, four as likely IPs,
five as likely polars, and three as unknowns. The IP-type sources display a
steep X-ray spectrum with $\Gamma\leq 1.1$ and an iron emission complex in the
6–7 keV band. The spectra of polars are much softer compared to IPs.

Table 1: Sources with possible Gaia optical counterparts.

XMM Name | Density | $G_{\rm mag}$ | Plx | Distance | $L_{\rm x}$
---|---|---|---|---|---
arcsec-2 | mas | kpc | erg s-1
4XMM J173058.9 | 0.009 | 20.05 | $0.59$ | $3.2^{+2.2}_{-1.3}$ | $1.9\times 10^{33}$
4XMM J174033.8 | 0.009 | 19.23 | $0.23$ | $4.3^{+1.5}_{-1.3}$ | $6.4\times 10^{33}$
4XMM J174009.0 | 0.03 | 18.79 | $0.71$ | $1.8^{+1.2}_{-0.5}$ | $2.3\times 10^{33}$
4XMM J174954.6 | 0.1 | 18.97 | $0.61$ | $1.6^{+2.8}_{-2.5}$ | $5.3\times 10^{32}$
4XMM J174816.9 | 0.01 | 21.13 | $0.63$ | $2.9^{+2.3}_{-1.4}$ | $8.6\times 10^{32}$
4XMM J175511.6 | 0.02 | 18.64 | $0.40$ | $4.9^{+3.2}_{-2.4}$ | $8.8\times 10^{32}$
4XMM J175328.4 | 0.02 | 8.40 | $0.55$ | $1.8^{+0.4}_{-0.2}$ | $2.2\times 10^{33}$
666Sources with Gaia counterparts found within the 3$\sigma$ positional error
of XMM-Newton. The density was calculated by drawing a circle with a radius of
1 arcmin at the source position.

###### Acknowledgements.

SM and GP acknowledge financial support from the European Research Council
(ERC) under the European Union’s Horizon 2020 research and innovation program
HotMilk (grant agreement No. 865637). SM and GP also acknowledge support from
Bando per il Finanziamento della Ricerca Fondamentale 2022 dell’Istituto
Nazionale di Astrofisica (INAF): GO Large program and from the Framework per
l’Attrazione e il Rafforzamento delle Eccellenze (FARE) per la ricerca in
Italia (R20L5S39T9). KM is partially supported by the NASA ADAP program
(NNH22ZDA001N-ADAP). We thank the referee for the comments, corrections, and
suggestions that significantly improved the manuscript.

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##  Appendix A Additional figures and tables

Table 2: Various details of the X-ray pulsators in the GC plus Galactic disk.

lat | lon | ObsID | $\rm Pos_{err}$ | Period (s) | PF | XMM Name | PGaia | Sig | Type | References
---|---|---|---|---|---|---|---|---|---|---
arcsec | Our | Previous | 2–10 keV | ¿$\sigma$
351.4972 | -0.3537 | 0886070601 | 0.14 | $414.5\pm 1.4$ | 414.8 | $68.4\pm 1.6$ | 4XMM J172511.3–361657 | $4\times 10^{-4}$ | 3 | NS HMXRB | Zurita Heras et al. (2006)
353.1032 | -0.6956 | 0861171201 | 0.73 | $607.5\pm 3.7$ |  | $45.6\pm 10.8$ | 4XMM J173058.9–350812 | $1.5\times 10^{-2}$ | 3 | Likely IP |
354.7018 | 0.4782 | 0916800201 | 0.35 | $517.6\pm 3.6$ |  | $41.3\pm 12.5$ | XMMU J173029.8–330920 | $1\times 10^{-3}$ | 3 | Unknown | Mondal et al. in prep
357.1486 | -1.6563 | 0865510101 | 0.48 | $614.2\pm 1.4$ | 614 | $28.2\pm 4.5$ | 4XMM J174517.0–321358 | $4\times 10^{-3}$ | 3 | IP | Vermette et al. (2023)
357.3792 | -2.0600 | 0743980401 | 0.57 | $4768.3\pm 151.5$ | 4860 | $74.2\pm 22.4$ | 4XMM J174728.9–321441 | $2.1\times 10^{-2}$ | 3 | Polar | Wevers et al. (2017)
357.6116 | 0.3037 | 0886020101 | 0.79 | $5067.3\pm 260.5$ |  | $57.9\pm 16.2$ | 4XMM J173837.0–304818 | $1.2\times 10^{-2}$ | 3 | Likely Polar |
358.3043 | 0.2442 | 0886010601 | 0.48 | $432.4\pm 1.9$ | 432.1 | $54.8\pm 7.7$ | 4XMM J174033.8–301501 | $6.8\times 10^{-3}$ | 3 | IP | Mondal et al. (2022)
359.2786 | 0.9300 | 0764191201 | 0.17 | $623.2\pm 2.6$ | 626 | $37.2\pm 3.8$ | 4XMM J174016.0–290337 | $1.4\times 10^{-3}$ | 3 | IP | Farrell et al. (2010)
359.2882 | -1.0793 | 0152920101 | 0.65 | $2179.6\pm 18.2$ |  | $64.1\pm 19$ | 4XMM J174809.8–300616 | $2.5\times 10^{-2}$ | 3 | Unknown |
359.4941 | 1.0946 | 0764191101 | 0.6 | $725.9\pm 3.5$ | 729 | $45.1\pm 6.8$ | 4XMM J174009.0–284725 | $3.2\times 10^{-2}$ | 3 | IP | Kaur et al. (2010)
359.8061 | -1.2096 | 0801683401 | 0.56 | $997.7\pm 7.6$ | 1001.5 | $31\pm 8.9$ | 4XMM J174954.6–294336 | $7.7\times 10^{-2}$ | 3 | IP | Mondal et al. (2023)
0.0036 | -1.9883 | 0801682901 | 0.76 | $2917.0\pm 68.5$ |  | $54.9\pm 12.7$ | 4XMM J175327.8–295716 | $1.8\times 10^{-1}$ | 2 | Likely Polar |
0.1413 | -0.1089 | 0762250301 | 0.32 | $1737.7\pm 5.4$ | 1745 | $15.8\pm 3.7$ | 4XMM J174622.7–285218 | $2.5\times 10^{-3}$ | 2 | IP | Muno et al. (2009)
0.1476 | -2.2564 | 0402280101 | 0.59 | $552.1\pm 1.4$ |  | $45.1\pm 9.8$ | 4XMM J175452.0–295758 | $9.9\times 10^{-2}$ | 3 | Unknown |
0.5849 | -1.5353 | 0801682801 | 0.49 | $672.5\pm 3.0$ |  | $77.8\pm 18.9$ | 4XMM J175301.3–291324 | $7.7\times 10^{-2}$ | 3 | Likely IP |
0.7407 | -0.4932 | 0801681301 | 0.44 | $1209.7\pm 11.7$ | 1212.4 | $44.3\pm 8.3$ | 4XMM J174917.7–283329 | $1.3\times 10^{-2}$ | 3 | IP | Mondal et al. (2023)
0.7632 | -1.3587 | 0801682601 | 0.97 | $6784.4\pm 439.5$ |  | $73\pm 20.4$ | 4XMM J175244.4–285851 | $2.8\times 10^{-1}$ | 3 | Likely Polar | Bahramian et al. (2021)
0.9918 | -0.0821 | 0783160101 | 0.47 | $593.6\pm 0.7$ | 593 | $41.3\pm 7.3$ | 4XMM J174816.9–280750 | $6.8\times 10^{-2}$ | 3 | IP | Sidoli et al. (2006)
1.3573 | 1.0522 | 0201200101 | 1.9 | $3195.8\pm 116.3$ | 3200 | $28.9\pm 6.5$ | 4XMM J174445.6–271344 | $6.2\times 10^{-1}$ | 3 | NS/WD ? | Lopes de Oliveira et al. (2006)
1.4196 | -2.2266 | 0782770201 | 0.51 | $354.68\pm 0.7$ |  | $47.6\pm 10.9$ | 4XMM J175740.5–285105 | $9.2\times 10^{-1}$ | 3 | Likely IP |
1.5909 | 0.0633 | 0510010401 | 1.9 | $132.1\pm 0.3$ | 132 | $28.9\pm 3.4$ | 4XMM J174906.8–273233 | $2.0\times 10^{-1}$ | 3 | NS HMXRB | Karasev et al. (2008)
2.6999 | -0.7212 | 0886081101 | 0.59 | $1.36652\pm 2\times 10^{-5}$ |  | $48.1\pm 12.6$ | XMMU J175441.9–265919 | $7.7\times 10^{-3}$ | 3 | Likely NS XRB |
3.5626 | -0.3453 | 0886081301 | 0.53 | $1134.5\pm 12.4$ |  | $79.1\pm 14.4$ | 4XMM J175511.6–260315 | $1.7\times 10^{-2}$ | 3 | Likely IP |
3.5766 | -0.3953 | 0886081301 | 0.51 | $392.5\pm 1.5$ |  | $28.2\pm 5$ | 4XMM J175525.0–260402 | $1.5\times 10^{-2}$ | 3 | Likely NS XRB |
4.4701 | 0.6376 | 0840910501 | 0.36 | $4093.3\pm 182.3$ |  | $34.8\pm 6.1$ | 4XMM J175328.4–244627 | $8.8\times 10^{-2}$ | 3 | Likely Polar |
6.3308 | -0.4448 | 0886110501 | 0.64 | $5215.9\pm 7.5$ |  | $61.8\pm 19.3$ | XMMU J180140.3–234221 | $1.7\times 10^{-2}$ | 2 | Likely Polar |

777The details of the X-ray pulsators, including the positional information,
XMM-Newton detection ID, and the source association in the 4XMM catalog. The
source XMMU J173029.8–330920, XMMU J175441.9–265919, and XMMU J180140.3–234221
are the first time detected by XMM-Newton in our ongoing _Heritage_ survey of
the Galactic disk. We also searched for Swift-XRT and Chandra counterparts of
these two sources, but no counterparts were found. The PGaia represents the
probability of spurious association with a Gaia source.

Table 3: Details of the pulsators for which more than one observation is
available.

XMM Name | ObsID | Date | Flux | Period | Exposure
---|---|---|---|---|---
erg s-1 cm-2 | ks
4XMM J172511.3–361657 | 0206380401 | 2004-03-21 | $(7.42\pm 0.04)\times 10^{-11}$ | $413.8\pm 3.7$ | 10.9
0405640201 | 2006-08-29 | $(3.94\pm 0.17)\times 10^{-13}$ |  | 22.9
0405640301 | 2006-08-31 | $(6.46\pm 0.04)\times 10^{-11}$ | $414.5\pm 3.8$ | 11.3
0405640401 | 2006-09-04 | $(2.55\pm 0.02)\times 10^{-11}$ | $414.4\pm 3.3$ | 12.5
0405640501 | 2006-09-06 | $(3.21\pm 0.08)\times 10^{-12}$ | $409.8\pm 3.5$ | 11.9
0405640601 | 2006-09-08 | $(5.53\pm 0.27)\times 10^{-13}$ |  | 13.9
0405640701 | 2006-09-15 | $(1.81\pm 0.02)\times 10^{-11}$ | $414.4\pm 1.7$ | 22.9
0405641001 | 2006-09-27 | $(2.19\pm 0.21)\times 10^{-13}$ |  | 12.4
0405640901 | 2006-09-28 | $(2.77\pm 0.02)\times 10^{-11}$ | $413.9\pm 2.6$ | 15.2
0405640801 | 2006-10-01 | $(4.12\pm 0.02)\times 10^{-11}$ | $413.2\pm 2.5$ | 15.7
| 0886070601 | 2006-10-01 | $(3.56\pm 0.03)\times 10^{-11}$ | $414.5\pm 1.4$ | 26.6
4XMM J174517.0–321358 | 0553950201 | 2010-10-09 | $(2.10\pm 0.06)\times 10^{-12}$ | $613.8\pm 0.9$ | 86.4
0870990201 | 2021-02-28 | $(1.24\pm 0.02)\times 10^{-12}$ | $614.2\pm 2.4$ | 31.6
0865510101 | 2021-03-02 | $(1.26\pm 0.02)\times 10^{-12}$ | $614.2\pm 1.4$ | 62.9
4XMM J174033.8–301501 | 0823030101 | 2018-09-29 | $(1.96\pm 0.09)\times 10^{-12}$ | $433.0\pm 5.6$ | 8.0
0886010601 | 2021-03-18 | $(2.88\pm 0.06)\times 10^{-12}$ | $432.4\pm 1.9$ | 23.0
4XMM J174016.0–290337 | 0304220101 | 2005-09-29 | $(3.61\pm 0.13)\times 10^{-12}$ | $624.8\pm 9.4$ | 8.5
0764191201 | 2016-03-05 | $(5.14\pm 0.06)\times 10^{-12}$ | $623.2\pm 2.6$ | 33
0764191101 | 2016-03-05 | $(8.78\pm 0.30)\times 10^{-12}$ | $622.6\pm 2.6$ | 33
4XMM J174809.8–300616 | 0152920101 | 2003-04-02 | $(2.96\pm 0.17)\times 10^{-13}$ | $2179.6\pm 18.2$ | 52.2
0801683301 | 2018-04-06 | $(2.08\pm 0.39)\times 10^{-13}$ |  | 29.8
4XMM J174009.0–284725 | 0511010701 | 2008-02-27 | $(3.75\pm 0.09)\times 10^{-12}$ | $733.1\pm 14.6$ | 9.3
0764191501 | 2016-02-25 | $(5.89\pm 0.22)\times 10^{-12}$ | $725.5\pm 3.7$ | 30.5
0764191101 | 2016-03-05 | $(5.90\pm 0.14)\times 10^{-12}$ | $725.9\pm 3.5$ | 33
0764191601 | 2016-03-10 | $(5.63\pm 0.21)\times 10^{-12}$ | $725.5\pm 6.0$ | 19
4XMM J174954.6–294336 | 0801681401 | 2017-10-07 | $(1.28\pm 0.05)\times 10^{-12}$ |  | 28
0801683401 | 2018-04-06 | $(1.73\pm 0.06)\times 10^{-12}$ | $997.7\pm 7.6$ | 29.2
4XMM J175327.8–295716 | 0085580501 | 2000-10-11 | $(3.35\pm 0.48)\times 10^{-13}$ |  | 8.0
0085581501 | 2001-03-24 | $(1.93\pm 1.22)\times 10^{-13}$ |  | 7.5
0085581601 | 2001-09-07 | $(3.77\pm 1.07)\times 10^{-13}$ |  | 8.2
0085581801 | 2002-03-13 | $(1.21\pm 0.53)\times 10^{-13}$ |  | 8.2
0801682901 | 2018-09-07 | $(4.29\pm 0.38)\times 10^{-13}$ | $2917.0\pm 68.5$ | 27.9
0801683601 | 2018-09-25 | $(3.05\pm 0.31)\times 10^{-13}$ | $2870.2\pm 58.1$ | 31.7
4XMM J175452.0–295758 | 0085580501 | 2000-10-11 | $(1.88\pm 0.45)\times 10^{-13}$ |  | 8.0
0206590201 | 2004-09-05 | $(2.46\pm 0.27)\times 10^{-13}$ |  | 20.9
0402280101 | 2006-09-10 | $(2.68\pm 0.19)\times 10^{-13}$ | $552.1\pm 1.4$ | 44.1
4XMM J175301.3–291324 | 0801682501 | 2018-09-03 | $(3.29\pm 1.29)\times 10^{-13}$ |  | 29.0
0801682801 | 2018-09-09 | $(1.86\pm 0.11)\times 10^{-13}$ | $673.1\pm 3.1$ | 32.9
0801683501 | 2018-09-25 | $(4.12\pm 0.57)\times 10^{-13}$ | $672.5\pm 3.0$ | 31.5

888The XMM-Newton ObsIDs details for sources in which more than one
observation is available. The flux is taken from the 4XMM catalog. In many
cases, the pulsation was not detected if the exposure was short or the source
flux was below a certain limit.

Table 4: Table 3 Continued.

XMM Name | ObsID | Date | Flux | Period | Exposure
---|---|---|---|---|---
4XMM J174917.7–283329 | 0410580401 | 2006-09-22 | $(1.85\pm 0.50)\times 10^{-13}$ |  | 32.9
0410580501 | 2006-09-26 | $(3.07\pm 0.63)\times 10^{-13}$ |  | 32.4
0801681301 | 2017-10-07 | $(1.32\pm 0.03)\times 10^{-12}$ | $1209.7\pm 11.7$ | 28.0
4XMM J174816.9–280750 | 0112970101 | 2000-09-23 | $(7.89\pm 0.43)\times 10^{-13}$ | $595.1\pm 5.5$ | 16.3
0112970201 | 2000-09-23 | $(9.34\pm 0.79)\times 10^{-13}$ | $587.8\pm 4.5$ | 18.1
0144220101 | 2003-03-12 | $(9.42\pm 0.61)\times 10^{-13}$ | $592.9\pm 1.4$ | 52.4
0205240101 | 2005-02-26 | $(6.37\pm 0.19)\times 10^{-13}$ | $592.7\pm 1.4$ | 51
0694640801 | 2012-10-06 | $(8.04\pm 0.37)\times 10^{-13}$ | $590.9\pm 1.7$ | 41.9
0694641501 | 2012-10-06 | $(6.67\pm 0.17)\times 10^{-13}$ | $592.9\pm 1.4$ | 51.8
0694640701 | 2012-10-02 | $(6.47\pm 0.19)\times 10^{-13}$ | $592.5\pm 1.7$ | 44.4
0694641401 | 2012-09-30 | $(7.43\pm 0.41)\times 10^{-13}$ |  | 51.8
0783160101 | 2016-10-02 | $(8.51\pm 0.23)\times 10^{-13}$ | $593.6\pm 0.7$ | 106
0862471201 | 2020-10-04 | $(1.01\pm 0.05)\times 10^{-12}$ |  | 46.9
4XMM J174445.6–271344 | 0201200101 | 2004-02-26 | $(2.03\pm 0.04)\times 10^{-12}$ | $3195.8\pm 116.3$ | 17.8
0691760101 | 2012-09-08 | $(1.33\pm 0.02)\times 10^{-12}$ |  | 22.9
4XMM J174906.8–273233 | 0510010401 | 2007-03-31 | $(1.13\pm 0.01)\times 10^{-11}$ | $132.1\pm 0.3$ | 12.2
0511010301 | 2008-03-04 | $(2.44\pm 0.12)\times 10^{-12}$ |  | 8.9
4XMM J175511.6–260315 | 0148090101 | 2003-03-17 | $(1.07\pm 0.14)\times 10^{-12}$ |  | 12.1
0148090501 | 2003-09-11 | $(1.53\pm 0.17)\times 10^{-12}$ |  | 11.2
0886081301 | 2023-04-06 | $(3.06\pm 0.21)\times 10^{-12}$ | $1134.5\pm 12.4$ | 24
4XMM J175525.0–260402 | 0148090101 | 2003-03-17 | $(3.03\pm 0.30)\times 10^{-12}$ |  | 12.1
0148090501 | 2003-09-11 | $(4.47\pm 0.43)\times 10^{-12}$ |  | 11.2
0886081301 | 2023-04-06 | $(1.84\pm 0.10)\times 10^{-12}$ | $392.5\pm 1.5$ | 24

999Same columns as Table 3.

Table 5: Details of the spectral fit.

XMM Name | $N_{\rm H}$ | $\Gamma$ | $N_{\rm po}$ | $N_{\rm 6.4}$ | $\rm EW_{6.4}$ | $N_{\rm 6.7}$ | $\rm EW_{6.7}$ | $N_{\rm 6.9}$ | $\rm EW_{6.9}$ | $\chi^{2}/\rm d.o.f$ | Flux
---|---|---|---|---|---|---|---|---|---|---|---
$\times 10^{22}\rm\ cm^{-2}$ | eV | eV | eV | 0.2–10 keV
4XMM J172511.3 | $10.2^{+0.3}_{-0.3}$ | $0.82^{+0.04}_{-0.04}$ | $3.1^{+0.3}_{-0.3}\times 10^{-3}$ | $3.9^{+0.6}_{-0.6}\times 10^{-5}$ | $70^{+6}_{-7}$ | $<5\times 10^{-6}$ | $<8$ | $<3\times 10^{-6}$ | $<7$ | 2364/2299 | $3.56\times 10^{-11}$
4XMM J173058.9 | $1.4^{+0.6}_{-0.4}$ | $0.6^{+0.3}_{-0.2}$ | $5^{+3}_{-2}\times 10^{-5}$ | $3^{+3}_{-2}\times 10^{-6}$ | $129^{+70}_{-58}$ | $5^{+3}_{-3}\times 10^{-6}$ | $236^{+105}_{-83}$ | $<1\times 10^{-6}$ | $<50$ | 69/82 | $1.52\times 10^{-12}$
XMMU J173029.8 | $5^{+5}_{-4}$ | $0.0^{+0.7}_{-0.6}$ | $9^{+28}_{-6}\times 10^{-6}$ | $2^{+1}_{-1}\times 10^{-6}$ | $270^{+135}_{-135}$ | $<2\times 10^{-6}$ | $<160$ | $2^{+3}_{-1}\times 10^{-6}$ | $204^{+306}_{-102}$ | 11/17 | $6.70\times 10^{-13}$
4XMM J174517.0 | $2.8^{+0.3}_{-0.3}$ | $1.1^{+0.1}_{-0.1}$ | $9^{+2}_{-2}\times 10^{-5}$ | $2.8^{+0.8}_{-0.8}\times 10^{-6}$ | $192_{-31}^{+32}$ | $1.4^{+0.8}_{-0.8}\times 10^{-6}$ | $78^{+34}_{-19}$ | $1.7^{+0.8}_{-0.8}\times 10^{-6}$ | $129^{+51}_{-35}$ | 339/326 | $9.59\times 10^{-13}$
4XMM J174728.9 | $0.2^{+0.1}_{-0.1}$ | $1.8^{+0.4}_{-0.4}$ | $7^{+4}_{-2}\times 10^{-6}$ | $<4\times 10^{-7}$ | $<1500$ | $<5\times 10^{-7}$ | $<2000$ | $<5\times 10^{-7}$ | $<2200$ | 16/28 | $4.00\times 10^{-14}$
4XMM J173837.0 | $0.9^{+0.7}_{-0.5}$ | $0.8^{+0.4}_{-0.4}$ | $1.1^{+1.1}_{-0.5}\times 10^{-5}$ | $<8\times 10^{-7}$ | $<351$ | $7^{+91}_{-7}\times 10^{-8}$ | $29^{+329}_{-29}$ | $1^{+9}_{-1}\times 10^{-7}$ | $59^{+320}_{-59}$ | 23/42 | $2.06\times 10^{-13}$
4XMM J174033.8 | $0.9^{+0.2}_{-0.2}$ | $0.5^{+0.1}_{-0.1}$ | $6^{+1}_{-1}\times 10^{-5}$ | $7^{+2}_{-2}\times 10^{-6}$ | $226^{+50}_{-46}$ | $4^{+2}_{-2}\times 10^{-6}$ | $103^{+42}_{-40}$ | $3^{+2}_{-2}\times 10^{-6}$ | $95^{+50}_{-49}$ | 185/189 | $2.88\times 10^{-12}$
4XMM J174016.0 | $0.44^{+0.03}_{-0.03}$ | $0.63^{+0.04}_{-0.02}$ | $1.36^{+0.09}_{-0.07}\times 10^{-4}$ | $9^{+2}_{-2}\times 10^{-6}$ | $182^{+34}_{-20}$ | $6^{+2}_{-2}\times 10^{-6}$ | $91^{+23}_{-16}$ | $7^{+2}_{-2}\times 10^{-6}$ | $135^{+27}_{-29}$ | 736/607 | $5.14\times 10^{-12}$
4XMM J174809.8 | $2^{+2}_{-1}$ | $0.3^{+0.5}_{-0.5}$ | $6^{+9}_{-3}\times 10^{-6}$ | $5^{+7}\times 10^{-7}$ | $175^{+71}$ | $<7\times 10^{-7}$ | $<172$ | $3^{+14}_{-3}\times 10^{-7}$ | $99^{+264}_{-99}$ | 50/48 | $2.96\times 10^{-13}$
4XMM J174009.0 | $0.8^{+0.3}_{-0.2}$ | $0.1^{+0.1}_{-0.1}$ | $6^{+2}_{-1}\times 10^{-5}$ | $1.8^{+0.5}_{-0.5}\times 10^{-5}$ | $359^{+65}_{-120}$ | $1.6^{+0.6}_{-0.6}\times 10^{-5}$ | $265^{+78}_{-115}$ | $7^{+6}_{-6}\times 01^{-6}$ | $109^{+66}_{-56}$ | 170/166 | $5.90\times 10^{-12}$
4XMM J174954.6 | $2.4^{+0.9}_{-0.7}$ | $0.4^{+0.3}_{-0.3}$ | $3^{+2}_{-1}\times 10^{-5}$ | $2^{+1}_{-1}\times 10^{-6}$ | $107^{+69}_{-59}$ | $3^{+2}_{-2}\times 10^{-6}$ | $172^{+103}_{-85}$ | $<3\times 10^{-6}$ | $<150$ | 90/91 | $1.73\times 10^{-12}$
4XMM J175327.8 | $7^{+7}_{-5}\times 10^{-2}$ | $1.3^{+0.3}_{-0.3}$ | $2.8^{+0.9}_{-0.7}\times 10^{-5}$ | $<2\times 10^{-6}$ | $<774$ | $<2\times 10^{-6}$ | $<660$ | $<2\times 10^{-6}$ | $<751$ | 64/41 | $4.29\times 10^{-13}$
4XMM J174622.7 | $4.3^{+0.9}_{-0.7}$ | $0.7^{+0.2}_{-0.2}$ | $2.0^{+1.0}_{-0.6}\times 10^{-5}$ | $1^{+5}_{-1}\times 10^{-7}$ | $20^{+46}_{-20}$ | $1.1^{+0.6}_{-0.6}\times 10^{-7}$ | $207^{+76}_{-71}$ | $2^{+6}_{-2}\times 10^{-7}$ | $30^{+67}_{-30}$ | 192/187 | $5.13\times 10^{-13}$
4XMM J175452.0 | $2^{+3}_{-1}$ | $0.9^{+0.9}_{-0.7}$ | $1.2^{+4.4}_{-0.9}\times 10^{-5}$ | $1.4^{+0.7}_{-0.7}\times 10^{-6}$ | $690^{+445}_{-445}$ | $<9\times 10^{-7}$ | $<206$ | $<1\times 10^{-6}$ | $<529$ | 33/41 | $2.68\times 10^{-13}$
4XMM J175301.3 | $0.3^{+0.5}_{-0.2}$ | $0.5^{+0.3}_{-0.3}$ | $5^{+3}_{-2}\times 10^{-6}$ | $10^{+6}_{-6}\times 10^{-7}$ | $521^{+288}_{-225}$ | $4^{+8}_{-4}\times 10^{-7}$ | $105^{+189}_{-105}$ | $4^{+10}_{-4}\times 10^{-7}$ | $179^{+387}_{-179}$ | 63/54 | $1.86\times 10^{-13}$
4XMM J174917.7 | $3.2^{+0.7}_{-0.6}$ | $1.0^{+0.2}_{-0.2}$ | $9^{+5}_{-3}\times 10^{-5}$ | $2^{+1}_{-1}\times 10^{-6}$ | $94^{+50}_{-47}$ | $4^{+2}_{-2}\times 10^{-6}$ | $246^{+88}_{-67}$ | $9^{+16}_{-9}\times 10^{-7}$ | $52^{+112}_{-52}$ | 107/108 | $1.31\times 10^{-12}$
4XMM J175244.4 | $0.1^{+0.2}_{-0.1}$ | $0.7^{+0.3}_{-0.3}$ | $7^{+3}_{-2}\times 10^{-6}$ | $<8\times 10^{-7}$ | $<388$ | $6^{+95}_{-6}\times 10^{-8}$ | $17^{+231}_{-17}$ | $1.4^{+0.9}_{-0.9}\times 10^{-6}$ | $796^{+443}_{-405}$ | 29/51 | $2.11\times 10^{-13}$
4XMM J174816.9 | $27^{+7}_{-6}$ | $1.3^{+0.6}_{-0.6}$ | $2.0^{+4}_{-1}\times 10^{-4}$ | $1.8^{+0.7}_{-0.7}\times 10^{-6}$ | $221^{+86}_{-86}$ | $1.0^{+0.7}_{-0.7}\times 10^{-6}$ | $109^{+76}_{-76}$ | $6^{+9}_{-6}\times 10^{-7}$ | $71^{+107}_{-71}$ | 159/164 | $8.50\times 10^{-13}$
4XMM J174445.6 | $0.37^{+0.02}_{-0.02}$ | $1.75^{+0.06}_{-0.06}$ | $3.3^{+0.2}_{-0.2}\times 10^{-4}$ | $2^{+1}_{-1}\times 10^{-6}$ | $80^{+39}_{-40}$ | $7^{+2}_{-2}\times 10^{-6}$ | $371^{+88}_{-76}$ | $2^{+1}_{-1}\times 10^{-6}$ | $109^{+43}_{-47}$ | 399/412 | $2.03\times 10^{-12}$
4XMM J175740.5 | $3^{+2}_{-1}$ | $0.2^{+0.4}_{-0.3}$ | $1.1^{+1.1}_{-0.5}\times 10^{-5}$ | $1.6^{+0.8}_{-0.8}\times 10^{-6}$ | $176^{+84}_{-49}$ | $1.3^{+0.9}_{-0.9}\times 10^{-6}$ | $112^{+61}_{-41}$ | $1.3^{+1.0}_{-1.0}\times 10^{-6}$ | $151^{+19}_{-64}$ | 113/92 | $7.63\times 10^{-13}$
4XMM J174906.8 | $21^{+1}_{-1}$ | $1.3^{+0.1}_{-0.1}$ | $2.6^{+0.8}_{-0.6}\times 10^{-3}$ | $6^{+4}_{-4}\times 10^{-6}$ | $40^{+15}_{-15}$ | $3^{+40}_{-3}\times 10^{-7}$ | $2^{+13}_{-2}$ | $3^{+44}_{-3}\times 10^{-7}$ | $2^{+16}_{-2}$ | 400/442 | $1.12\times 10^{-11}$
XMMU J175441.9 | $3^{+2}_{-1}$ | $0.3^{+0.4}_{-0.4}$ | $1.1^{+1.3}_{-0.6}\times 10^{-5}$ | $2^{+1}_{-1}\times 10^{-6}$ | $274^{+137}_{-137}$ | $3^{+14}_{-3}\times 10^{-7}$ | $28^{+131}_{-28}$ | $1^{+1}_{-1}\times 10^{-6}$ | $172^{+172}_{-172}$ | 33/26 | $4.31\times 10^{-13}$
4XMM J175511.6 | $0.35^{+0.11}_{-0.09}$ | $1.1^{+0.1}_{-0.1}$ | $2.8^{+0.6}_{-0.5}\times 10^{-5}$ | $1^{+6}_{-1}\times 10^{-7}$ | $31^{+109}_{-31}$ | $6^{+7}_{-6}\times 10^{-7}$ | $194^{+162}_{-188}$ | $<1\times 10^{-6}$ | $<263$ | 69/76 | $3.06\times 10^{-13}$
4XMM J175525.0 | $10^{+1}_{-1}$ | $1.4^{+0.2}_{-0.2}$ | $5^{+2}_{-2}\times 10^{-4}$ | $<1\times 10^{-6}$ | $<33$ | $9^{+16}_{-9}\times 10^{-7}$ | $34^{+36}_{-34}$ | $<1\times 10^{-6}$ | $<55$ | 171/170 | $1.84\times 10^{-12}$
4XMM J175328.4 | $0.43^{+0.03}_{-0.02}$ | $1.40^{+0.04}_{-0.04}$ | $6.3^{+0.4}_{-0.3}\times 10^{-4}$ | $3^{+2}_{-2}\times 10^{-6}$ | $71^{+28}_{-29}$ | $2^{+2}_{-2}\times 10^{-6}$ | $28^{+26}_{-28}$ | $4^{+3}_{-3}\times 10^{-6}$ | $90^{+38}_{-36}$ | 883/914 | $5.56\times 10^{-12}$
XMMU J180140.3 | $3^{+2}_{-1}$ | $2.1^{+1.3}_{-0.9}$ | $1.1^{+5.0}_{-0.8}\times 10^{-4}$ | $2^{+3}_{-1}\times 10^{-6}$ | $670^{+1005}_{-335}$ | $<4\times 10^{-8}$ | $<8$ | $<6\times 10^{-6}$ | $<1500$ | 25/27 | $2.31\times 10^{-13}$

101010The model used for spectral fit is power-law for the continuum plus
three Gaussian at 6.4, 6.7, and 6.9 keV for iron emission complex. The
spectral model is convolved with a Galactic absorption component tbabs. The
normalization of power-law model $N_{\rm po}$ is given in units of photons
keV-1 cm2 s-1 at 1 keV and the normalization of the lines are given in units
of photons cm2 s-1 in the line energy.

Figure 8: Lomb-Scargle periodogram of the sources listed in Table 2. The
periodograms are constructed using 2–10 keV EPIC-pn light curves. The
horizontal green and red lines indicate the $2\sigma$ and $3\sigma$ confidence
levels, respectively, computed from simulations, and the blue line indicates
the false alarm probability ($3\sigma$ confidence level) estimated from the
analytical approximation from Baluev (2008). The small inset shows the folded
light curve.

Figure 9: Fig. 8 Continued.

Figure 10: Spectral modeling of the sources in our sample using a model
composed of tbabs*(power-law+g1+g2+g3). The g1,g2, and g3 represent three
Gaussian lines, at 6.4, 6.7, and 6.9 keV, respectively. The black, red, and
green colors represent data from the EPIC-pn, MOS1, and MOS2 detectors,
respectively.

Figure 11: Fig. 10 Continued.