# Einstein Probe discovery of a super-soft outburst from CXOU J005245.0-722844: a rare BeWD binary in the Small Magellanic Cloud

A Marino, H Yang, F. Coti Zelati, N Rea, S Guillot, G.K Jaisawal, C Maitra, F Haberl, E Kuulkers, W Yuan, et al.

► To cite this version:

A Marino, H Yang, F. Coti Zelati, N Rea, S Guillot, et al. Einstein Probe discovery of a super-soft outburst from CXOU J005245.0-722844: a rare BeWD binary in the Small Magellanic Cloud. The Astrophysical Journal Letters, 2025, 980 (2), pp.L36. 10.3847/2041-8213/ad9580. hal-04679966

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# Einstein Probe Discovery of EP J005245.1–722843: A Rare Be–White Dwarf Binary in the Small Magellanic Cloud?

A. Marino<sup>1,2</sup> <sup>1</sup>&Dagger;, H. N. Yang<sup>3,4,5</sup> <sup>1</sup>&Dagger;, F. Coti Zelati<sup>1,2</sup> <sup>1</sup>&Dagger;, N. Rea<sup>1,2</sup> <sup>1</sup>&Dagger;, S. Guillot<sup>6</sup> <sup>1</sup>&Dagger;, G. K. Jaisawal<sup>7</sup> <sup>1</sup>&Dagger;, C. Maitra<sup>4</sup> <sup>1</sup>&Dagger;, J.-U. Ness<sup>8</sup> <sup>1</sup>&Dagger;, F. Haberl<sup>4</sup> <sup>1</sup>&Dagger;, E. Kuulkers<sup>9</sup> <sup>1</sup>&Dagger;, W. Yuan<sup>3,5</sup>, H. Feng<sup>10</sup>, L. Tao<sup>10</sup> <sup>1</sup>&Dagger;, C. Jin<sup>3,5</sup> <sup>1</sup>&Dagger;, H. Sun<sup>3</sup> <sup>1</sup>&Dagger;, W. Zhang<sup>3</sup>, W. Chen<sup>3,5</sup>, E. P. J. van den Heuvel<sup>11</sup> <sup>1</sup>&Dagger;, R. Soria<sup>12,13,14</sup> <sup>1</sup>&Dagger;, B. Zhang<sup>15,16</sup> <sup>1</sup>&Dagger;, S.-S. Weng<sup>17</sup> <sup>1</sup>&Dagger;, L. Ji<sup>18</sup> <sup>1</sup>&Dagger;, G. B. Zhang<sup>19</sup> <sup>1</sup>&Dagger;, X. Pan<sup>3</sup> <sup>1</sup>&Dagger;, Z. Lv<sup>3,5</sup>, C. Zhang<sup>3,5</sup>, Z. X. Ling<sup>3,5</sup>, Y. Chen<sup>10</sup> <sup>1</sup>&Dagger;, S. Jia<sup>5,10</sup> <sup>1</sup>&Dagger;, Y. Liu<sup>3</sup>, H. Q. Cheng<sup>3</sup>, D. Y. Li<sup>3</sup> <sup>1</sup>&Dagger;, K. Gendreau<sup>20</sup> <sup>1</sup>&Dagger;, M. Ng<sup>21</sup> <sup>1</sup>&Dagger;, and T. Strohmayer<sup>20,22</sup> <sup>1</sup>&Dagger;

<sup>1</sup> Institute of Space Sciences (ICE, CSIC), Campus UAB, Carrer de Can Magrans s/n, 08193, Barcelona, Spain; marino@ice.csic.es

<sup>2</sup> Institut d'Estudis Espacials de Catalunya (IEEC), Esteve Terradas 1, RDIT Building, Of. 212 Mediterranean Technology Park (PMT), 08860, Castelldefels, Spain

<sup>3</sup> National Astronomical Observatories, Chinese Academy of Sciences, Beijing, 100101, People's Republic of China; hnyang@nao.cas.cn

<sup>4</sup> Max-Planck-Institut für extraterrestrische Physik, Gießenbachstraße 1, D-85748 Garching bei München, Germany

<sup>5</sup> School of Astronomy and Space Sciences, University of Chinese Academy of Sciences, Beijing, 100049, People's Republic of China

<sup>6</sup> IRAP, CNRS, 9 avenue du Colonel Roche, BP 44346, 31028 Toulouse Cedex 4, France

<sup>7</sup> DTU Space, Technical University of Denmark, Elektrovej 327-328, DK-2800 Lyngby, Denmark

<sup>8</sup> European Space Agency (ESA), European Space Astronomy Centre (ESAC), Camino Bajo del Castillo s/n, E-28692 Villanueva de la Cañada, Madrid, Spain

<sup>9</sup>European Space Agency (ESA), European Space Research and Technology Centre (ESTEC), Keplerlaan 1, 2201 AZ Noordwijk, The Netherlands

<sup>10</sup> Key Laboratory of Particle Astrophysics, Institute of High Energy Physics, Chinese Academy of Sciences, Beijing 100049, People's Republic of China

<sup>11</sup> Anton Pannekoek Institute for Astronomy, University of Amsterdam, P.O. Box 94249, 1090 GE Amsterdam, The Netherlands

<sup>12</sup> INAF—Osservatorio Astrofisico di Torino, Strada Osservatorio 20, I-10025 Pino Torinese, Italy

<sup>13</sup> Sydney Institute for Astronomy, School of Physics A28, The University of Sydney, Sydney, NSW 2006, Australia

<sup>14</sup> College of Astronomy and Space Sciences, University of the Chinese Academy of Sciences, Beijing 100049, People's Republic of China

<sup>15</sup> Nevada Center for Astrophysics, University of Nevada Las Vegas, NV 89154, USA

<sup>16</sup> Department of Physics and Astronomy, University of Nevada Las Vegas, NV 89154, USA

<sup>17</sup> Department of Physics and Institute of Theoretical Physics, Nanjing Normal University, Nanjing, People's Republic of China

<sup>18</sup> School of Physics and Astronomy, Sun Yat-sen University, Zhuhai, 519082, People's Republic of China

<sup>19</sup> Yunnan Observatories, Chinese Academy of Sciences, Kunming 650216, People's Republic of China

<sup>20</sup> Astrophysics Science Division, NASA's Goddard Space Flight Center, Greenbelt, MD 20771, USA

<sup>21</sup> MIT Kavli Institute for Astrophysics and Space Research, Massachusetts Institute of Technology, Cambridge, MA 02139, USA

<sup>22</sup> Joint Space-Science Institute, NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA

Received 2024 July 29; revised 2024 November 19; accepted 2024 November 20; published 2025 February 18

## Abstract

On 2024 May 27, the Wide-field X-ray Telescope on board the Space Sciences, University of Chinese Academy of Einstein Probe (EP) mission detected enhanced X-ray emission from a new transient source in the Small Magellanic Cloud during its commissioning phase. Prompt follow-up with the EP Follow-up X-ray Telescope, the Swift X-ray Telescope. and NICER have revealed a very soft, thermally emitting source ($kT$ ~ 0.1 keV at the outburst peak) with an X-ray luminosity of $L \sim 4 \times 10^{38}$ erg s$^{-1}$, labeled EP J005245.1–722843. This supersoft outburst faded very quickly in a week's time. Several emission lines and absorption edges were present in the X-ray spectrum, including deep nitrogen (0.67 keV) and oxygen (0.87 keV) absorption edges. The X-ray emission resembles the supersoft source phase of typical nova outbursts from an accreting white dwarf (WD) in a binary system, despite the X-ray source being historically associated with an O9-B0e massive star exhibiting a 17.55 day periodicity in the optical band. The discovery of this supersoft outburst suggests that EP J005245.1–722843 is a BeWD X-ray binary: an elusive evolutionary stage where two main-sequence massive stars have undergone a common envelope phase and experienced at least two episodes of mass transfer. In addition, the very short duration of the outburst and the presence of Ne features hint at a rather massive, i.e., close to the Chandrasekhar limit, Ne–O WD in the system.

Unified Astronomy Thesaurus concepts: Accretion (14); X-ray binary stars (1811); High mass x-ray binary stars (733); White dwarf stars (1799); Small Magellanic Cloud (1468)

## 1. Introduction

Be–white dwarf X-ray binary systems (BeWDs) are a subclass of X-ray binaries where a white dwarf (WD) accretes matter from a Be-type main-sequence star. Binary evolution models predict that BeWDs should be about 7 times more common than Be–neutron star systems (N. V. Raguzova 2001). Population synthesis models (C.-H. Zhu et al. 2023) indicate

that BeWDs originate from binaries initially composed of a Be star and a subdwarf O or B star. A substantial fraction of these systems can evolve either into red giants (via merging of the WD with a nondegenerated star; 60%–70% of the cases) or into double WDs detectable by the LISA mission as sources of gravitational wave emission (30%–40%). However, BeWDs have been so far relatively elusive. On the one hand, $\gamma$ Cas analogs, a class of Be stars showing an excess in hard X-rays, are largely considered to be BeWD binaries, although the matter is still controversial (see, e.g., D. R. Gies et al. 2023, and references therein). On the other hand, a few confirmed Be binaries, which show instead rather soft X-ray outbursts, have

<figure><img src="image_3.png" /></figure>

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also been identified as BeWDs (P. Kahabka et al. 2006; R. Sturm et al. 2012; M. J. Coe et al. 2020; F. Haberl et al. 2020; J. A. Kennea et al. 2021). The emission from the latter sources in outburst is typically characterized by very soft X-ray spectra with no detectable X-ray emission above >2 keV, indicative of nuclear burning on the WD surface. In this regime, BeWDs can be classified as supersoft sources (SSSs; V. Cracco et al. 2018), a group of sources that includes various types of X-ray systems powered by thermonuclear burning onto the surface of an accreting WD (J. Greiner et al. 1991). This category also encompasses novae during their SSS phase (e.g., L. Chomiuk et al. 2021). The lack of sensitive wide-field soft X-ray monitors has made the detection of outbursts from BeWDs rather challenging, especially during the rising phase. On 2024 May 27, the Wide-field X-ray Telescope (WXT) on board the Einstein Probe (EP) mission detected enhanced X-ray emission from a new transient source, EP J005245.1−722843 (J0052 hereafter), in the Small Magellanic Cloud (SMC) during its commissioning phase (H. N. Yang et al. 2024). This source was promptly identified as CXOU J005245.0−722844, following a serendipitous 52 s observation of the field by the Neil Gehrels Swift Observatory (Swift) X-ray Telescope (XRT) a few hours later (T. M. Gaudin et al. 2024; J. A. Kennea et al. 2024), as part of the Swift Small Magellanic Cloud Survey (J. A. Kennea et al. 2018). Two follow-up NICER observations on May 28 and 29 revealed a rapid decrease in the X-ray count rate between the two observations (G. K. Jaisawal et al. 2024). The source has a bright optical counterpart (average magnitude of $G \sim 14.9$), which has been first classified as an O9-B0e star by V. Antoniou et al. (2009) and later suggested to belong to the B0 class based also on its spectrum (H. A. Sheets et al. 2013; G. Maravelias et al. 2014; T. M. Gaudin et al. 2024). The companion star has an estimated age of ~40 Myr, according to the age of the OB population in that region of the SMC (V. Antoniou et al. 2019). The long-term optical light curve shows a periodicity at ~17.55 days, which is likely related to the orbital period of the system (I. Sarraj et al. 2012; H. Treiber et al. 2024). This Letter reports on the classification of this source as a BeWD following the discovery of its supersoft outburst by EP and its X-ray evolution. In the following, we adopt the position listed in the Chandra Source Catalog Release 2.1 (I. N. Evans et al. 2010): R.A. = 00h52m45s.13, decl. = −72°28′43″87 (J2000.0), and a distance of 62 kpc (D. Graczyk et al. 2020). All parameter uncertainties are quoted at a 1σ confidence level unless otherwise stated.

## 2. Observations and Data Reduction

In this Letter, we have used observations from various X-ray telescopes, including the WXT and Follow-up X-ray Telescope (FXT) on board EP, Swift, NICER, and XMM-Newton, as well as archival Chandra observations (see Table 1 for details).

### 2.1. EP

As a space observatory to discover and characterize X-ray transients, EP (W. Yuan et al. 2022) carries two instruments, the WXT and FXT. EP was launched on 2024 January 9 into a low-Earth circular orbit of 592 km height with an orbital period of 97 minutes. The WXT employs Lobster-eye micropore optics (MPO) to provide an instantaneous field of view of ~3600 square degrees. The instrument comprises 12 modules, each covering about 300 square degrees. Each module includes a mirror

assembly with 36 MPO, an optical baffle, and a focal-plane detector array of CMOS sensors. The in-orbit calibrations of WXT have been completed, whose results are in excellent agreement with those of the on-ground calibrations (H. Cheng 2024, in preparation). J0052 was simultaneously detected by the CMOS detectors 14 and 37 of the WXT during observations on 2024 May 27 from 08:41:28 to 12:47:16 UTC. Data were processed using the *wxtpipeline* tool version 0.1.0. The source region was determined as a circle with a radius of 67 pixels (1 pixel ~ 0.136") centered on the source. In the field of view of WXT, there are multiple bright objects together with their crossarms. Therefore, we carefully selected a circle of radius 268 pixels nearby as the background region to minimize the influence of other bright objects.

The FXT was configured in PW mode and FF mode for the observations performed on May 28 and June 4, respectively. Data were processed using the *fxtchain* tool available within the FXT Data Analysis Software. Source and background photons were extracted using a circle of radius 70" centered on the source and a close-by circle of radius 128", respectively.

### 2.2. Swift

The XRT observations were processed using *xrtpipe-line*, employing the latest version of CALDB (v.20240522). We extracted photons using a 20" circular region centered on the source. In the first observation (WT mode), we estimated the background with a similar-sized region far from the source. In the second observation (PC mode), we used an annular region with inner and outer radii of 45" and 150" for background estimation. We then extracted light curves and spectra through the *xrtproducts* pipeline. For the last four observations, we combined the data (totaling 8 ks of exposure) and set a 3σ upper limit on the net count rate of 0.002 counts s⁻¹ (calculated over the 0.3–10 keV energy range using the *sosta* command). In the first two Swift observations considered in this work, the UVOT was operated in image mode, using the UVW1 filter in the first one and cycling through all available filters in the second one. In the last four observations, the UVOT was operated in event mode using the UVM2 filter. We used a circle with a radius of 5" radius for the source and a nearby circle with a radius of 10" devoid of optical sources for the background. For image-mode data, we performed aperture photometry (properly corrected to infinite aperture) with *uvotsource*. The optical brightness in the UVM2 filter varied from ~12.6 to 13.2 mag (Vega) between May 27 and June 19, without showing any clear monotonic trend. For event-mode data, we processed the files with *coordinator* and *uvotscreen* and searched for periodic signals using the unbinned event files (see Section 3.4).

### 2.3. NICER

NICER data were processed using the *nicer12* pipeline with default screening settings (following the same prescriptions as in, e.g., A. Marino et al. 2023). Several NICER light curves showed irregular count rate increases during the outburst. To minimize charged particle contamination, we reprocessed the data with *nicer12*, setting the *over-only_range* to 0–5 and *COR_RANGE* to >1.5 GeV/c. Any observation segment still affected by non-X-ray flares after this procedure was neglected from the following analysis. The *nicer13-spect* pipeline was used to produce spectra and
<table><caption>**Table 1**
Journal of the 2024 X-Ray Observations of J0052 Presented in this Work</caption><thead><tr><th>Telescope/Instrument (Setup)</th><th>Obs ID</th><th>Start—End Time (UTC)<br>(Mmm DD hh:mm:ss)</th><th>Exposure<br>(ks)</th><th>Epochs<sup>a</sup></th></tr></thead><tbody><tr><td>EP/WXT</td><td>13600006246</td><td>May 25 11:06:46–May 27 06:28:12</td><td>67.9</td><td></td></tr><tr><td>EP/WXT</td><td>13600006247</td><td>May 27 08:02:48–May 27 09:39:06</td><td>2.0</td><td></td></tr><tr><td>EP/WXT</td><td>13600006248</td><td>May 27 09:39:06–May 29 03:22:17</td><td>63.0</td><td></td></tr><tr><td>Swift/XRT (PC)</td><td>03111173076</td><td>May 27 22:29:57–May 27 22:30:50</td><td>0.052</td><td></td></tr><tr><td>EP/WXT</td><td>08500000102</td><td>May 28 04:51:10–May 28 06:26:53</td><td>1.6</td><td></td></tr><tr><td>EP/FXT (PW)</td><td>08500000102</td><td>May 28 05:31:05–May 28 06:23:21</td><td>1.5</td><td>1</td></tr><tr><td>Swift/XRT (WT)</td><td>00033745002</td><td>May 28 15:40:13–May 28 17:19:37</td><td>1.9</td><td>2</td></tr><tr><td>NICER/XTI</td><td>7204550101</td><td>May 28 16:56:06–May 28 23:24:20</td><td>4.1</td><td>3,4,5,6</td></tr><tr><td>NICER/XTI</td><td>7204550102</td><td>May 29 11:26:10–May 30 00:00:00</td><td>6.0</td><td>7,8,9,10,11,12</td></tr><tr><td>NICER/XTI</td><td>7204550103</td><td>May 30 01:21:10–May 30 23:22:00</td><td>8.0</td><td>13,14,15,16</td></tr><tr><td>NICER/XTI</td><td>7204550104</td><td>May 31 00:33:30–May 31 21:01:21</td><td>6.4</td><td>17,18</td></tr><tr><td>NICER/XTI</td><td>7204550105</td><td>Jun 1 04:32:49–Jun 1 17:16:20</td><td>1.9</td><td>19</td></tr><tr><td>NICER/XTI</td><td>7204550106</td><td>Jun 2 05:09:40–Jun 2 17:59:00</td><td>5.8</td><td>20</td></tr><tr><td>NICER/XTI</td><td>7204550107</td><td>Jun 2 23:46:40–Jun 3 23:09:40</td><td>3.4</td><td>21</td></tr><tr><td>NICER/XTI</td><td>7204550108</td><td>Jun 4 00:50:20–Jun 4 23:59:20</td><td>1.9</td><td></td></tr><tr><td>EP/FXT (FF)</td><td>08500000107</td><td>Jun 4 12:15:38–Jun 4 16:19:28</td><td>9.1</td><td></td></tr><tr><td>NICER/XTI</td><td>7204550109</td><td>Jun 5 10:33:20–Jun 5 13:49:20</td><td>1.3</td><td></td></tr><tr><td>NICER/XTI</td><td>7204550110</td><td>Jun 6 02:24:40–Jun 6 22:23:00</td><td>0.8</td><td></td></tr><tr><td>NICER/XTI</td><td>7204550111</td><td>Jun 6 23:43:40–Jun 7 21:49:40</td><td>8.5</td><td></td></tr><tr><td>NICER/XTI</td><td>7204550112</td><td>Jun 8 10:14:20–Jun 8 18:03:00</td><td>0.3</td><td></td></tr><tr><td>Swift/XRT (PC)</td><td>00033745004</td><td>Jun 13 20:35:38–Jun 13 22:32:52</td><td>2.5</td><td></td></tr><tr><td>XMM-Newton/EPIC-pn (SW)</td><td>0935191301</td><td>Jun 14 19:12:05–Jun 15 03:59:07</td><td>31.6</td><td></td></tr><tr><td>XMM-Newton/EPIC-MOS1 (T)</td><td>0935191301</td><td>Jun 14 19:26:34–Jun 15 03:54:49</td><td>27.4</td><td></td></tr><tr><td>XMM-Newton/EPIC-MOS2 (T)</td><td>0935191301</td><td>Jun 14 19:26:51–Jun 15 03:54:39</td><td>29.4</td><td></td></tr><tr><td>Swift/XRT (PC)</td><td>00033745005</td><td>Jun 15 02:31:09–Jun 15 21:52:53</td><td>1.6</td><td></td></tr><tr><td>Swift/XRT (PC)</td><td>00033745006</td><td>Jun 17 19:09:39–Jun 17 21:12:52</td><td>2.6</td><td></td></tr><tr><td>Swift/XRT (PC)</td><td>00033745007</td><td>Jun 19 18:35:06–Jun 19 18:56:53</td><td>1.3</td><td></td></tr></tbody></table>

Notes. FF: full frame; PW: partial window; PC: photon counting; WT: windowed timing; SW: small window; T: timing.

background files. In particular, the scorpeon model<sup>23</sup> with default settings was adopted for the creation of the background file.

## 2.4. XMM-Newton

The EPIC-pn on board XMM-Newton was configured in the SW mode, while both EPIC-MOS cameras operated in the timing mode. Data were processed using the Science Analysis System. We applied intensity filters to the time series of the entire field of view in the EPIC-pn data to remove several background flaring episodes (which were also visible at low energies), resulting in a net on-source exposure time of ~15 ks. No X-ray emission was detected at the source position. Using the eupper tool and assuming a circular region with a radius of 10" for the source photons along with an annulus with inner and outer radii of 10" and 20", respectively, for the background, we estimated a $3\sigma$ upper limit on the net count rate of 0.007 counts s$^{-1}$ over the 0.3–10 keV energy range. Given the nondetection in the EPIC-pn data, we did not perform further analysis on the EPIC-MOS and Reflection Grating Spectrometer data.

## 2.5. Chandra

The field of J0052 was observed with the ACIS-I instrument on board the Chandra X-ray Observatory three times on 2002

July 20 (7.7 ks), 2006 April 25 (49.3 ks), and 2006 April 26 (47.4 ks). We processed the data using the CIAO software package. In 2002, J0052 was detected at a count rate of $(1.9^{+1.0}_{-0.8}) \times 10^{-3}$ counts s$^{-1}$ (0.3–7 keV) at the 90% c.l., corresponding to a detection significance exceeding 99%. For the stacked observations from 2006, the count rate was $(1.3^{+0.6}_{-0.7}) \times 10^{-4}$ counts s$^{-1}$ (0.3–7 keV), representing a ~$3\sigma$ detection.

# 3. Data Analysis and Results

## 3.1. The EP/WXT Light Curve

The WXT light curve is shown in Figure 1. After removing spurious drops in count rates caused by instrumental artifacts, the analysis of the observations before and after the source detection revealed a rapid increase in emission starting around MJD 60456.8, with a linear rise at ~0.3 counts s$^{-1}$ day$^{-1}$, peaking around MJD 60457.5 after approximately 0.6 days. This was followed by a slow linear decay at ~0.02 counts s$^{-1}$ day$^{-1}$ until around MJD 60458.8 and then a faster decay at ~0.23 counts s$^{-1}$ day$^{-1}$.

## 3.2. X-Ray Spectral Analysis

The X-ray flux and spectral shape of J0052 varied significantly over hours. Therefore, we extracted time-resolved spectra over intervals with approximately constant count rates. For NICER observations 7204550101, 7204550102, 7204550103, and

<footnote><sup>a</sup> The observations considered for the spectral analysis in this Letter.</footnote>

<footnote n="23"><sup>23</sup> https://heasarc.gsfc.nasa.gov/docs/nicer/analysis_threads/scorpeon-overview/</footnote>
<figure><img src="image_4.png" /><figcaption>**Figure 1.** EP/WXT light curves of J0052 around the epoch of the outburst detection in the soft (0.5–2 keV; top) and hard (2–4 keV; middle) energy bands. The drop in the middle of the light curve of CMOS 37 is caused by instrumental effects. The time evolution of the hardness ratio is also plotted (bottom).</figcaption></figure>

7204550104, we divided them into four, six, four, and two epochs, respectively, extracting one spectrum for each. The signal was largely or completely background dominated in all NICER, EP/FXT, and Swift/XRT observations performed after June 2, so we excluded these spectra from further analysis. The final spectral sample consists therefore of 1 FXT, 1 XRT, and 19 NICER time-resolved spectra, hereafter referred to as epochs (see Table 1 for a summary). We grouped all the spectra ensuring a minimum of 25 counts per bin, allowing the use of the $\chi^2$ statistics. Figure 2(a) shows the evolution of the extracted spectra of the source. We included data only where the source emission was above the background, which varied from spectrum to spectrum but was always within the 0.3–1.3 keV range. We note that despite the not-optimal calibration of NICER data below ~0.4 keV,<fnref n="24" /> we decided to retain data between 0.3 and 0.4 keV due to the source bright soft emission.

We performed the spectral analysis with `xspec` v.12.14.b (K. A. Arnaud 1996). A systematic error of 1% was included, as recommended by the NICER team.<fnref n="25" /> The interstellar absorption was modeled through the `TBabs` component, with

abundances from J. Wilms et al. (2000) and cross sections from D. A. Verner et al. (1996). We derived the unabsorbed bolometric fluxes between 0.001 and 10 keV using the `cflux` convolution model. We initially adopted a simple absorbed blackbody model (`BBodyrad`) to fit the spectra. Although it matched the continuum reasonably well, the overall fit was poor, with significant unmodeled residuals (see, e.g., Figure 2, panel (c)). For NICER, artifacts like fluorescence lines from Earth's atmosphere can result from improper background modeling. To check for these residuals, we reextracted the spectrum from the brightest NICER snapshot using the `nicer13-spec` pipeline with the `bkgformat` option set to "script" instead of "file," allowing a proper background fit with the continuum. The residuals in the best-fit model were identical regardless of the method used, indicating that the observed spectral features are real and not due to the background. These residuals include two edges, at 0.32 keV in NICER data and at 0.47 keV in FXT data, likely due to calibration issues<fnref n="26" /> and are not discussed. Instead, we detected intrinsic absorption edges from the source at energies of ~0.63, ~0.88, and ~1.0 keV, although they almost never appeared all together in a single spectrum. Two localized excesses at about 0.5–0.7 and 0.9 keV were visible in Epochs 1–14 and 1, 3–6, respectively.

We tested four models to describe the spectra: (a) a `BBodyrad` component plus a set of absorption edges and Gaussians; (b) a `BBodyrad` component plus a set of absorption edges and an `APEC` model, which describes the emission spectrum from a collisionally ionized diffuse gas; (c) a NLTE stellar atmosphere model, obtained employing the grid 003 of the `TMAP`<fnref n="27" /> public repository (T. Rauch 2003) plus `APEC`; (d) `SSS_atm_SMC` (V. F. Suleimanov et al. 2024)<fnref n="28" /> plus `APEC`. In all fits including `APEC`, the abundance parameter was unconstrained, so we consistently kept it fixed to 0.1, i.e., the SMC abundances. Models (a) and (b) describe the spectra equally well. Acceptable fits are obtained with model (c) only for spectra corresponding to the decay phase of the outburst, i.e., from Epoch 11 onward. At the peak of the outburst, the model falls short in describing the emission above 0.8 keV. The maximum temperature allowed by the `TMAP` model might be indeed too low to account for the relatively hard spectra in Epochs 1–10. For similar reasons, model (d) is also inadequate to describe our data, especially at the peak. In the following, we will mainly consider the results obtained with model (b), being more physically motivated than model (a) and the only model able to describe coherently our entire data set.

We consistently obtained $N_H$ values within the range of $(1-2) \times 10^{21} \text{ cm}^{-2}$. This range is significantly larger than the expected foreground value for the SMC, which is only ~6 × $10^{20} \text{ cm}^{-2}$ (J. M. Dickey & F. J. Lockman 1990). The discrepancy between our obtained and the expected $N_H$ values might indicate the presence of local absorption along the line of sight, as has been suggested for other `SSSs` in the SMC (M. J. Coe et al. 2020; J. A. Kennea et al. 2021). From the blackbody normalization, we estimated the radius of the blackbody emitting regions in each time-resolved spectrum, finding a range of about 3–6 × $10^3 \text{ km}$. The best-fit values

<footnote n="24">²⁴ https://heasarc.gsfc.nasa.gov/docs/nicer/analysis_threads/cal-recommend/</footnote>

<footnote n="25">²⁵ https://heasarc.gsfc.nasa.gov/docs/nicer/analysis_threads/spectrum-systematic-error/</footnote>

<footnote n="26">²⁶ See https://heasarc.gsfc.nasa.gov/docs/nicer/analysis_threads/arf-rmf/.</footnote>

<footnote n="27">²⁷ Tübingen NLTE model atmosphere package: http://astro.uni-tuebingen.de/~rauch/TMAF/flux_HHeCNONeMgSiS_gen.html.</footnote>

<footnote n="28">²⁸ `SSS_atm_SMC` models the X-ray supersoft emission from a hot WD atmosphere with SMC abundances.</footnote>
<figure><img src="image_5.png" /><figcaption>**Figure 2.** Summary plot for the spectral analysis performed for J0052. Panel (a): single spectra divided by each instrument's effective area taken by EP/FXT (triangles), Swift/XRT (squares), and NICER (circles). The time for each spectrum is referenced from MJD 60458.23. Panel (b): best-fit unabsorbed models of selected observations. Panel (c): unfolded NICER spectrum taken during Epoch 4 along with the best-fit model (top panel). The *bbodyrad* component is represented by a dashed black line, while the *SPEC* component is drawn with dotted silver line. The energies of the identified edges are marked with vertical maroon lines. The unmodeled and modeled residuals are shown in the middle and bottom panel.</figcaption></figure>

obtained for both the continuum and the edges are listed in Tables A1 and A2, respectively. The simultaneous evolution of the best-fit parameters can be observed in Figure 3.

### 3.3. Preoutburst Detections and Postoutburst Upper Limits

The upper limit on the net count rate measured in the last four stacked Swift/XRT observations and in the XMM-Newton/EPIC-pn observation converts to unabsorbed bolometric fluxes of $<6 \times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$, respectively, assuming the same parameters measured in the latest NICER spectrum (see Table A1). On the other hand, the count rate recorded in the 2002 Chandra observation, which has a higher signal-to-noise ratio compared to the 2006 observations, corresponds to an observed flux of $\approx 3 \times 10^{-14}$ erg cm$^{-2}$ s$^{-1}$ (0.3–8 keV; computed using the *srcflux* tool, assuming an absorbed power-law model with an absorption column density of $2 \times 10^{21}$ cm$^{-2}$ and a photon

index in the range of 0.5–1.5). This results in a luminosity of $\approx 1.5 \times 10^{34}$ erg s$^{-1}$ (0.3–8 keV). For the 2006 observations, assuming the same spectral shape as above and scaling based on the count rate, we estimate a luminosity that is about 1 order of magnitude lower, $\approx 10^{33}$ erg s$^{-1}$.

### 3.4. Searches for Periodic Signals

To search for periodic signals, we used the data collected when the source was brightest and corrected the photon times of arrival to the barycenter of the Solar System. We considered data of the first two NICER observations and the first EP/FXT observation, all restricted to the 0.3–2 keV energy band to maximize the signal-to-noise ratio. Swift/UVOT data taken in event mode during the longest observation on June 17 were also searched. We then extracted power spectra from the unbinned event files from each data set separately. No statistically significant signals were detected in any of the
<figure><img src="image_6.png" /><figcaption>**Figure 3.** Best-fit spectral parameters obtained in this work and their evolution during the outburst. The color sequence is the same as the one used in Figure 2. See also Tables A1 and A2.</figcaption></figure>

data. The 3σ upper limits on the pulsed fraction were calculated (G. L. Israel & L. Stella 1996), resulting in <4% up to 2 kHz in the NICER data, between 10% and 16% up to 250 Hz in the EP/FXT data, and between 7% and 10% up to 50 Hz in the Swift/UVOT data. For the NICER data, which had the highest signal-to-noise ratio, we also employed Fourier-domain acceleration search techniques using PRESTO<sup>29</sup> (S. M. Ransom et al. 2002) to search for signals up to 2 kHz. This analysis was conducted on the entire data set of the first two observations as well as on 300 s data chunks. No statistically significant signal was detected. These results are consistent with those reported by G. K. Jaisawal et al. (2024).

## 4. Discussion and Conclusions

In this Letter, we have presented the EP/WXT discovery of an X-ray outburst from the BeWD binary EP J005245.1−722843 (V. Antoniou et al. 2009) in the SMC and the follow-up X-ray observing campaign performed with EP/FXT, NICER, Swift/XRT, and XMM-Newton/EPIC. The source experienced a rapid increase in bolometric luminosity from ~4 × 10<sup>33</sup> erg s<sup>−1</sup> before the onset of the outburst to ~4 × 10<sup>38</sup> erg s<sup>−1</sup> at the outburst peak. A slower decay over 1 order of magnitude in luminosity was then observed in about 6 days, after which the system was no longer detectable by NICER, likely due to being too faint, too soft, or both. An XMM-Newton/EPIC-pn observation set a stringent upper limit of <5 × 10<sup>35</sup> erg s<sup>−1</sup> on the bolometric luminosity of the source 12 days after the last NICER detection. This indicates a reduction in luminosity by at least 3 orders of magnitude during that period, marking the end of the outburst. We therefore conclude that the total duration of the outburst from J0052 ranged between 6 and 12 days, making it remarkably short. Throughout the entire outburst, the source consistently exhibited a soft spectral shape, with no emission detected beyond ~1.5 keV. However, the observed spectrum was not as soft as a typical SSS, whose emission usually cuts off at ~0.7 keV (e.g., M. Orio et al. 2022), making this source somewhat peculiar. The spectra are well modeled using a combination of a blackbody spectrum and an apec model. The source reached a blackbody temperature of ~120 eV in about 1 day since the EP/WXT detection and then exhibited a rapid decay to ~50 eV in less than a week. The blackbody radius showed an almost opposite trend, going from a minimum of ~2700 km coincident with the kT peak to a maximum of ~6700 km after about 2 days and then plateauing for the remainder of the outburst. Similar trends have been reported in the SSS phase of several novae, analyzed as well with blackbody models (see, e.g., K. L. Page et al. 2015, 2020), although the peak temperatures were significantly lower than those seen in J0052. A more remarkable similarity in both the peak temperatures and the spectral shape, i.e., not soft enough for a canonical SSS, can be spotted in the outbursting BeWD binary MAXI J0158−744 (K. L. Li et al. 2012). The initial temporary decrease in the temperature observed between Epochs 1 and 3 could also be due to variability in the early SSS phase, as reported by J. P. Osborne et al. (2011), or could be an artifact due to intercalibration issues between FXT, XRT, and NICER. We caution that blackbody models are only a crude approximation of the actual physics in SSSs. Therefore, while the trends observed in the parameters are reliable, the specific values obtained for these parameters are less so. The attempt to model the thermal component with more physically reliable WD atmosphere models (T. Rauch 2003; V. F. Suleimanov et al. 2024) was instead not as successful. At the peak of the outburst, these models can not describe our data (as in MAXI J0158−744; K. L. Li et al. 2012), possibly because of the upper limit on the WD temperature being too stringent and/or some of the model assumptions being not applicable to this source. Good fits are instead obtained in the decline phase of the outburst, a stage where possibly the emission of our source resembles more that of a hot compact WD.

Almost all the spectra analyzed in this Letter showed a dense and intricate array of discrete spectral features in absorption and emission. The ~0.63, ~0.88 keV, and ~1.0 keV edges can be associated<sup>30</sup> with absorption edges from hydrogen-like species

<footnote n="29"><sup>29</sup> https://github.com/scottransom/presto</footnote>

<footnote n="30"><sup>30</sup> We mostly refer to https://xdb.lbl.gov/Section1/Sec_1-8.html for the identification of these lines.</footnote>
<figure><img src="image_7.png" /><figcaption>**Figure 4.** Schematic evolutionary path of J0052 as described in the text. This is the lower-mass analog of the evolutionary model by G. M. H. J. Habets (1986a) for forming a Be X-ray binary, here adapted for the case in which the accreting object is a massive Ne-O WD.</figcaption></figure>

of N, O, and Ne, respectively. Nitrogen, oxygen, and neon features are often found in the X-ray spectra of novae and SSSs (see, e.g., J. U. Ness et al. 2013; J. J. Drake et al. 2021; Y. Bhargava et al. 2024), and they can be considered as a signature of a WD accretor. The additional emission lines are instead well modeled by including in the fits a relatively cold, i.e., $kT_{\text{SPEC}} \sim 100-200$ eV, shocked gas emitting an APEC component. The presence of this component might suggest the presence of shocked ejecta around an emerging hot WD, as sometimes observed in novae with early turned-on SSS phases (see, e.g., K. L. Page et al. 2020). We conclude that the majority of the observed properties of the X-ray outburst of J0052, with the exception of the rather hard spectral shape at the outburst peak and its incompatibility with WD atmosphere models, are in line with the SSS phase of a nova eruption. In this scenario, the observed X-ray outburst would be explained by thermonuclear burning on the surface of the WD and not by the initial thermonuclear runaway phase of a small-amplitude nova, as previously suggested by T. M. Gaudin et al. (2024) based only on the first Swift observation of the source. However, the duration of the observed supersoft outburst, i.e., about a week, stands out with respect to most novae in the SSS phase, which typically lasts tens to hundreds of days (e.g., M. Henze et al. 2010, but see M. Orio et al. 2023 for exceptions). The duration of an SSS phase depends primarily on the mass and chemical composition of the WD. WDs near

the Chandrasekhar limit have shorter SSS phases since less accreted material is needed for a thermonuclear runaway. Conversely, higher hydrogen abundance in the accreted material slows the process (I. Hachisu et al. 2006). The combination of a ~10 day duration, a peak effective temperature of ~120 eV, and a peak luminosity above $10^{38}$ erg s$^{-1}$ suggests the presence of a very massive object (W. M. Wolf et al. 2013; M. D. Soraisam et al. 2016). However, we caution that the existing literature may not be specifically tailored to the scenario of a WD accreting from a Be star. In addition, the presence of O and Ne features could be the signature of an Ne-O WD in the system, which is among the most massive WDs, typically exceeding $1.1 M_{\odot}$ (K. Taguchi et al. 2023). Future high-resolution optical and/or X-ray spectra might shed light on the chemical composition of the WD and potentially confirm the accreting star in J0052 as an Ne-O WD.

It is noteworthy that in novae, the SSS phase typically begins after an optical-to-UV brightening phase caused by the expanding WD photosphere, which usually lasts for weeks (e.g., M. J. Darnley et al. 2016). In the case of J0052, however, only a 0.5 magnitude increase in the optical band was observed 6 hr before the X-ray outburst (T. M. Gaudin et al. 2024), which would be unusual for a thermonuclear runaway-driven explosion. Similarly, a short-lived optical flare occurred just 11 hr before the X-ray outburst of the BeWD MAXI J0158-744 (K. L. Li et al. 2012). It is possible that the optical brightening in both that source and in
J0052 might have occurred but went undetected due to the dominant emission from the optically bright Be companion star (as suggested by K. L. Li et al. 2012). Alternatively, an intrinsically short optical flare could be due the whole nova explosion being quick, possibly because of exceptionally low ejecta mass (M. Morii et al. 2013), as also suggested by T. M. Gaudin et al. (2024).

The ensemble of the X-ray spectral clues gathered in this work indicates that the compact object in J0052 is a massive, possibly Ne-O, WD. We therefore propose that J0052 is a BeWD binary. Such probably common but very elusive systems provide valuable insights into the evolution of binary stars, showcasing observable effects of the common envelope (CE) phase and mass transfer in massive star binaries. A binary system initially composed of two main-sequence stars of masses of about 8 and 6 $M_\odot$ with an orbital period of ~3 days is a possible progenitor for J0052. Following conservative mass transfer as computed by Y. Götzberg et al. (2018), the system evolves into a $2.1 M_\odot$ He star and a $11.9 M_\odot$ Be main-sequence star with an orbital period of about 21 days. According to the computations by S. E. Woosley (2019), such a He star terminates its evolution as an Ne-O WD with a mass of $1.23 M_\odot$, confirming results by G. M. H. J. Habets (1986b). During the evolution of the He star, mass transfer to the Be star resumes, but the mass transfer is nonconservative due to the He star's convective envelope, which expands on a dynamical timescale, leading to the formation of a CE surrounding the Be star and the degenerate Ne-O core. When the envelope has expanded to the second Lagrangian point, it leads to the formation of a circumbinary disk with a mass of about $0.9 M_\odot$. This causes the orbital period to decrease by a factor of about 0.8, resulting in a final orbital period of about 17 days (see Section 4.3.8 of T. M. Tauris & E. P. J. van den Heuvel 2023 for orbital period effects of a circumbinary disk). The disk will finally be expelled from the system. We sketch the proposed evolutionary scenario for J0052 in Figure 4. It is noteworthy that, despite the above scenario being only one of the possible models for the secular evolution of the binary system, the expected age of the system fits very well with the J0052 age, i.e., of about 40 Myr (V. Antoniou et al. 2019). Indeed, the initial $8 M_\odot$ star takes approximately ~35 Myr before evolving into a system with a Be star of about $12-13 M_\odot$, characterized by lifetimes of typically ~20 Myr (G. Schaller et al. 1992).

If the nature of the source is confirmed, this work may represent one of the very few cases where an X-ray outburst in such systems has been thoroughly monitored for most of its duration, from rise to decay. The unprecedented monitoring capabilities of EP in the soft X-ray range are particularly well suited to discovering supersoft outbursts from these sources, offering the opportunity to address many unresolved questions about them.

## Acknowledgments

We thank the anonymous referee for their insightful and useful comments. This work is based on data obtained with

Einstein Probe, a space mission supported by Strategic Priority Program on Space Science of Chinese Academy of Sciences, in collaboration with ESA, MPE, and CNES (grant No. XDA15310000), the Strategic Priority Research Program of the Chinese Academy of Sciences (grant No. XDB0550200), and the National Key R&D Program of China (2022YFF0711500). This work is also based on data obtained by the Neil Gehrels Swift Observatory (a NASA/UK/ASI mission) supplied by the UK Swift Science Data Centre at the University of Leicester; NICER, a 0.2–12 keV X-ray telescope operating on the International Space Station, funded by NASA; and XMM-Newton, an ESA science mission with instruments and contributions directly funded by ESA Member States and NASA. This Letter employs a list of Chandra data sets, obtained by the Chandra X-ray Observatory, contained in the Chandra Data Collection (CDC) with DOI: 10.25574/cdc.294. We thank N. Schartel for approving Target of Opportunity observations with XMM-Newton in the Director's Discretionary Time and the XMM-Newton Science Operation Center for scheduling and carrying out the observations. Similarly, we thank the Swift team for approving and performing our Target of Opportunity observations.

We acknowledge the support by the National Natural Science Foundation of China (grant Nos. 12321003, 12103065, 12203071, 12333004, 12373040, and 12021003), the China Manned Space Project (grant Nos. CMS-CSST-2021-A13 and CMS-CSST-2021-B11), and the Youth Innovation Promotion Association of the Chinese Academy of Sciences. We acknowledge the data resources and technical support provided by the China National Astronomical Data Center, the Astronomical Science Data Center of the Chinese Academy of Sciences, and the Chinese Virtual Observatory. A.M. and N.R. are supported by the European Research Council (ERC) under the European Union's Horizon 2020 research and innovation program (ERC Consolidator grant "MAGNESIA" No. 817661, PI: NR). F.C.Z. is supported by a Ramón y Cajal fellowship (grant agreement RYC2021-030888-I). A.M., F.C.Z., and N.R. acknowledge support from grant SGR2021-01269 from the Catalan Government (PI: Graber/Rea). HY acknowledges support from the China Scholarship Council (No. 202310740002). N.R. acknowledges support from the ESA Science Faculty Visitor program to ESTEC (funding reference ESA-SCI-E-LE-054), where most of this work has been carried out. S.G. acknowledges the support of the CNES. We thank A. Zesas, D. De Martino, S. Scaringi, T. Maccarone, and D. Buckley for the fruitful discussion.

**Facilities:** EP, Swift, NICER, XMM, CXO.

**Software:** CIAO v4.16.0 (A. Fruscione et al. 2006), fxtsoftware v1.05, HEASOFT v6.33.2 (Nasa High Energy Astrophysics Science Archive Research Center Heasarc, 2014), Matplotlib v3.9 (J. D. Hunter 2007), NICERDAS v12, SAS v21.0.0 (C. Gabriel et al. 2004), XSPEC v12.14.0h (K. A. Arnaud 1996).
## Appendix

In this Appendix, we present detailed results from the spectral analysis of all the observations considered in this work. In

Table A1 we report the results of the spectral analysis for the spectral continuum, while in Table A2 we provide the best-fit parameters for the absorption edges.

<table><caption>**Table A1**
Spectral Analysis Results Using a Blackbody Model for the Continuum Emission and *vapec*</caption><thead><tr><th>Time<br>(days)</th><th>nH<br>(×10<sup>22</sup>)</th><th>kT<sub>bb</sub><br>(eV)</th><th>R<sub>bb</sub><br>(×10<sup>3</sup> km)</th><th>kT<sub>apec</sub><br>(eV)</th><th>N<sub>apec</sub><br>...</th><th>L<sub>bol</sub><br>(×10<sup>38</sup> erg s<sup>-1</sup>)</th><th>χ<sup>2</sup><br>(d.o.f.)</th></tr></thead><tbody><tr><td>0.01</td><td>0.140<sup>+0.007</sup><sub>-0.006</sub></td><td>100<sup>+1</sup><sub>-2</sub></td><td>4.9 ± 0.1</td><td>240 ± 4</td><td>0.28 ± 0.04</td><td>3.1 ± 0.2</td><td>411(287)</td></tr><tr><td>0.43</td><td>0.180<sup>+0.016</sup><sub>-0.017</sub></td><td>110<sup>+4</sup><sub>-3</sub></td><td>4.0<sup>+0.7</sup><sub>-0.5</sub></td><td>230<sup>+16</sup><sub>-18</sub></td><td>(0.30)</td><td>3.90 ± 0.07</td><td>87(73)</td></tr><tr><td>0.55</td><td>0.140 ± 0.003</td><td>120 ± 2</td><td>2.8<sup>+0.2</sup><sub>-0.1</sub></td><td>210 ± 6</td><td>0.29 ± 3</td><td>2.80 ± 0.02</td><td>98(90)</td></tr><tr><td>0.61</td><td>0.140<sup>+0.004</sup><sub>-0.003</sub></td><td>120 ± 2</td><td>2.9<sup>+0.2</sup><sub>-0.1</sub></td><td>210 ± 6</td><td>0.27<sup>+0.03</sup><sub>-0.02</sub></td><td>2.70 ± 0.02</td><td>103(90)</td></tr><tr><td>0.67</td><td>0.140 ± 0.003</td><td>120 ± 2</td><td>2.9<sup>+0.2</sup><sub>-0.1</sub></td><td>210<sup>+6</sup><sub>-7</sub></td><td>0.25 ± 0.03</td><td>2.60 ± 0.02</td><td>107(90)</td></tr><tr><td>0.74</td><td>0.140 ± 0.004</td><td>110 ± 2</td><td>3.2 ± 0.2</td><td>190 ± 8</td><td>0.25 ± 0.03</td><td>2.90 ± 0.02</td><td>124(90)</td></tr><tr><td>1.25</td><td>0.170 ± 0.006</td><td>90 ± 2</td><td>5.5<sup>+0.5</sup><sub>-0.4</sub></td><td>170 ± 5</td><td>0.30 ± 0.02</td><td>3.80 ± 0.03</td><td>100(79)</td></tr><tr><td>1.45</td><td>0.150<sup>+0.006</sup><sub>-0.005</sub></td><td>94 ± 2</td><td>4.5<sup>+0.5</sup><sub>-0.4</sub></td><td>170 ± 8</td><td>0.25 ± 0.03</td><td>2.90 ± 0.02</td><td>87(69)</td></tr><tr><td>1.51</td><td>0.150 ± 0.006</td><td>92 ± 2</td><td>4.7<sup>+0.5</sup><sub>-0.4</sub></td><td>150 ± 5</td><td>0.30<sup>+0.06</sup><sub>-0.05</sub></td><td>3.30 ± 0.03</td><td>84(65)</td></tr><tr><td>1.58</td><td>0.160<sup>+0.008</sup><sub>-0.007</sub></td><td>89 ± 2</td><td>5.3<sup>+0.7</sup><sub>-0.5</sub></td><td>160<sup>+9</sup><sub>-10</sub></td><td>0.22 ± 0.04</td><td>3.00 ± 0.02</td><td>85(64)</td></tr><tr><td>1.64</td><td>0.160 ± 0.008</td><td>87<sup>+3</sup><sub>-2</sub></td><td>5.5<sup>+0.7</sup><sub>-0.6</sub></td><td>160<sup>+11</sup><sub>-10</sub></td><td>0.18<sup>+0.04</sup><sub>-0.03</sub></td><td>2.80 ± 0.02</td><td>79(63)</td></tr><tr><td>1.71</td><td>0.160<sup>+0.008</sup><sub>-0.007</sub></td><td>86 ± 2</td><td>5.4<sup>+0.7</sup><sub>-0.5</sub></td><td>140<sup>+10</sup><sub>-13</sub></td><td>0.24<sup>+0.10</sup><sub>-0.06</sub></td><td>3.20 ± 0.03</td><td>71(63)</td></tr><tr><td>2.03</td><td>0.140<sup>+0.007</sup><sub>-0.006</sub></td><td>80 ± 3</td><td>5.0<sup>+0.9</sup><sub>-0.6</sub></td><td>(140)</td><td>0.15 ± 0.03</td><td>2.00 ± 0.02</td><td>58(37)</td></tr><tr><td>2.16</td><td>0.140 ± 0.008</td><td>81 ± 3</td><td>4.8<sup>+0.8</sup><sub>-0.6</sub></td><td>(140)</td><td>0.06 ± 0.02</td><td>1.50 ± 0.01</td><td>39(42)</td></tr><tr><td>2.29</td><td>0.140<sup>+0.011</sup><sub>-0.010</sub></td><td>78 ± 4</td><td>4.9<sup>+1.3</sup><sub>-0.8</sub></td><td>(140)</td><td>0.06<sup>+0.03</sup><sub>-0.02</sub></td><td>1.34 ± 0.02</td><td>40(41)</td></tr><tr><td>2.42</td><td>0.140<sup>+0.011</sup><sub>-0.010</sub></td><td>77 ± 4</td><td>4.8<sup>+1.2</sup><sub>-0.7</sub></td><td>(140)</td><td>0.04 ± 0.02</td><td>1.20 ± 0.01</td><td>30(42)</td></tr><tr><td>3.06</td><td>0.130 ± 0.010</td><td>74<sup>+3</sup><sub>-2</sub></td><td>3.7<sup>+0.6</sup><sub>-0.5</sub></td><td>...</td><td>...</td><td>0.46 ± 0.01</td><td>32(40)</td></tr><tr><td>3.13</td><td>0.110 ± 0.010</td><td>76 ± 2</td><td>3.1 ± 0.3</td><td>...</td><td>...</td><td>0.35 ± 0.01</td><td>35(42)</td></tr><tr><td>4.29</td><td>(0.130)</td><td>56 ± 2</td><td>5.6<sup>+0.6</sup><sub>-0.5</sub></td><td>...</td><td>...</td><td>0.38 ± 0.02</td><td>24(25)</td></tr><tr><td>5.06</td><td>(0.130)</td><td>53 ± 1</td><td>5.2<sup>+0.5</sup><sub>-0.4</sub></td><td>...</td><td>...</td><td>0.25 ± 0.01</td><td>35(26)</td></tr><tr><td>5.96</td><td>(0.130)</td><td>50 ± 2</td><td>5.1<sup>+1.0</sup><sub>-0.7</sub></td><td>...</td><td>...</td><td>0.19 ± 0.02</td><td>16(21)</td></tr></tbody></table>

Note. The parameters that have been fixed are indicated in parentheses. We consider MJD 60458.23 as reference time. Abundances in *vapec* fixed to 0.1.

<table><caption>**Table A2**
Spectral Analysis of the Absorption Edges Included in the Fits</caption><caption>Absorption Edges</caption><thead><tr><th>Time<br>(days)</th><th>E<sub>e,1</sub><br>(keV)</th><th>τ<sub>e,1</sub><br>...</th><th>E<sub>e,2</sub><br>(keV)</th><th>τ<sub>e,2</sub><br>...</th><th>E<sub>e,3</sub><br>(keV)</th><th>τ<sub>e,3</sub><br>...</th><th>E<sub>e,4</sub><br>(keV)</th><th>τ<sub>e,4</sub><br>...</th></tr></thead><tbody><tr><td>0.01</td><td>0.370 ± 0.006</td><td>(1.0)</td><td>...</td><td>...</td><td>0.910 ± 0.004</td><td>1.0 ± 0.6</td><td>...</td><td>...</td></tr><tr><td>0.43</td><td>...</td><td>...</td><td>...</td><td>...</td><td>0.870 ± 0.001</td><td>3.0<sup>+2.0</sup><sub>-0.8</sub></td><td>...</td><td>...</td></tr><tr><td>0.55</td><td>0.330 ± 0.005</td><td>(1.0)</td><td>...</td><td>...</td><td>0.870 ± 0.003</td><td>1.0 ± 0.7</td><td>0.990<sup>+0.006</sup><sub>-0.005</sub></td><td>1.20 ± 0.08</td></tr><tr><td>0.61</td><td>0.330 ± 0.005</td><td>(1.0)</td><td>...</td><td>...</td><td>0.870 ± 0.003</td><td>1.2 ± 0.1</td><td>0.990 ± 0.006</td><td>1.20 ± 0.08</td></tr><tr><td>0.67</td><td>0.330 ± 0.005</td><td>(1.0)</td><td>...</td><td>...</td><td>0.870 ± 0.004</td><td>1.2 ± 0.1</td><td>0.980 ± 0.007</td><td>1.20 ± 0.10</td></tr><tr><td>0.74</td><td>0.340<sup>+0.005</sup><sub>-0.006</sub></td><td>(1.0)</td><td>...</td><td>...</td><td>0.870 ± 0.005</td><td>1.1 ± 0.1</td><td>0.980 ± 0.009</td><td>1.20 ± 0.12</td></tr><tr><td>1.25</td><td>0.340<sup>+0.005</sup><sub>-0.006</sub></td><td>(1.0)</td><td>...</td><td>...</td><td>0.870 ± 0.002</td><td>2.8<sup>+0.3</sup><sub>-0.2</sub></td><td>...</td><td>...</td></tr><tr><td>1.45</td><td>0.330 ± 0.007</td><td>(1.0)</td><td>0.670 ± 0.008</td><td>0.39<sup>+0.08</sup><sub>-0.07</sub></td><td>0.850 ± 0.003</td><td>3.1<sup>+0.4</sup><sub>-0.3</sub></td><td>...</td><td>...</td></tr><tr><td>1.51</td><td>0.340 ± 0.005</td><td>(1.0)</td><td>0.680 ± 0.008</td><td>0.45 ± 0.01</td><td>0.850 ± 0.003</td><td>2.9 ± 0.3</td><td>...</td><td>...</td></tr><tr><td>1.58</td><td>0.340 ± 0.006</td><td>(1.0)</td><td>0.680 ± 0.006</td><td>0.62 ± 0.08</td><td>0.850<sup>+0.003</sup><sub>-0.004</sub></td><td>3.3<sup>+0.5</sup><sub>-0.4</sub></td><td>...</td><td>...</td></tr><tr><td>1.64</td><td>0.340 ± 0.006</td><td>(1.0)</td><td>0.670 ± 0.005</td><td>0.84<sup>+0.10</sup><sub>-0.09</sub></td><td>0.840<sup>+0.004</sup><sub>-0.005</sub></td><td>3.4<sup>+0.8</sup><sub>-0.5</sub></td><td>...</td><td>...</td></tr><tr><td>1.71</td><td>0.340 ± 0.006</td><td>(1.0)</td><td>0.670 ± 0.004</td><td>1.10<sup>+0.08</sup><sub>-0.07</sub></td><td>0.840 ± 0.003</td><td>3.4<sup>+0.6</sup><sub>-0.4</sub></td><td>...</td><td>...</td></tr><tr><td>2.03</td><td>0.300<sup>+0.013</sup><sub>-0.030</sub></td><td>(1.0)</td><td>0.670 ± 0.003</td><td>1.80<sup>+0.12</sup><sub>-0.12</sub></td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>2.16</td><td>0.320<sup>+0.007</sup><sub>-0.009</sub></td><td>(1.0)</td><td>0.660 ± 0.002</td><td>2.50 ± 0.11</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>2.29</td><td>0.310<sup>+0.009</sup><sub>-0.012</sub></td><td>(1.0)</td><td>0.660 ± 0.003</td><td>2.60<sup>+0.16</sup><sub>-0.15</sub></td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>2.42</td><td>0.320<sup>+0.008</sup><sub>-0.009</sub></td><td>(1.0)</td><td>0.660 ± 0.002</td><td>2.90<sup>+0.18</sup><sub>-0.16</sub></td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>3.06</td><td>(0.300)</td><td>(1.0)</td><td>0.650 ± 0.004</td><td>2.60 ± 0.20</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>3.13</td><td>(0.300)</td><td>(1.0)</td><td>0.660 ± 0.003</td><td>2.90 ± 0.20</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>4.29</td><td>(0.300)</td><td>(1.0)</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr></tbody></table>

Note. The parameters that have been fixed are indicated in parentheses. We consider MJD 60458.23 as reference time.
## ORCID iDs

A. Marino https://orcid.org/0000-0001-5674-4664

H. N. Yang https://orcid.org/0000-0002-7680-2056

F. Coti Zelati https://orcid.org/0000-0001-7611-1581

N. Rea https://orcid.org/0000-0003-2177-6388

S. Guillot https://orcid.org/0000-0002-6449-106X

G. K. Jaisawal https://orcid.org/0000-0002-6789-2723

C. Maitra https://orcid.org/0000-0002-0766-7313

J.-U. Ness https://orcid.org/0000-0003-0440-7193

F. Haberl https://orcid.org/0000-0002-0107-5237

E. Kuulkers https://orcid.org/0000-0002-5790-7290

L. Tao https://orcid.org/0000-0002-2705-4338

C. Jin https://orcid.org/0000-0002-2006-1615

H. Sun https://orcid.org/0000-0002-9615-1481

E. P. J. van den Heuvel https://orcid.org/0000-0002-6417-2018

R. Soria https://orcid.org/0000-0002-4622-796X

B. Zhang https://orcid.org/0000-0002-9725-2524

S.-S. Weng https://orcid.org/0000-0001-7595-1458

L. Ji https://orcid.org/0000-0001-9599-7285

G. B. Zhang https://orcid.org/0000-0001-8630-5435

X. Pan https://orcid.org/0000-0002-6938-3594

Y. Chen https://orcid.org/0000-0001-9834-2196

S. Jia https://orcid.org/0000-0002-5203-8321

D. Y. Li https://orcid.org/0000-0002-4562-7179

K. Gendreau https://orcid.org/0000-0001-7115-2819

M. Ng https://orcid.org/0000-0002-0940-6563

T. Strohmayer https://orcid.org/0000-0001-7681-5845

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