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# Optical and Infrared Study of the Obscured B[e] Supergiant High-mass X-Ray Binary IGR J16318-4848

Francis Fortin, Sylvain Chaty, Andreas Sander

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Francis Fortin, Sylvain Chaty, Andreas Sander. Optical and Infrared Study of the Obscured B[e] Supergiant High-mass X-Ray Binary IGR J16318-4848. The Astrophysical Journal, 2020, 894 (2), pp.86. 10.3847/1538-4357/ab881c hal-02564549

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# Optical and Infrared Study of the Obscured B[e] Supergiant High-mass X-Ray Binary IGR J16318-4848*

Francis Fortin<sup>1,2</sup> ●, Sylvain Chaty<sup>1,2</sup>, and Andreas Sander<sup>3</sup>

<sup>1</sup>Université de Paris, CNRS, Astroparticule et Cosmologie, F-75013 Paris, France

<sup>2</sup>AIM, CEA, CNRS, Université de Paris, Université Paris-Saclay, F-91191 Gif-sur-Yvette, France

<sup>3</sup>Armagh Observatory and Planetarium, College Hill, Armagh BT61 9DG, Northern Ireland

Received 2020 March 20; accepted 2020 April 7; published 2020 May 8

## Abstract

The supergiant high-mass X-ray binary IGR J16318-4848 was the first source detected by the INTERNATIONAL Gamma-Ray Astrophysics Laboratory satellite in 2003 and distinguishes itself by its high intrinsic absorption and B[e] phenomenon. It is the perfect candidate to study both binary interaction and the environment of supergiant B[e] stars. This study targets the local properties of IGR J16318-4848. We aim to clarify the geometry of this system and distinguish different key emitting regions in the binary. We provide optical to near-infrared (NIR) spectra from the Very Large Telescope/X-Shooter and analyze both fine structures of the lines and the broadband spectral energy distribution by adding archival mid-infrared Spitzer and Herschel data. We also perform a stellar atmosphere and wind modeling of the optical to NIR spectrum using the Potsdam Wolf-Rayet code. We determine the contribution of the irradiated inner edge of the dusty circumbinary disk, derive the velocity of an equatorial stellar wind, and suggest the compact object orbits within the cavity between the star and the disk. We report on flat-topped lines originating from a spherically symmetric disk wind, along with the first detection of what is likely the polar component of the stellar wind. Stellar atmosphere and wind modeling shows that the central star may have a helium-enhanced atmosphere, likely because of its intense wind shedding part of its hydrogen envelope. Finally, we compare the properties of IGR J16318-4848 with a similar source, CI Cameleopardis.

*Unified Astronomy Thesaurus concepts: Near infrared astronomy (1093); Optical observation (1169); High mass x-ray binary stars (733); Stellar atmospheres (1584); Stellar winds (1636)*

## 1. Introduction

Since 2002, the INTERNATIONAL Gamma-Ray Astrophysics Laboratory (INTEGRAL) has been observing the sky looking for gamma-ray sources of various nature. On top of significantly increasing the number of known X-ray binaries, INTEGRAL was able to discover a new type of highly obscured supergiant high-mass X-ray binary (sgHMXB), as reviewed in Chaty (2013) and Walter et al. (2015). These peculiar binaries host either a neutron star (NS) or a black hole in orbit around an early-type supergiant star. Depending on the configuration of the binary, the compact object can be fed through the intense stellar wind of its giant companion or by Roche lobe overflow. The study of such extreme objects is crucial for understanding both the environment of supergiant stars and the products of binary interaction.

The first source detected by INTEGRAL, IGR J16318-4848 is one of the most absorbed sgHMXBs in the Galaxy known to this day. Discovered on 2003 January 29 (Courvoisier et al. 2003) with INTEGRAL/IBIS in the 15-40 keV band, the X-ray column density is so high ($N_H \simeq 2 \times 10^{24}$ cm$^{-2}$; Matt & Guainazzi 2003; Walter et al. 2003) that its flux drops drastically below 5 keV. It is known to be a Galactic persistent X-ray source with recurrent outbursts that last up to ~20 days.

Filliatre & Chaty (2004) use optical and near-infrared (NIR) spectra to derive an absorption of $A_V = 17.4$ mag, which is far greater than the line-of-sight value of 11.4, while still 100 times lower than in X-rays. This led the authors to suggest a concentration of X-ray absorbing material local to the compact

object and the presence of a shell around the whole binary absorbing optical/NIR wavelengths. The NIR spectrum in Filliatre & Chaty (2004) shows many prominent emission features in the same way as CI Cameleopardis (CI Cam), the first HMXB to be detected with an sgB[e] companion. Forbidden [Fe II] lines and P Cygni profiles, also present in the NIR spectrum, are evidence of a complex and rich environment, local to the binary.

Later, Kaplan et al. (2006) used photometry to show evidence of mid-infrared (MIR) excess in IGR J16318-4848. The authors found that an ~1000 K blackbody can be fit to the MIR excess in the spectral energy distribution (SED), and while they could not further characterize the exact nature and amount of the emitting material, they associated it with the presence of warm dust around the central star. Moon et al. (2007) provided Spitzer spectra from 5 to 40 µm that reveal a rich environment composed of an ionized stellar wind, a lower-density region giving birth to forbidden lines, a photodisassociated region, and two-component circumstellar dust ($T > 700$ and ~180 K). The actual geometry of this component was yet unknown, and if organized spherically around the central star, it would not contribute significantly to the absorption, meaning a much colder component ($T < 100$ K) located at the outermost regions of the binary could be responsible for the extreme extinction. Rahoui et al. (2008) performed photometry on IGR J16318-4848 with the Very Large Telescope (VLT)/VISIR and reach a similar conclusion, i.e., that warm circumstellar dust is responsible for the MIR excess.

However, Ibarra et al. (2007) suggested that the column density is inhomogeneous and that the circumstellar matter could very well be concentrated in the equatorial plane, seen almost edge-on, hence the very high $N_H$. The outflow might

<footnote>* Based on observations collected at the European Organisation for Astronomical Research in the Southern Hemisphere under ESO program 089.D-0056.</footnote>
<table><caption>**Table 1**
Log of Our Observations</caption><thead><tr><th rowspan="2">Target</th><th rowspan="2">Status</th><th rowspan="2">R.A.<br>J2000</th><th rowspan="2">Decl.<br>J2000</th><th rowspan="2">Date<br>(UTC)</th><th colspan="2">Airmass</th><th colspan="2">Exposure</th></tr><tr><th>Start</th><th>End</th><th>VIS</th><th>NIR</th></tr></thead><tbody><tr><td>IGR J16318-4848</td><td>SCIENCE</td><td>16:31:48.41</td><td>-48:49:03.54</td><td>2012 Jul 08T01:00:42</td><td>1.129</td><td>1.108</td><td>1200 s</td><td>200 s</td></tr><tr><td>HD 145412</td><td>CALIB</td><td>16:13:11.73</td><td>-49:53:05.71</td><td>2012 Jul 08T00:26:00</td><td>1.156</td><td>1.155</td><td>5 s</td><td>5 s</td></tr></tbody></table>

then be bimodal, with a fast polar wind and a slow, dense equatorial outflow.

Chaty & Rahoui (2012) used VLT/VISIR MIR along with New Technology Telescope (NTT)/Sofl and Spitzer spectra to fit the SED of IGR J16318–4848. They reported the presence of an irradiated rim around the star at $T_{rim} = 3500–5500$ K and a warm dust component at $T_{dust} = 767$ K in the outer regions of the binary using models of Herbig AeBe-forming stars, which have circumstellar material analogous to IGR J16318–4848.

Jain et al. (2009) suggested a possible 80 day period based on Swift-BAT and INTEGRAL data. Recently, Iyer & Paul (2017) provided the results of a long-term observation campaign on IGR J16318–4848 with Swift/BAT. They derived an orbital period of 80.09 ± 0.01 days. The folded light curves reveal two distinct peaks separated by low-intensity phases. Several flares are detected and preferentially happen in the same phase as the main peak, which may indicate that the compact object is crossing a denser and more inhomogeneous medium. There is also a correlation between the intensity of the flares and the time gap between them. This is reminiscent of disk-fed systems, where the material is regularly depleted and then restructured.

Concerning the nature of the compact object, several pointed observations with XMM-Newton (Walter et al. 2003; Iyer & Paul 2017), Suzaku (Barragán et al. 2009), and NuSTAR (Iyer & Paul 2017) show no detectable pulsation in IGR J16318–4848. The absence of radio jets reported by Filiatre & Chaty (2004) could indicate the presence of an NS whose poles do not cross our line of sight, which is what we assume in the rest of this paper.

## 2. Observations and Data Reduction

The observations of IGR J16318–4848 and standard star HD 145412 were performed in 2012 July at the European Southern Observatory (ESO) in Chile under program ID 089.D-0056 (see summary in Table 1). Spectra from 300 to 2480 nm were acquired on the 8 m VLT Unit 2 Cassegrain (UT2) on three different arms (UVB, VIS, and NIR) of the X-Shooter instrument. Because of the high intrinsic absorption of the source, the UVB spectrum only allows one to set an upper limit to the flux. According to the phase diagram provided in Figure 3 in Iyer & Paul (2017), our X-Shooter observations took place at phase 0.944 ± 0.005, which corresponds to a low-intensity phase in X-rays.

### 2.1. Optical (UVB, VIS) and NIR Data

Optical (UVB) echelle spectra were obtained through a 0"/5 × 11" slit giving a spectral resolution of $R = 9700$ (31 km s⁻¹) over a spectral range of 300–560 nm with a dispersion of 0.02 nm pixel⁻¹. Four exposures of 300 s were taken, for a total integration time of 1200 s.

Optical (VIS) echelle spectra were obtained through a 0"/7 × 11" slit giving a spectral resolution of $R = 11,400$ (26 km s⁻¹) over a spectral range of 533–1020 nm with a dispersion of 0.02 nm

pixel⁻¹. Four exposures of 300 s were taken, for a total integration time of 1200 s.

The NIR echelle spectra were obtained through a 0"/6 × 11" slit giving a spectral resolution of $R = 8100$ (37 km s⁻¹) over a spectral range of 994–2580 nm with a dispersion of 0.06 nm pixel⁻¹. Twenty exposures of 10 s were taken, for a total integration time of 200 s.

All of the acquisitions followed the standard ESO nodding pattern. The data reduction was performed with ESOReflex, using the dedicated X-Shooter pipeline. It consists of an automated echelle spectrum extraction along with standard bias, dark and sky subtraction, and airmass correction. Median stacking was used to add individual spectra in order to correct for cosmic rays.

The wavelength calibration was done during the reduction using calibration lamp frames and OH sky lines (NIR). The rms of the solution is 0.003 and 0.009 nm for VIS and NIR spectra, respectively.

The spectral response was obtained along with the flux calibration using the standard star spectrum (see Section 2.2). Telluric absorption features were corrected using Molecfit (Kausch et al. 2015; Smette et al. 2015), a software that fits atmospheric features using a radiation transfer code and various parameters from the local weather.

### 2.2. Detector Response and Flux Calibration

We use the spectrum of the standard star HD 145412 to extract a response curve that allows us to both correct for spectral response and derive a flux calibration. The flux $F_ν(ν)$ in the standard star spectrum can be written as

$$F_{\nu}(\nu) = R_{\nu} \times \left( F_{\text{feat}} + \left( \frac{R_{*}}{D_{*}} \right)^{2} B_{\nu}(\nu, T) \right) \times A_{\nu} \times A_{\nu}^{\text{atm}}, \quad (1)$$

where $R_ν$ is the response curve we need to isolate, $F_{feat}$ is the flux of the emission and/or absorption features of the standard star, $B_ν(ν, T)$ is its blackbody continuum, $R_*$ and $D_*$ are the radius of the star and its distance, and $A_ν$ and $A_ν^{atm}$ are, respectively, interstellar and atmospheric absorption.

First, $A_ν^{atm}$ is computed by Molecfit. In the optical, absorption mainly comes from H₂O vapor and O₂, while CO₂, H₂O, and CH₄ dominate in the NIR.

Second, $A_ν$ is computed using the formula in Cardelli et al. (1989) for the optical and NIR. For HD 145412, we use the following values found in Morales Durán et al. (2006): $A_V = 0.77$ and $R_V = 3.208$.

Third, the term $(R_*/D_*)^2 B_ν(ν, T_{eff})$ is obtained by fitting a blackbody emission to the SED of HD 145412, where

$$B_{\nu}(\nu, T_{\text{eff}}) = \frac{2h\nu^3}{c^2} \frac{1}{\exp\left(\frac{h\nu}{kT_{\text{eff}}} - 1\right)}. \quad (2)$$
<table><caption>**Table 2**
Photometric Data for HD 145412</caption><thead><tr><th>Wavelength (nm)</th><th>Flux (Jy)</th><th>Flux Err. (Jy)</th><th>Reference</th></tr></thead><tbody><tr><td>357.06</td><td>2.240</td><td>0.245</td><td>UVB; Mermilliod 1991<sup>a</sup></td></tr><tr><td>428.00</td><td>6.384</td><td>0.088</td><td>Tycho<sup>b</sup></td></tr><tr><td>437.81</td><td>6.261</td><td>0.685</td><td>UVB; Mermilliod 1991</td></tr><tr><td>534.00</td><td>6.889</td><td>0.063</td><td>Tycho</td></tr><tr><td>546.61</td><td>6.483</td><td>0.709</td><td>UVB; Mermilliod 1991</td></tr><tr><td>585.76</td><td>5.519</td><td>0.604</td><td>Gaia DR1<sup>c</sup></td></tr><tr><td>1 235.00</td><td>4.595</td><td>0.097</td><td>2MASS<sup>d</sup></td></tr><tr><td>1 662.00</td><td>3.022</td><td>0.086</td><td>2MASS</td></tr><tr><td>2 146.50</td><td>2.202</td><td>0.183</td><td>DENIS<sup>e</sup></td></tr><tr><td>2 159.00</td><td>2.079</td><td>0.046</td><td>2MASS</td></tr><tr><td>3 352.60</td><td>0.950</td><td>0.081</td><td>WISE<sup>f</sup></td></tr><tr><td>3 507.51</td><td>0.880</td><td>0.046</td><td>GLIMPSE<sup>g</sup></td></tr><tr><td>4 436.58</td><td>0.523</td><td>0.023</td><td>GLIMPSE</td></tr><tr><td>4 602.80</td><td>0.565</td><td>0.016</td><td>WISE</td></tr><tr><td>5 628.10</td><td>0.356</td><td>0.010</td><td>GLIMPSE</td></tr><tr><td>7 589.16</td><td>0.207</td><td>0.004</td><td>GLIMPSE</td></tr><tr><td>8 228.36</td><td>0.188</td><td>0.004</td><td>AKARI<sup>h</sup></td></tr><tr><td>11 560.80</td><td>0.088</td><td>0.002</td><td>WISE</td></tr><tr><td>22 088.30</td><td>0.030</td><td>0.003</td><td>WISE</td></tr></tbody></table>

**Notes.**

a Mermilliod (2006).

b Høg et al. (2000).

c Collaboration et al. (2016).

d Skrutskie et al. (2006).

e Epchtein et al. (1999).

f Wright et al. (2010).

g Churchwell et al. (2009).

h Ishihara et al. (2010).

The Gaia archive (Collaboration et al. 2018) provides a parallax of $5.48 \pm 0.05$ mas for HD 145412, which corresponds to a distance $D_* = 183 \pm 2$ pc. Using the calibrated photometric points available in the literature (see Table 2), we find a blackbody of temperature $T_{\text{eff}} = 8615 \pm 133$ K and radius $R_* = 6.1 \pm 0.1 R_\odot$. Considering the uncertainties of the parameters $T_{\text{eff}}$ and $R_*$, our flux calibration uncertainty varies from 8% to 4% on the full spectral range (533–2478 nm).

Finally, each individual spectral feature $F_{\text{feat}}$ of the standard star is fitted in relation to the local continuum and subtracted. These features mostly come from hydrogen series in absorption (Balmer, Paschen, and Brackett) and are Lorentzian-shaped.

Because of the various residuals of feature fitting (both stellar and atmospheric) and a poor signal-to-noise ratio toward the blue part of the spectra, we apply a median filter with a 21 pixel window width to the cleaned response curve. This allows us to obtain a smoother response at small scales while keeping its overall shape. The boundary effects of such a filter on the edge of the detectors can be neglected behind the signal-to-noise ratio drop at the edges and the coverage of both the VIS and NIR arms in the region 994–1020 nm. A normalized version of the final optical (VIS) to NIR spectrum is shown in Figure 1. As for UVB, the standard deviation of the data in the middle of the wavelength range (400–475 nm) provides an upper limit on the flux of $4.7 \times 10^{-6}$ Jy.

### 3. Spectral Feature Analysis

Here we describe the features in the X-Shooter spectrum (see Figure 1) and derive the associated parameters that we will

discuss in Section 4. The lines are identified using previous studies of IGR J16318–4848 (Filliatre & Chaty 2004) or studies of similar P Cygni, sgB[e], or even T Tauri sources (Edwards et al. 1987; Hillier et al. 1998; Hynes et al. 2002; Clark et al. 2013). We mostly detect hydrogen and helium in emission, along with iron and other metals. All of the velocity shifts mentioned are given in the heliocentric rest frame.

#### 3.1. Hydrogen Lines

The hydrogen lines are by far the most prominent features in the X-Shooter spectrum. The Paschen, Brackett, and Pfund series are visible up to quantum numbers of ~20–25, above which they become too faint and/or blended with one another. Balmer's $H\alpha$ is also present; however, its local background is under the detection limit. All values derived on this line will thus be lower limits. The list of hydrogen lines is available in Table A1.

The Paschen and Brackett lines show blueshifted absorption features, which are reminiscent of P Cygni profiles, although the Pfund lines do not show such features. In the rest of the paper, the P Cygni velocities are derived using the centers of each component (emission, absorption) and thus yield a lower limit on the expansion velocity of the medium. Fitting the hydrogen P Cygni profiles with a double Gaussian returns a mean velocity difference between emission and absorption of $264.8 \pm 4.3$ km s$^{-1}$. The FWHM of the emission component is $340.7 \pm 4.5$ km s$^{-1}$, on average, while it is $204.9 \pm 9.4$ km s$^{-1}$ for the absorption component. The center of the emission line is blueshifted at $-49 \pm 20$ km s$^{-1}$, on average. Several outliers may impact the estimate of the blueshift because of a poor signal-to-noise ratio or a line profile that deviates from a single Gaussian (discussed in Section 4.2); at this point, median statistics may give a more realistic velocity estimate of $-48 \pm 10$ km s$^{-1}$.

Balmer $H\alpha$ shows a thin emission line on top of a P Cygni profile. This is highlighted by subtracting the average P Cygni profile from the $H\alpha$ line (see Figure 2). The remaining line is then fitted with a single Gaussian. We estimate its width at $25.6 \pm 3.3$ km s$^{-1}$, an order of magnitude lower than the other hydrogen lines in the spectrum (see Table A1).

#### 3.2. Helium Lines

All helium lines are detected in emission and are atomic He I (see Table A2). They show blueshifted absorption features like hydrogen lines. Some He I lines may have extra components on top of the common P Cygni profile. For instance, He I $\lambda$1700 and $\lambda$2058 nm (see Figure 3) appear to have an extra wide emission component. The P Cygni emission and absorption components in these transitions are separated by $150 \pm 30$ km s$^{-1}$, on average. The mean FWHM of the emission is $290 \pm 20$ km s$^{-1}$, while it is $250 \pm 10$ km s$^{-1}$ for the absorption. The third, wide component can be reproduced by an emission line of width $840 \pm 20$ km s$^{-1}$. Its center is compatible with the center of the P Cygni emission line in both the $\lambda$1700 and $\lambda$2058 nm transitions.

The $\lambda$2112 nm line has a rather particular profile (Figure 3(b)). Given that the median ratio of equivalent width (EQW) between emission and absorption is 1.6 for helium lines, it is 1.2 for this particular line, making the P Cygni absorption stronger in comparison. We note that the emission line is centered at a similar velocity as other helium lines ($-119$ km s$^{-1}$). However, we measure a greater P Cygni velocity at $-352$ km s$^{-1}$.
<figure><img src="image_4.png" /><figcaption>**Figure 1.** Optical to NIR X-Shooter spectrum of IGR J16318-4848. Identified lines are indicated with gray ticks.</figcaption></figure>

<figure><img src="image_5.png" /><figcaption>**Figure 2.** Evidence for the double component in Hα. The line is plotted (a) before and (b) after subtracting the average H I P Cygni profile at a velocity of 264.8 km s⁻¹.</figcaption></figure>

<figure><img src="image_6.png" /><figcaption>**Figure 3.** (a) Helium profile with extra broad emission wings. (b) Helium profile with strong P Cygni absorption.</figcaption></figure>

## 3.3. Metallic Lines

### 3.3.1. Iron Flat-topped Lines

Lines associated with Fe II (and [Fe II], although much fainter) present a unique, flat-topped profile with narrow symmetrical wings. We show the average Fe II profile in

<figure><img src="image_7.png" /><figcaption>**Figure 4.** Average flat-topped profile of Fe II lines in IGR J16318-4848.</figcaption></figure>

Figure 4. To fit these lines, we use a model that consists of a rectangle function convolved with a Gaussian. All of the identified flat-topped lines from Fe II and [Fe II] are listed in Table A3.

For the allowed transitions of Fe II, the average half-width of the rectangle component is 250 ± 20 km s⁻¹, and 81 ± 22 km s⁻¹ is the half-width of the Gaussian broadening. The average heliocentric velocity is -75 ± 7 km s⁻¹.

For the forbidden transitions of [Fe II], the average half-width of the rectangle component is 285 ± 6 km s⁻¹, and 43 ± 6 km s⁻¹ is the half-width of the Gaussian broadening. The average heliocentric velocity is -47 ± 6 km s⁻¹.
<table><caption>**Table 3**
List of DIBs Found in the Spectrum of IGR J16318–4848 Compared to DIBs Found in Cox et al. (2014)</caption><thead><tr><th>λ<sub>DIB</sub> (nm)</th><th>λ<sub>fit</sub> (nm)</th><th>FWHM (nm)</th><th>EQW (nm)</th></tr></thead><tbody><tr><td>1069.7</td><td>1069.687</td><td>0.738</td><td>0.13</td></tr><tr><td>1078.0</td><td>1077.899</td><td>0.375</td><td>0.07</td></tr><tr><td>1179.7</td><td>1179.560</td><td>0.500</td><td>0.09</td></tr><tr><td>1317.5</td><td>1317.385</td><td>0.664</td><td>0.24</td></tr><tr><td>1780.3</td><td>1779.843</td><td>1.120</td><td>0.07</td></tr></tbody></table>

### 3.3.2. Other Metals

Forbidden [N II], [O I], and [S II] can be found close to Hα. They all have very small widths (<25 km s⁻¹), and their average blueshift is -32 ± 1 km s⁻¹. These two parameters are significantly different from all other lines visible in the spectrum.

Across the optical to NIR spectrum, Mg II doublets are found with intensities comparable to helium or even hydrogen lines. They are compiled in Table A4. Without taking into account the polluted lines with bad signal, their average FWHM is 339 ± 40 km s⁻¹, and the average blueshift is -100 ± 30 km s⁻¹.

### 3.4. Diffuse Interstellar Bands

A set of diffuse interstellar bands (DIBs) in X-Shooter spectra are presented in Cox et al. (2014). The authors showed a correlation between the EQWs of the DIBs with the absorption in the line of sight. We measured the EQW of the DIBs we detected in the spectrum of IGR J16318–4848 (see Table 3). Our EQW measurements of DIBs are all compatible with $A_V > 10.9$ according to the correlation coefficients in Cox et al. (2014; see comparison in Figure 5). If we compute the absorption with the relation they derived, we obtain $A_V = 38.6 \pm 2.2$ for λ1180 nm and $A_V = 23.6 \pm 5.6$ for λ1317 nm. The first is not a realistic estimation, as the correlation domain does not go above $A_V = 11$ ($E_{B-V} \sim 3.5$), and we suggest that it is no longer valid at higher absorption values. On the other hand, the value computed for the λ1180 nm transition is more realistic, although the correlation has a single point above $A_V = 11$, and as such, the confidence interval is rather large. We note that the value is still compatible with the one derived by Chaty & Rahoui (2012) of 18.3 ± 0.4. However, considering the high uncertainty of the correlation in this domain, we suggest that we cannot improve the estimation of the absorption, and that it is preferable to keep the value of $A_V = 18.3$.

## 4. Spectral Feature Interpretation

In this section, we use all of the information derived from the spectrum and conform it to the previously inferred geometry of IGR J16318–4848 to discuss the origin of the various line profiles.

### 4.1. Hydrogen-to-helium Abundance Ratio

Following the method presented in Allen et al. (1985) to derive the hydrogen-to-helium abundance ratio, we use the intensity ratio of He I λ1700 and λ1200 nm versus Brγ. This method provides a lower limit on the abundance, since it is probing H⁺ that recombines into H I (same for helium). With an absorption of

<figure><img src="image_8.png" /><figcaption>**Figure 5.** Correlation of the EQWs of DIBs with extinction (lines; adapted from Cox et al. 2014) vs. our measurements (dots). Red: λ1180 nm; blue: λ1317 nm.</figcaption></figure>

$A_V = 18.3$ mag, the first ratio, He I λ1700 nm/Brγ, gives 0.248, while the He I λ1200 nm/Brγ ratio provides 0.296. Both measurements are compatible with an evolved environment, as it is at least 3.7 times the solar abundance.

This could very well come from the properties of the birthplace of IGR J16318–4848. Especially, the disk could be at least partly formed from leftover material after the formation of the system. However, the first supernova event could also have fed the medium around the binary with heavy elements. Also, the central supergiant star likely emits a strong wind that could also participate in enhancing the environment with helium and metals, although the modeling of the stellar atmosphere and wind we present later in Section 7 may suggest that the helium mainly comes from the star itself.

### 4.2. Further Investigating the Profile of the Hydrogen Lines

Hydrogen lines with a sufficient signal-to-noise ratio are visibly different from a regular single Gaussian profile. This is particularly the case for Brγ and H I λ1280 nm, which have symmetric distortions around both sides of the emission line.

In their study of the supergiant A[e] binary HD 62623, Millour et al. (2011) presented spectro-interferometric data that are compatible with the Brγ line originating from the inner rim of the equatorial dusty disk of the system, which produces a double-peaked profile that probes the orbital velocity of the medium.

We thus fitted a double-peaked emission to the P Cygni profile of Brγ and H I λ1280 nm (Figure 6). It is possible to reproduce the data with a double peak, with the condition of adding an extra emission component of lower amplitude centered on the double-peaked profile.

For Brγ (Figure 6), the P Cygni velocity is measured to be 373 ± 11 km s⁻¹. The double peak is centered at -69 ± 5 km s⁻¹ and separated by 232 ± 10 km s⁻¹, which corresponds to an orbital velocity $v\sin(i) = 116 \pm 5$ km s⁻¹. The extra emission component is centered at -76 ± 5 km s⁻¹. Residuals on the red part of the line may correspond to the wing of the Brγ line due to electron scattering.

For H I λ1280 nm (Figure 6), the P Cygni velocity is measured to be 340 ± 12 km s⁻¹. The double peak is centered at -63 ± 7 km s⁻¹ and separated by 219 ± 13 km s⁻¹, which
<figure><img src="image_9.png" /><figcaption>**Figure 6.** Hydrogen line profiles fitted with a P Cygni profile on top of a double-peaked emission from a Keplerian rim. The continuum-subtracted data are in black, the model is in red, and the individual components are in dotted gray.</figcaption></figure>

corresponds to an orbital velocity $v\sin(i) = 110 \pm 7$ km s$^{-1}$.
The extra emission component is centered at $-71 \pm 5$ km s$^{-1}$.
The residual on the red part of the line is from an extra helium emission line also presenting a P Cygni profile.

While the two P Cygni velocities we derive are not fully compatible, the double-peaked structures are similar in both transitions and provide a mean orbital velocity of $v\sin(i) = 113 \pm 4$ km s$^{-1}$. The center of the double-peaked structure is located at $-66 \pm 4$ km s$^{-1}$. We note that for both lines, the center of the double-peaked profile is consistent with the center of the extra emission component.

The double-peaked profile and the P Cygni absorption are features expected in the geometry of IGR J16318–4848; however, we have yet to explain the origin of the extra emission component needed to reproduce the data. Because its center is compatible with the one of the double-peaked profile, we would first suggest that it comes from circumbinary material. But it could very well be a coincidence and in fact originate from the central star, as the maximum orbital velocity amplitude of the star is 6–30 km s$^{-1}$ (for a central star of 25–50 $M_{\odot}$ and a compact object of 1.4–10 $M_{\odot}$ in circular Keplerian orbit seen purely edge-on).

Concerning the asymmetry of the double-peaked component, we can quantify it by calculating the peak-to-peak amplitude asymmetry defined as $A = |\Delta I|/I$, where *I* is the amplitude of the peaks in which continuum was subtracted. This results in an amplitude asymmetry of 24% ± 7% for Brγ and 21% ± 9% for HⅠλ1280 nm.

This asymmetry can come from either a deviation from a circular orbit or the interaction between the compact object and the rim, bringing its heated material closer into the central cavity in the form of a wake. One way to pinpoint the origin of the asymmetry would be to observe the line profile change over time. If the rim is affected by the compact object, then the signal should be modulated by the period of the binary (~80 days; Iyer & Paul 2017). If ellipticity is responsible, then the modulation should come from the precession of the circumbinary material, which would be at a longer, superorbital period (see, e.g., Charles & Coe 2003). Although our current spectroscopic data do not allow us to choose one case over another, the results of the broadband SED discussed in

<figure><img src="image_10.png" /><figcaption>**Figure 7.** X-Shooter spectrum around the region containing narrow lines, indicated by red ticks.</figcaption></figure>

Section 6 allow cases in which the compact object orbits close enough to the rim so that its gravitational pull is of the same order of magnitude as the central star. This could set up the conditions for the rim to be modulated by the periodic passing of the compact object.

### 4.3. Origin of the Narrow Emission Lines

As shown by the previous results on SED fitting (Chaty & Rahoui 2012), the NIR continuum is dominated by the emission of the irradiated rim. In our X-Shooter spectrum, the NIR features are typically broadened at velocities on the order of FWHM = 250 km s$^{-1}$, which could be associated with the orbital velocities of the medium around the center of mass. This is the case for all the different elements we identify (H, He, Mg, Fe, etc.) and is further supported by the discussion in Section 4.2. However, below 700 nm, we find features that are noticeably different, mostly because of their much narrower width, as shown in Figure 7.

#### 4.3.1. Hα Line

The most prominent of these features is the second component of the Hα line, which lies on top of a P Cygni line profile analogous to other HⅠ lines. Its center is shifted at $-33$ km s$^{-1}$, and its FWHM is $25.6$ km s$^{-1}$. While the width is definitely much smaller than all the other lines in the spectrum (by a factor ~15), we cannot firmly tell if the velocity shift of the line center is significantly different from other hydrogen lines.

In the sgB[e] star rms 82, Seriacopi et al. (2017) reported a slight depolarization along the Hα line compared to its local continuum. This is compatible with the scenario where Hα is produced in a large volume around the star, thus suffering less scattering from the star’s envelope, hence the lower polarization.

This rules out the hypothesis that the narrow Hα line originates from the central star itself. It cannot originate from orbiting material close to the central star (i.e., the rim) because of the Keplerian velocity being much larger than its width. Further away in the disk, where the width of the line could be compatible with the lower Keplerian velocity, the medium is too cool, and the conditions to form Hα are not met. If it were directly produced in the equatorial wind itself, we would expect it to be broadened at HWHM ~ $v_{p-cyg}$, close to 400 km s$^{-1}$. We can thus exclude an equatorial origin for this component.

This leaves only the polar wind of the central star, in analogy with the results in Seriacopi et al. (2017). If this is the case, it would suggest that the polar wind is collimated and seen almost edge-on (see discussion in Section 4.5).
### 4.3.2. Forbidden Metallic Lines

Five other narrow lines are found within the range 630–680 nm. We identify them as the forbidden transitions from [O I], [N II], and [S II]. We provide their characteristics in Table A5; however, because the continuum is under the detection limit in this part of the spectrum, these values should be taken with great caution. Their FWHMs are all measured to be less than 25 km s⁻¹, and they are centered at −31 ± 5 km s⁻¹, on average.

We note that the same [O I] and [N II] transitions were identified in another sgB[e] star, LHA 115-S 18 (Clark et al. 2013), while the [S II] lines were not. The significantly narrower profile of these lines indicates that they do not come from the same region as the other, wider lines that come from the rim.

A study from Edwards et al. (1987) reports the detection of all of the aforementioned forbidden lines in what the authors argued to be the bipolar wind of T Tauri stars. The authors also referred to a book by Pottasch (1984) in which a relation between the ratio of [S II] and the electron density of the medium is provided. Since the local continuum in our spectrum is not well defined, it is hard to have a good estimate of their EQW in IGR J16318–4848. With our data, the ratio I(673)/I (671) is 0.6 ± 0.4, corresponding to an upper limit in the electron density of $n_e < 6 \times 10^2$ cm⁻³.

However, the presence of similar forbidden lines may be spurious. This is further supported by a recent spectroscopic study of a sample of sgB[e] stars (Maravelias et al. 2018). The authors reported on the profile of forbidden emission lines, one of them being in common with IGR J16318–4848 ([O I] λ630.0 nm). The analysis reveals double-peaked profiles that indicate that they originate from ringlike structures in orbit around the central star; for confirmed binaries, the profiles change over time. We do not resolve the profiles of the forbidden lines in the X-Shooter spectrum of IGR J16318–4848, so we cannot confirm if they indeed show similar structures. But their low velocity shift and width compared to the other lines is compatible with them coming from the equatorial plane of the dusty disk, as long as they originate more than 120–260 au away from the central star, which is the closest Keplerian orbit they could have, considering their width.

### 4.3.3. Nebular Origin of Narrow Lines

Another possibility is for all of the optical narrow lines to come from foreground emission of the interstellar medium. Ferland et al. (2012) measured the [O I] line widths in the Orion Nebula to be 12.6 ± 2.4 km s⁻¹ in FWHM, which is the same order of magnitude as our FWHM measurement of the λ630 nm [O I] line in the IGR J16318–4848 optical spectrum. If the narrow lines are nebular, this would set the minimal distance of IGR J16318–4848 to be 2.4 ± 0.3 kpc (see the discussion on distance in Section 5), as the velocity shift of those lines is compatible with the velocity of a star-forming region (SFR) located close to the line of sight, in the Carina-Sagittarius arm.

Given that the widths of the optical narrow lines are all very close to the instrumental resolution of X-Shooter, we suggest that performing high-resolution spectroscopy (R > 20,000) on the forbidden lines and the Hα line is necessary to clear up their origin.

### 4.4. Origin of the Flat-topped Lines

In Bertout & Magnan (1987), different ways of forming flat-topped lines are discussed, and all concern material in spherical expansion. An optically thin medium for which the turbulent velocity is much smaller than the macroscopic outflow velocity (i.e., the expansion velocity) can generate flat-topped lines. In the case of an optically thick medium, the turbulent velocity has to be much greater than the outflow velocity to form the flat top. In the first case, the half-width of the profile probes the expansion velocity; in the second case, it probes the turbulent velocity.

Out of the 15 iron lines we detect in the spectrum of IGR J16318–4848, 12 are permitted transitions, while three are forbidden. The forbidden lines are much fainter than the former; because of that, it is possible that the difference in their characteristics (blueshift, width, etc.) can be attributed to noise rather than an actual intrinsic difference.

Because the forbidden [Fe II] lines are very likely to be optically thin, regular Fe II lines sharing the same properties means the medium itself is thin. This favors the case in which the expanding medium that gives rise to flat-topped lines in IGR J16318–4848 is optically thin, and we can estimate its terminal velocity to be 250 ± 20 km s⁻¹. If we associate the broadening of the wings with the orbital motion of the medium emitting the lines, their Keplerian distance to a 25–50 M☉ central star would be 3.4–6.8 au for a circular orbit.

### 4.5. Inclination of the System

Several studies of IGR J16318–4848 provide hints of a very high inclination system; however, there is yet to be a quantitative value. According to Matt & Guainazzi (2003), the shape of the Fe and Ni Kα lines can be best reproduced considering an absorbing medium with a small covering factor, i.e., a disk seen at high inclination. Similarly, Barragán et al. (2009) argued that despite the high absorption, the lines in the Suzaku spectrum do not need Compton scattering to be reproduced, which is consistent with a nonspherical and inhomogeneous distribution of the absorbing material. In Chaty & Rahoui (2012), the authors reported on additional absorption components from silicates, which could be due to either an inaccurate absorption law or auto-absorption from a circumstellar disk seen edge-on at high inclination.

In our X-Shooter spectrum, another argument in favor of a high inclination angle comes from the Hα line. Its narrow component is likely to come from the polar wind of the central star. The profile shows no sign of deviation from a single Gaussian, and its FWHM of 25 km s⁻¹ favors a narrow opening angle. Because of the typically high velocity of such winds, a small deviation from an edge-on line of sight should greatly impact the shape of the profile. For a fiducial terminal velocity of 1000 km s⁻¹ (order of magnitude in CI Cam polar wind), and considering our resolution of 26 km s⁻¹ in the optical, we would be able to detect the blue and red components of a double-peaked Hα, corresponding to the symmetrical polar wind, if the inclination angle was ~1°.5 off the edge-on configuration (i.e., *i* < 88°5).

## 5. Distance and X-Ray Luminosity

Filliatre & Chaty (2004) estimated the distance to be between 0.9 and 6.3 kpc using the typical bolometric luminosity and temperature of sgB[e] stars. Later, Chaty &
<figure><img src="image_11.png" /><figcaption>**Figure 8.** Radial velocities of CO regions from Dame et al. (2001) around the position of IGR J16318–4848.</figcaption></figure>

Rahoui (2012) used MIR data with VLT/VISIR, along with a Herbig Ae/Be model consisting of a hot star (~20,000 K), an irradiated rim (~5500 K), and a dusty disk (~900 K) compatible with a distance of ~1.6 kpc as derived in Rahoui et al. (2008).

The spectral features in the X-Shooter spectrum hint that the local environment of the binary could be a strong contributor to the overall SED. If this is the case, the MIR data would thus probe the region where the irradiated rim contributes roughly as much as the outer dusty disk. The sheer size of these emitting regions could make them much brighter than the central star in MIR. This means that the aforementioned distance of 1.6 kpc is likely to be underestimated.

## 5.1. Nearby SFRs

As an sgB[e] HMXB, IGR J16318–4848 is a young object. It is likely to be located close to an SFR, in one of the spiral arms in the line of sight (see, e.g., Coleiro & Chaty 2013), as suggested in Filliatre & Chaty (2004).

We use velocity maps from Dame et al. (2001) around the position of IGR J16318–4848 ($l = 335.61, b = -0.447$) to look for the radial velocity of the spiral arms in the line of sight. Two cuts around the position ($b1 = -0.5$ and $b2 = -0.375$) are available (see Figure 8). The two cuts show the same three regions with different velocities. We fit each region with a Gaussian to derive the central velocity and standard deviation. We found peaks at (1) $-39 \pm 4.7$, (2) $-78 \pm 4.7$, and (3) $-115 \pm 6.4$ km s$^{-1}$.

We then retrieve the online data of Russeil (2003) to find the distances of the SFRs we could associate with these regions. Two SFRs are found close to the line of sight. The first lies in the Carina-Sagittarius arm ($l = 334.7, b = -0.1$) at $2.4 \pm 0.3$ kpc, has a radial velocity of $-33$ km s$^{-1}$, and is likely to be associated with region 1 from Figure 8. The second is in the Scutum-Crux arm ($l = 335.9, b = 0.2$) at $4.9 \pm 0.2$ kpc, has a radial velocity of $-78$ km s$^{-1}$, and is associated with region 2.

In terms of angular separation from the SFRs, IGR J16318–4848 is closer to the second (0.7°) than the first (1°); however, considering the precision of the SFR coordinates, we reckon the difference in separation cannot be a decisive argument to associate IGR J16318–4848 with either of the two SFRs.

## 5.2. Finding the Best Estimator of the Radial Velocity

Considering that the masses of the central star and compact object are likely to be 25–50 and 1.4–10 $M_{\odot}$, respectively, the maximum orbital velocity of the star around the center of mass is $29$ km s$^{-1}$ for a circular orbit. If we assume the narrow component of $H\alpha$ comes from the polar wind of the star, then its shift of $-33$ km s$^{-1}$ can yield a systemic velocity between $-62$ and $-4$ km s$^{-1}$. Without any prior knowledge of the full radial velocity curve, all of the values of systemic velocity are equally probable within this range. We reckon this estimator is rather uncertain and thus compatible with IGR J16318–4848 being associated with both SFRs 1 and 2.

If we instead consider the velocity shift from the center of the double-peaked profile of hydrogen coming from the rim, we can assume that it is not affected by any significant orbital motion around the center of mass other than its Keplerian orbit. Thus, it would provide a systemic velocity of $-66 \pm 4$ km s$^{-1}$, which favors an association with SFR 2 at 4.9 kpc.

## 5.3. Using the Latest Results from Gaia DR2

Gaia observed IGR J16318–4848, and it appears in the second data release (DR2). However, while its proper motion has been successfully measured, the parallax provided in Gaia DR2 is negative ($p = -0.5 \pm 0.3$ mas). This value cannot be explained by the suggested zero-point systematic of $-0.03$ mas and thus cannot be used to derive a distance by simply inverting the parallax.

Bailer-Jones et al. (2018) provided Bayesian inferences of distances for Gaia DR2 sources using a distance prior that assumes an exponential decrease in the space density of sources in the line of sight, with a length scale that depends on the location in the plane of the sky. This method allows one to retrieve a distance estimate, even for sources with negative parallaxes. Their method returns a distance for IGR J16318–4848 of $5.2_{-1.8}^{+2.7}$ kpc, using a length scale of 1.375 kpc. This result is compatible with locating IGR J16318–4848 close to the SFR 2 discussed earlier, at around 4.9 kpc.

The prior was chosen so that it can produce a consistent catalog of distances over all Gaia sources. However, the authors argued that any additional relevant constraint on the distance of a specific source can be used to refine the estimation. We suggest that we can use the two SFRs discussed earlier to produce a custom prior for IGR J16318–4848 and compute a more representative distance estimate.

We base ourselves on the prior used in Bailer-Jones et al. (2018) and use the same length scale as they did for IGR J16318–4848 (1.375 kpc). On top of that, we add two Gaussian priors corresponding to the two SFRs in the line of sight, weighted by their separation from IGR J16318–4848 in the plane of the sky. We produce a posterior distribution of the distance probability density shown in Figure 9.

The negative Gaia parallax measurement makes it very unlikely that the source is located below 3 kpc, even with the added SFR prior at 2.4 kpc. The integration of the distribution around its maximum returns a slightly more constrained distance estimation of $4.9_{-1.5}^{+1.9}$ kpc.

## 5.4. X-Ray Luminosity

Iyer & Paul (2017) compiled the X-ray fluxes (4–11 keV) of IGR J16318–4848 through its orbital phases, obtained by
<figure><img src="image_12.png" /><figcaption>**Figure 9.** Distance probability density with custom prior for IGR J16318–4848.</figcaption></figure>

XMM-Newton, Swift, ASCA, NuSTAR, and Suzaku, each fitted by an absorbed power law with an extra Fe-Kα emission line. The 4–11 keV flux varies from 2.29 to $11.59 \times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$. For an isotropic emission, that would correspond to an X-ray luminosity of $(1.5–8.0) \times 10^{33}$ erg s$^{-1}$ at 2.4 kpc, $(0.7–3.3) \times 10^{34}$ erg s$^{-1}$ at 4.9 kpc, and up to $(1.2–6.4) \times 10^{34}$ erg s$^{-1}$ at 6.8 kpc.

In all of the above cases, this puts IGR J16318–4848 toward the category of low-luminosity HMXBs ($L_X < 4 \times 10^{36}$ erg s$^{-1}$) and is compatible with a wind-fed system.

## 6. Fitting the SED

In this section, we model the SED of IGR J16318–4848 using the various results from spectroscopy with X-Shooter in the scope of an sgB[e] donor star surrounded by a circumstellar disk. However, we note that the spectral type of the companion was never directly tested by either its effective temperature (via $H\beta$, for instance) or stellar atmosphere modeling. For instance, Filliatre & Chaty (2004) argued that an Of/WN companion should not display as much hydrogen; however, a cool WN star may have a hydrogen mass fraction up to 50%–60% (Hainich et al. 2014; Sander et al. 2014). Moreover, if the P Cygni lines form in the dense stellar wind, the wind would produce a nonblackbody continuum that may contribute significantly to the SED.

In this section, we thus work with the sgB[e] companion hypothesis, and we briefly test the dense wind donor case in Section 7.

### 6.1. The Broadband Data Set

To complete the X-Shooter data, we retrieved archival Spitzer spectra and Herschel data to build the broadband SED from optical to MIR. Because the spectra from X-Shooter and Spitzer count more than 100,000 spectral bins in total, and we do not fit any synthetic stellar spectra, we binned the data in order to reduce the computing time and smooth out the spectral features. Since the full SED spans over 2–3 orders of magnitude in wavelength, we chose to evenly sample the X-Shooter and Spitzer data in the logarithmic space, so that we end up with 50 binned data points for each of the two spectra. We did not bin the Herschel data, since we only have two data

points, at 70 and 100 µm. This binning of the X-Shooter spectrum allows us to partially overcome the high reddening and recover the signal of the continuum down to 600 nm, as presented in Figure 10. In the following, all fits ignore the data below this wavelength. We assume that the orbit of the system is circularized and that the compact object is a $1.4 M_{\odot}$ NS.

### 6.2. Source Model, Geometry, and Absorption

#### 6.2.1. Dusty Disk Contribution

We use the 2D flat model presented in Lachaume et al. (2007) for the dusty disk, along with the temperature evolution law across the radius $T_{\text{disk}}(r)$ at a fixed temperature index $q$:

$$T_{\text{disk}}(r) = T_{\text{in}} \left( \frac{r}{R_{\text{in}}} \right)^{-q} \qquad (3)$$

Its inner radius is set at $R_{\text{rim}}$, where the temperature is $T_{\text{disk}}(R_{\text{rim}})=T_{\text{in}}$. Its total contribution is the sum of the flux radiated by annuli of radius $r$ and width $dr$ at the temperature $T_{\text{disk}}(r)$:

$$F_{\nu, \text{disk}} = \frac{2\pi \cos(i)}{D^2} \int_{R_{\text{in}}}^{R_{\text{out}}} r B_{\nu}(\nu, T(r)) dr. \qquad (4)$$

Following the results in Chaty & Rahoui (2012), we fixed $q = 0.75$, which corresponds to a viscous disk.

#### 6.2.2. Irradiated Rim Contribution

We consider the inner rim to be a portion of a cylinder of uniform temperature $T_{\text{rim}}$. Its effective area depends heavily on the viewing angle, its radius $R_{\text{rim}}$, and half-height $H_{\text{rim}}$ because of self-occultation; it is calculated using the formula given in Dullemond et al. (2001). The authors defined the parameter $\delta = \tan(i) \times H_{\text{rim}}/R_{\text{rim}}$; if $\delta > 1$ (high inclination, high $H_{\text{rim}}$-to-$R_{\text{rim}}$ ratio), the surface area is given in Equation (5), and if $\delta \le 1$ (lower inclination, low $H_{\text{rim}}$-to-$R_{\text{rim}}$ ratio), the surface area is given by Equation (6):

$$S_{\text{rim}} = 2R_{\text{rim}}^2 \cos(i)[\delta\sqrt{1 - \delta^2} + \arcsin(\delta)], \qquad (5)$$

$$S_{\text{rim}} = \pi R_{\text{rim}}^2 \cos(i). \qquad (6)$$

The total contribution of the irradiated rim is thus given by Equation (7):

$$F_{\nu, \text{rim}} = \frac{S_{\text{rim}}}{D^2} B_{\nu}(\nu, T_{\text{rim}}). \qquad (7)$$

#### 6.2.3. Stellar Contribution

The central star of radius $R_*$ has a fixed temperature of $T_* = 20,000$ K, following the same assumption made in Chaty & Rahoui (2012). Its effective area is calculated by taking into account the occultation by the rim.

The effective area is obtained by considering two different cases, depending on whether $R_*$ or $H_{\text{rim}}$ is larger. For readability, we use $\alpha = \frac{\pi}{2} - i$ in the following equations. If $R_* < H_{\text{rim}}$, there are three cases to consider. The star is fully visible (low inclination), fully occulted (high inclination), or partially occulted by the rim. Depending on the relative size of the star and disk, we find the following specific angles:

$$\alpha_m = 2 \arctan \left( \frac{-R_{\text{rim}} + \sqrt{H_{\text{rim}}^2 + R_{\text{rim}}^2 - R_*^2}}{R_* + H_{\text{rim}}} \right), \qquad (8)$$
<figure><img src="image_13.png" /><figcaption>Figure 10. Broadband SED of IGR J16318–4848 obtained from the X-Shooter, Spitzer, and Herschel data.</figcaption></figure>

<figure><img src="image_14.png" /><figcaption>Figure 11. Edge-on view of the adopted geometry for the fit of the SED. The central star is in blue, the irradiated rim in orange, and the dusty disk in red. The disk, which probably extends much further to scale, is only shown truncated on the right side for the sake of readability.</figcaption></figure>

$$ \alpha_M = 2 \arctan \left( \frac{R_{\text{rim}} - \sqrt{H_{\text{rim}}^2 + R_{\text{rim}}^2 - R_*^2}}{R_* - H_{\text{rim}}} \right). $$

If $\alpha \ge \alpha_M$, the star is fully visible, and its effective surface is a disk. If $\alpha \le \alpha_m$, the star is fully occulted by the rim. If $\alpha_m < \alpha < \alpha_M$, the star is partially occulted, and its effective surface $S^*$ is determined by

$$ \begin{aligned} h(\alpha) &= R^* - \cos(\alpha)(H_{\text{rim}} - R_{\text{rim}} \tan(\alpha)) \\ S^* &= R^{*2} \arccos\left(1 - \frac{h(\alpha)}{R^*}\right) \\ & \quad - (R^* - h)\sqrt{2 R^* h(\alpha) - h^2(\alpha)}. \end{aligned} \quad (9) $$

If $R_* > H_{\text{rim}}$, the central star is either fully visible if $\alpha > \alpha_M$ or partially occulted by the rim and disk if $\alpha < \alpha_M$. In the latter case, the effective area is again obtained with Equation (9). The stellar contribution is

$$ F_{\nu}^{*} = \frac{S^{*}}{D^{2}} B_{\nu}(\nu, T^{*}). \quad (10) $$

A complete summary of the adopted geometry is shown in Figure 11.

### 6.2.4. Absorption in the Line of Sight

The latest measurement of the absorption toward IGR J16318–4848 comes from Chaty & Rahoui (2012) at $A_V = 18.3$. As discussed in Section 3.4, we are not able to update that value with our current data. To fit the data, all of the source functions will be reddened using the formula from Cardelli et al. (1989) for optical/NIR and Chiar & Tielens

(2006) for MIR, taking into account the silicate absorption features at 9.7 and 18 µm.

### 6.3. Different Approaches for SED Fitting

For the sake of completion, we explore the two cases for which IGR J16318–4848 is located at either $D = 2.4$ or 4.9 kpc, corresponding to the SFRs we discuss in Section 5.1. The model we use to fit the SED is then governed by the following parameters: $R_*$, $T_*$, $R_{\text{rim}}$, $H_{\text{rim}}$, $T_{\text{rim}}$, $R_{\text{out}}$, $T_{\text{in}}$, $i$, $D$, and $A_V$. We perform the fit with $D$ and $A_V$ fixed. Even then, the problem is highly degenerated, and we thus use additional constraints for the fit to converge, which we discuss in the following sections.

#### 6.3.1. Fixing $T_{\text{rim}}$ at the Dust Sublimation Temperature

In this section, the rim emission is constrained by the central star. Its temperature is fixed at 1500 K, the dust sublimation temperature. Its covering factor, $H_{\text{rim}}/R_{\text{rim}}$ (hereafter $H/R$), is a variable, and its radius is computed via conservation of stellar flux to match the rim temperature.

In this case, the fit converges toward a star of radius $38 R_\odot$ kpc$^{-1}$ independently of the distance chosen, with a temperature of $6600 \pm 300$ K. The rim covering factor $H/R$ converges toward zero, suggesting that the rim has close to zero contribution in the flux. The dusty disk is poorly constrained, with an inner temperature of $800 \pm 400$ K and an outer radius of $80 \pm 70$ au kpc$^{-1}$. The inclination angle is $89^\circ 4 \pm 0^\circ 9$, which is degenerated, as it is compatible with the disk both contributing to the SED and being completely invisible.
<table><caption>**Table 4**

Results of the Fit of the SED with the Rim Radius Constrained by Spectroscopy</caption><thead><tr><th>Mass (M<sub>☉</sub>)</th><th>H/R</th><th>R<sub>out</sub> (au)</th><th>i (deg)</th><th>R<sub>*</sub> (R<sub>☉</sub>)</th><th>R<sub>rim</sub> (au)</th></tr></thead><tbody><tr><td colspan="6" style="text-align: center;">2.4 kpc</td></tr><tr><td>25</td><td>0.19 ± 0.01</td><td>110 ± 50</td><td>87.07 ± 0.02</td><td>38.7 ± 0.2</td><td>1.74 ± 0.01</td></tr><tr><td>50</td><td>0.11 ± 0.01</td><td>200 ± 100</td><td>89.298 ± 0.005</td><td>80.5 ± 0.4</td><td>3.47 ± 0.01</td></tr><tr><td colspan="6" style="text-align: center;">4.9 kpc</td></tr><tr><td>25</td><td>0.40 ± 0.02</td><td>100 ± 50</td><td>76 ± 1</td><td>33.8 ± 0.1</td><td>1.64 ± 0.01</td></tr><tr><td>50</td><td>0.18 ± 0.01</td><td>200 ± 100</td><td>86.93 ± 0.02</td><td>77.8 ± 0.3</td><td>3.46 ± 0.01</td></tr><tr><td colspan="6" style="text-align: center;">3.4 kpc</td></tr><tr><td>25</td><td>0.33 ± 0.01</td><td>110 ± 50</td><td>84.0 ± 0.3</td><td>36.4 ± 0.2</td><td>1.72 ± 0.01</td></tr><tr><td>50</td><td>0.27 ± 0.01</td><td>200 ± 100</td><td>88.53 ± 0.08</td><td>75.4 ± 0.3</td><td>3.47 ± 0.01</td></tr><tr><td colspan="6" style="text-align: center;">6.8 kpc</td></tr><tr><td colspan="6" style="text-align: center;">Converges to unrealistic parameters</td></tr><tr><td>25</td><td>0.22 ± 0.01</td><td>200 ± 100</td><td>83.97 ± 0.04</td><td>75.8 ± 0.5</td><td>3.44 ± 0.01</td></tr></tbody></table>

Note. $T_{\text{rim}} = 6700 \pm 200$ K, $T_{\text{in}} = 1370 \pm 50$ K.

The stellar parameters are by themselves not compatible with the sgB[e] hypothesis, as the stellar temperature is too cool to match an early supergiant star. The absence of any rim contribution is also in contradiction with the results of spectroscopy. This suggests that the central star does not dominate the emission in IGR J16318–4848.

### 6.3.2. Adding Extra Constraints from Spectroscopy

In this section, we make the hypothesis that the rim dominates the SED in the NIR. The parameters of the star are computed from the fitted values of the rim. The equation from Dullemond et al. (2001) that we use to compute the stellar radius is the following:

$$R^* = R_{\text{rim}} \left( \frac{T_{\text{rim}}}{T^*} \right)^2 \left( 1 + \frac{H_{\text{rim}}}{R_{\text{rim}}} \right)^{-\frac{1}{2}}. \quad (11)$$

It corresponds to the conservation of radiative flux from the star to the rim, assuming the entire flux received by the rim is used to heat it. The last term in the equation takes into account the self-irradiation of the rim, which becomes more important as its height becomes larger.

The stellar temperature is fixed at 20,000 K, as suggested in previous studies (Filliatre & Chaty 2004; Chaty & Rahoui 2012). We also note that this temperature is compatible with the models from Vink (2018) that predict slow wind velocities ($v \lesssim 500 \text{ km s}^{-1}$) for supergiant stars cooler than 21,000 K, which is the case for the wind of IGR J16318–4848.

The stellar radius will be determined by the fitted temperature, radius, and width of the rim. In this case, the rim contribution is degenerated, since three parameters influence its effective area of emission. Thus, we use spectroscopy results on the orbital velocity of the rim to constrain its absolute radius assuming a circular Keplerian orbit.

We retrieve the average projected orbital velocity $v_r \sin(i) = 113 \text{ km s}^{-1}$ from the double-peaked profiles identified in Section 4.2. Given a central mass (i.e., the stellar mass plus the compact object), we can derive a spectroscopic radius of the rim for a certain inclination angle. If we make the hypothesis that the compact object orbits within the cavity, it provides a minimum central mass so that the spectroscopic rim

radius is higher than the compact object orbit,

$$M_{\text{min}} = \frac{P v_r^3 \sin^3(i)}{2\pi G}, \quad (12)$$

with P the orbital period of the compact object and $v_r$ the orbital velocity of the rim. In the extreme case of a purely edge-on view, the rim orbital velocity of $113 \text{ km s}^{-1}$ gives an absolute minimal central mass of $12 M_{\odot}$. Below that, the rim orbits closer to the star than the compact object.

Then, for higher central masses, the lowest inclination angle that is compatible with the previous orbit considerations is

$$\sin(i_{\text{min}}) = v_r \left( \frac{P}{2\pi G M} \right)^{\frac{1}{3}}. \quad (13)$$

The central mass should not be much higher than $50 M_{\odot}$. For the compact object to orbit within the cavity, the rim orbital velocity measurement implies a minimum inclination angle of $i_{\text{min}} = 49^{\circ}$. As IGR J16318–4848 is expected to be wind-fed, this provides a lower limit on the inclination during the fit.

### 6.3.3. Results

The numerical results of the fit are presented in Table 4 for various distance values spanning 2.4–6.8 kpc. We illustrate the modeled SEDs in Figures 12 and 13 for 2.4 and 4.9 kpc only, as the 3.4 and 6.8 kpc cases cannot be distinguished from the 4.9 kpc SED.

The primary result comes from the fact that all of the fits we performed converge to a rim temperature of $T_{\text{rim}} = 6740 \pm 210$ K and an inner disk temperature of $T_{\text{in}} = 1374 \pm 47$ K.

Considering the period of $80.09 \pm 0.01$ days, the orbit of the compact object ranges from 1.06 au (at $M_{\star} = 25 M_{\odot}$) to 1.34 au ($50 M_{\odot}$). From our fit results, this places the compact object within the cavity between the central star and the rim. This further supports, along with the low X-ray luminosity (Section 5.4), that IGR J16318–4848 is a wind-fed system.

We note that for a source located at 2.4 kpc, the fit converges toward an inclination angle that is off from an edge-on view only by a few degrees, in compatibility with our assumption that $H\alpha$ comes from the polar wind of the central star (Section 4.5). However, this configuration allows a geometry that makes the central star visible and emits a continuum that is 100–1000 times fainter than the continuum of the irradiated rim at 1 µm.
<figure><img src="image_15.png" /><figcaption>**Figure 12.** Results of the SED fit for a distance of 2.4 kpc with a central mass of (a) 25 and (b) 50 $M_{\odot}$.</figcaption></figure>

As for the case of 4.9 kpc, the inclination can deviate significantly from an edge-on configuration if the central mass is too low; however, this implies a greater height for the rim, which makes the ratio $H/R$ grow outside the range of expected values (0.1–0.3; Dullemond et al. 2001). Thus, we suggest that a reasonable lower limit for the inclination of IGR J16318-4848 is $i_{min} = 76^{\circ}$, which is obtained when fitting a 25 $M_{\odot}$ star at 4.9 kpc. In this configuration, the geometry of the circumbinary material does not allow the star to be visible.

For the furthest distance estimate of 6.8 kpc, the fit converges toward very high values of $H/R$ (≥0.8) for a 25 $M_{\odot}$ star, and the rest of the parameters are not properly constrained.

Overall, the SED modeling tends to favor higher central masses because it would otherwise imply a too-large $H/R$ ratio, except if the source is close enough. We reckon that a lower limit for the inclination is 76°, but the best-fitting cases return a higher inclination in the range 86°–88°, which agrees with the conclusion we draw from the polar wind $H\alpha$ line. We note that in all cases, the inclination is high enough to allow the compact object to orbit within the cavity. This further supports that IGR J16318-4848 is a wind-fed system.

Among the supergiant B[e] X-ray binaries in the Galaxy, IGR J16318-4848 is one of the most absorbed. It has a very

<figure><img src="image_16.png" /><figcaption>**Figure 13.** Results of the SED fit for a distance of 4.9 kpc with a central mass of (a) 25 and (b) 50 $M_{\odot}$.</figcaption></figure>

complex and dynamical environment that is challenging to observe and interpret. While we managed to draw a clearer picture of its local medium, there is not yet an unambiguous explanation for the origin of the circumbinary material. It is one of the many exotic features of this source, and it could be either due to the natural evolution of the massive companion or result from binary interaction.

# 7. Modeling the Stellar Atmosphere and Wind

## 7.1. The PoWR Model

As an alternative approach to the simple 20,000 K sgB[e] companion with a minor contribution presented in Section 6, we test here the assumption of a major contribution from a massive donor with the same temperature and an optically thick wind responsible for the P Cygni lines. The stellar atmosphere modeling was performed using the Potsdam Wolf-Rayet (PoWR) model atmosphere code. This code solves the comoving-frame radiative transfer together with the solution of the population numbers and the temperature stratification in full non-LTE (e.g., Hamann & Gräfener 2003). It fully accounts for iron-line blanketing (Gräfener et al. 2002) and optically thin density inhomogeneities (“microclumping”; Hamann & Koesterke 1998).
<figure><img src="image_17.png" /><figcaption>**Figure 14.** Comparison of the computed SED from the PoWR code with available photometric data, assuming a 2.4 kpc distance.</figcaption></figure>

The quasi-hydrostatic part is treated self-consistently (Sander et al. 2015), while a β-law is assumed in the supersonic domain. The PoWR model provides a stellar atmosphere stratification, as well as an emergent spectrum over a wide wavelength range, from the (E)UV to the MIR.

## 7.2. Results of the Modeling

In Figure 14, we show the overall SED obtained with the model compared to the available Gaia and JHK photometry assuming a distance set to the minimal value we derive of 2.4 kpc. Moreover, Figures 15 and 16 highlight selected parts of the normalized model spectrum compared to the available observation to illustrate the contribution of the star to the hydrogen emission lines and prominent P Cygni lines formed in the stellar wind. We note that we were not able to reproduce the flat-topped Fe II lines. In Table 5, we present the results for the stellar and wind parameters. For these, we assumed a fixed effective temperature of T<sub>*</sub> = 20,500 K at a Rosseland continuum optical depth of 20, a hydrogen mass fraction in line with Section 4.1 as well as a radius and mass combination from the constraints in Section 6.3.3 yielding the result closest to a globally consistent energy budget.

The PoWR model results show P Cygni profiles for hydrogen lines; however, we are unable to reproduce their intensity as shown in Figure 15. We do recover the intensity for the helium emission lines, but the modeled P Cygni absorption is usually weaker than the observations (Figure 16). However, we did not expect to fully reproduce the line shapes, as this would require a full decomposition into a stellar and a disk component, which is beyond the scope of the present paper. According to these results, we infer that the terminal velocity of the wind can be reasonably constrained to be between 300 and 400 km s<sup>-1</sup>.

## 7.3. Discussion

In this first approach at SED modeling of IGR J16318–4848 with the PoWR code, we consider that the radiation mainly comes from the star and its close environment. We reckon that it is a sensible alternative to the minor donor contribution considered in Section 6 and that it brings interesting results, especially concerning the reproduction of line profiles.

The terminal velocity of the wind inferred with the model atmosphere is in agreement with the one derived in Section 4.2. Concerning the hydrogen lines, the fact that we are not able to

<figure><img src="image_18.png" /><figcaption>**Figure 15.** Hydrogen lines reconstructed from the PoWR model (red) vs. the normalized input spectrum (blue).</figcaption></figure>

$$T_{\text{rim}} = 6700 \pm 200 \text{ K}, T_{\text{in}} = 1370 \pm 50 \text{ K}.$$

<figure><img src="image_19.png" /></figure>

$$v_{\infty} = 300 \text{ km s}^{-1}$$

<figure><img src="image_20.png" /></figure>

$$v_{\infty} = 400 \text{ km s}^{-1}$$

<caption>**Figure 16.** Wind lines reconstructed from the PoWR model for prominent helium transitions with different terminal velocities (red) vs. the normalized input spectrum (blue).</caption>

reproduce them correctly supports our findings in Section 4.2 that there is considerable contribution from the irradiated rim.

We suggest that the hydrogen lines produced by the PoWR code can be attributed to the extra central line discussed in Section 4.2 for which we previously failed to identify the nature, and in fact come from the vicinity of the central star.
<table><caption>**Table 5**
Parameters Used to Model the Atmosphere and Wind of IGR J16318–4848 with PoWR</caption><tr><td>T<sup>*</sup><sub>eff</sub></td><td>20,500 K</td></tr><tr><td>T<sub>2/3</sub></td><td>18,200 K</td></tr><tr><td>X<sub>H</sub></td><td>0.5</td></tr><tr><td>R<sup>*</sup></td><td>51 R<sub>☉</sub></td></tr><tr><td>M<sup>*</sup></td><td>25 M<sub>☉</sub></td></tr><tr><td>log (L/L<sub>☉</sub>)</td><td>5.6</td></tr><tr><td>log (Ṁ)</td><td>-4.8</td></tr><tr><td>v<sub>∞</sub></td><td>400 km s<sup>-1</sup></td></tr><tr><td>D<sub>∞</sub></td><td>10</td></tr><tr><td>v<sub>mic</sub></td><td>14 km s<sup>-1</sup></td></tr></table>

Also, this would indicate that the line of sight allows us to see at least part of the central star. However, our multicomponent SED modeling in Section 6 does not allow us to see the star if it is located at 4.9 kpc. This may suggest that the source is closer or that our geometrical model is not entirely accurate.

As for the helium lines, the stellar atmosphere and wind modeling reproduce their overall profile rather well. The P Cygni absorption is still slightly more difficult to fully recover, probably because of the complex circumbinary environment in the line of sight that is not taken into account during the modeling (i.e., the dusty disk). Most of the helium lines are still much better reproduced than hydrogen in terms of intensity, width, and P Cygni absorption with the PoWR code. This result might indicate that the central star has blown off part of its hydrogen envelope through intense stellar wind. As such, the photosphere might be enhanced in helium, and the circumbinary medium might have been partly formed from that wind, hence the intense hydrogen emission lines from the irradiated rim. Following the discussion in Section 4.1, this rather suggests an intrinsic origin for the enhanced He/H ratio.

Despite calculating outward to $10^5 R_*$, the PoWR model could not reproduce any of the flat-topped Fe II lines. While this would support the hypothesis that these transitions arise not from the star or its wind but rather from the circumbinary disk, their absence in the synthetic spectrum may also be due to the different treatment of the large amount of iron lines compared to other elements in PoWR. This will need to be investigated in a separate project.

The mass and radius of the central star recovered from the PoWR code are compatible with the SED fitting. However, the mass is on the lower end of our estimations, and our fit at $25 M_☉$ converges toward a rim height to rim radius ratio that is higher than the expected range. The effective temperature of the photosphere $T_{\text{eff}}^*$ of 20,500 K is motivated by the value of 20,000 K suggested in Filliatre & Chaty (2004) and Chaty & Rahoui (2012).

This first attempt at modeling the spectrum of IGR J16318–4848 hints that the exact nature of its donor is still up for debate and that a dense wind donor (a late hydrogen-rich Of/WN) is a possibility. To clear it up, we reckon that further modeling with a complete parameter space exploration along with the subtraction of potential disk contribution is required.

## 8. Comparison with the sgB[e] HMXB CI Cam

The first sgB[e] HMXB to be identified was CI Cam, and it has been extensively studied since its discovery (Smith & Remillard 1998). While both systems show different behaviors, we reckon the comparison with IGR J16318–4848 is relevant, as these binaries might be of a very similar nature or correspond to

different phases of the same evolutionary path. Bartlett et al. (2019) used recent Swift/XRT observations of CI Cam spanning ~150 days and suggested a variability timescale between 75 and 100 days, the same order of magnitude as the orbital period of 80 days derived for IGR J16318–4848 by Iyer & Paul (2017). Robinson et al. (2002) suggested the presence of a high-velocity outflow (1000–2500 km s⁻¹) associated with a polar outflow seen almost pole-on (Hynes et al. 2002). Bartlett et al. (2019) derived a column density of $6 \times 10^{22}–2 \times 10^{24} \text{ cm}^{-2}$, which reaches the estimated column density values of IGR J16318–4848. However, the dust extinction in CI Cam only reaches up to $A_V = 4$ mag (Hynes et al. 2002), much lower than the extinction for IGR J16318–4848 ($A_V = 18.3$; Chaty & Rahoui 2012). It is possible that this difference only comes from the viewing angle of those systems, which is inferred to be almost pole-on for CI Cam and almost edge-on for IGR J16318–4848. Both having $N_H$ values much higher than their optical interstellar absorption suggests that both have an absorbing medium local to the accretion region, i.e., around the compact object. The much higher $A_V$ in IGR J16318–4848 compared to CI Cam would be due to the presence of the equatorial circumbinary disk in the line of sight. Another notable feature in both sgB[e] systems is the presence of Fe II flat-topped lines in their optical to infrared spectrum. However, while these profiles are present in quiescence for IGR J16318–4848, they were only observed in CI Cam during its 1998 outburst; instead, in quiescence, the profiles were shown to be double-peaked (or with a central depression; see Hynes et al. 2002). This could stress a significant difference between the two binaries, as CI Cam only displays a spherically symmetric outflow during outburst (flat-topped lines) and returns to an axisymmetric state in quiescence (double-peaked lines), while IGR J16318–4848 has no record of such behavior, as the iron outflow persists in spherical symmetry. Another difference lies in the projected velocity of the expanding medium, which was measured to be $32 \text{ km s}^{-1}$ in CI Cam (Robinson et al. 2002) and $250 \pm 20 \text{ km s}^{-1}$ in IGR J16318–4848 (this study), about an order of magnitude higher.

## 9. Conclusion

We presented spectroscopy performed on IGR J16318–4848 with unprecedented resolution and coverage toward the optical band. With the analysis of spectral features and the modeling of the broadband SED using archival Spitzer and Herschel data, we obtain the following results.

1. The inclination of the system is higher than 76° and reaches up to 86°–88° for the best-fitting models.

2. Based on SFR associations and the DR2 of Gaia, we infer the distance to IGR J16318–4848 to be $4.9^{+1.9}_{-1.5}$ kpc.

3. We confirm the presence of P Cygni profiles in H I and He I lines that likely probe the equatorial wind of the central star expanding at velocities up to $370 \text{ km s}^{-1}$.

4. The H I lines with a sufficient signal-to-noise ratio can be reproduced by a double-peaked profile originating from the orbital motion of the rim at $v_r \sin(i) = 113 \pm 4 \text{ km s}^{-1}$.

5. We resolve the flat-topped profiles of previously identified Fe II and [Fe II] lines in emission. This indicates that they originate from an optically thin medium undergoing spherical expansion at $250 \pm 20 \text{ km s}^{-1}$. If their wing broadening comes from orbital motion, it will locate the origin of the wind at 3.4–6.8 au away from the central star, meaning that the iron lines form in a disk wind from the dusty equatorial disk.
6. The Hα line has an extra narrow component that, given the high inclination, could be associated with a fast polar wind seen almost edge-on.

7. Optical forbidden lines from [O I], [N II], and [S II] display narrow profiles that suggest they could originate either from far away in the dusty disk (120–260 au, or 700–740 R<sub>*</sub>) or from the foreground nebular emission of a spiral arm.

8. In the case of an sgB[e] companion, the irradiated rim temperature is fitted to be T<sub>rim</sub> = 6740 ± 210 K, and the temperature of the inner viscous disk is T<sub>in</sub> = 1374 ± 47 K. In this model, those structures outshine by far the central star, which contributes at best 10% in optical flux.

9. In the case where the compact object is a 1.4 M<sub>☉</sub> NS, it is likely to orbit within the central cavity. Adding the low X-ray luminosity we derived, this would suggest that the system is wind-fed, though we cannot rule out tidal interactions between the compact object and the inner disk.

10. We perform stellar atmosphere and wind modeling of the X-Shooter optical to NIR spectrum with the PoWR code, assuming a donor with a dense wind. The results are overall in good agreement with the observations, but further modeling with complete parameter space exploration might shed new light on the exact nature of the donor.

11. We compare IGR J16318–4848 with CI Cam and note that both systems might be of a similar nature, and that the differences in some of their key features could be explained by the very different viewing angle (edge-on for IGR J16318–4848, pole-on for CI Cam).

This work was supported by the Centre National d'Etudes Spatiales (CNES), based on observations obtained with the Multi-wavelength INTEGRAL NeTwork (MINE). S.C. is also grateful to the LabEx UnivEarthS for the funding of Interface project “Galactic binaries toward merging.” This publication makes use of VOSA, developed under the Spanish Virtual Observatory project supported by the Spanish MICINN through grant AyA2011-24052; NASA’s Astrophysics Data System Bibliographic Services, operated by the Smithsonian Astrophysical Observatory under NASA Cooperative Agreement NNX16AC86A; and data from the European Space Agency (ESA) mission Gaia (https://www.cosmos.esa.int/gaia), processed by the Gaia Data Processing and Analysis Consortium (DPAC; https://www.cosmos.esa.int/web/gaia/dpac/consortium). Funding for the DPAC has been provided by national institutions, in particular the institutions participating in the Gaia Multilateral Agreement. A.S. would like to thank STFC for funding under grant No. ST/R000565/1.

## Appendix
List of Identified Transitions

Tables A1–A5 list the fitted properties of lines we identified in the X-Shooter spectrum. Each table provides the reference wavelength in air (λ<sub>0</sub>; nm), the measured heliocentric velocity of the line center (V<sub>h</sub>; km s<sup>-1</sup>), the flux density of the line (Jy), the FWHM (km s<sup>-1</sup>), and the EQW (nm). For H I and He I, we also provide the heliocentric velocity difference between the emission and absorption components of the line (V<sub>PC</sub>; km s<sup>-1</sup>). Typical errors are 5 km s<sup>-1</sup> in velocity and 0.002 nm in EQW.

<table><caption>**Table A1**
Identified H I Lines</caption><thead><tr><th>Line</th><th>λ<sub>0</sub></th><th>V<sub>h</sub></th><th>Flux</th><th>FWHM</th><th>EQW</th><th>V<sub>PC</sub></th></tr></thead><tbody><tr><td>Hα 3-2</td><td>656.279</td><td>-73</td><td>2.31e<sup>-4</sup></td><td>383</td><td>-27.317</td><td>165</td></tr><tr><td>Pa 23-3</td><td>834.554</td><td>-43</td><td>1.34e<sup>-4</sup></td><td>224</td><td>-0.097</td><td>...</td></tr><tr><td>Pa 22-3</td><td>835.900</td><td>-13</td><td>2.47e<sup>-4</sup></td><td>266</td><td>-0.207</td><td>...</td></tr><tr><td>Pa 21-3</td><td>837.448</td><td>5</td><td>2.57e<sup>-4</sup></td><td>299</td><td>-0.264</td><td>...</td></tr><tr><td>Pa 20-3</td><td>839.24</td><td>-15</td><td>2.45e<sup>-4</sup></td><td>322</td><td>-0.274</td><td>...</td></tr><tr><td>Pa 19-3</td><td>841.332</td><td>-17</td><td>3.99e<sup>-4</sup></td><td>274</td><td>-0.358</td><td>...</td></tr><tr><td>Pa 17-3</td><td>846.726</td><td>-28</td><td>6.65e<sup>-4</sup></td><td>358</td><td>-0.661</td><td>254</td></tr><tr><td>Pa 16-3</td><td>850.249</td><td>-37</td><td>5.82e<sup>-4</sup></td><td>262</td><td>-0.373</td><td>...</td></tr><tr><td>Pa 15-3</td><td>854.538</td><td>-46</td><td>9.04e<sup>-4</sup></td><td>308</td><td>-0.746</td><td>352</td></tr><tr><td>Pa 14-3</td><td>859.839</td><td>-37</td><td>7.87e<sup>-4</sup></td><td>275</td><td>-0.454</td><td>272</td></tr><tr><td>Pa 13-3</td><td>866.502</td><td>-41</td><td>1.30e<sup>-3</sup></td><td>277</td><td>-0.745</td><td>313</td></tr><tr><td>Pa 12-3</td><td>875.046</td><td>-48</td><td>1.49e<sup>-3</sup></td><td>356</td><td>-0.921</td><td>224</td></tr><tr><td>Pa 11-3</td><td>886.289</td><td>-33</td><td>2.15e<sup>-3</sup></td><td>301</td><td>-0.928</td><td>257</td></tr><tr><td>Pa 10-3</td><td>901.533</td><td>-83</td><td>3.17e<sup>-3</sup></td><td>377</td><td>-1.182</td><td>202</td></tr><tr><td>Pa 8-3</td><td>954.618</td><td>-33</td><td>9.80e<sup>-3</sup></td><td>324</td><td>-1.415</td><td>275</td></tr><tr><td>Pa 7-3</td><td>1004.937</td><td>-33</td><td>3.22e<sup>-2</sup></td><td>322</td><td>-2.424</td><td>277</td></tr><tr><td><b>Pa 5-3</b></td><td><b>1281.808</b></td><td><b>-67</b></td><td>...</td><td>...</td><td><b>-11.4</b></td><td><b>340</b></td></tr><tr><td>Br 25-4</td><td>1496.733</td><td>-43</td><td>4.70e<sup>-2</sup></td><td>295</td><td>-0.289</td><td>...</td></tr><tr><td>Br 24-4</td><td>1500.086</td><td>-37</td><td>4.38e<sup>-2</sup></td><td>301</td><td>-0.275</td><td>...</td></tr><tr><td>Br 23-4</td><td>1503.904</td><td>-48</td><td>5.56e<sup>-2</sup></td><td>318</td><td>-0.360</td><td>...</td></tr><tr><td>Br 22-4</td><td>1508.277</td><td>-81</td><td>8.26e<sup>-2</sup></td><td>203</td><td>-0.338</td><td>270</td></tr><tr><td>Br 21-4</td><td>1513.322</td><td>-77</td><td>7.10e<sup>-2</sup></td><td>380</td><td>-0.537</td><td>198</td></tr><tr><td>Br 20-4</td><td>1519.184</td><td>-37</td><td>7.61e<sup>-2</sup></td><td>281</td><td>-0.443</td><td>...</td></tr><tr><td>Br 19-4</td><td>1526.054</td><td>-84</td><td>1.03e<sup>-1</sup></td><td>211</td><td>-0.426</td><td>...</td></tr><tr><td>Br 18-4</td><td>1534.179</td><td>-51</td><td>1.08e<sup>-1</sup></td><td>297</td><td>-0.632</td><td>...</td></tr><tr><td>Br 17-4</td><td>1543.892</td><td>-54</td><td>1.21e<sup>-1</sup></td><td>330</td><td>-0.786</td><td>258</td></tr><tr><td>Br 16-4</td><td>1555.645</td><td>-84</td><td>1.57e<sup>-1</sup></td><td>410</td><td>-1.161</td><td>189</td></tr><tr><td>Br 15-4</td><td>1570.066</td><td>-47</td><td>1.72e<sup>-1</sup></td><td>333</td><td>-1.050</td><td>248</td></tr><tr><td>Br 14-4</td><td>1588.054</td><td>-48</td><td>2.24e<sup>-1</sup></td><td>331</td><td>-1.386</td><td>242</td></tr><tr><td>Br 13-4</td><td>1610.931</td><td>-49</td><td>2.60e<sup>-1</sup></td><td>323</td><td>-1.524</td><td>244</td></tr><tr><td>Pa 4-3</td><td>1875.101</td><td>-61</td><td>1.01e<sup>1</sup></td><td>343</td><td>-40.337</td><td>...</td></tr><tr><td>Br 8-4</td><td>1944.556</td><td>-48</td><td>1.12</td><td>326</td><td>-4.167</td><td>304</td></tr><tr><td><b>Br 7-4</b></td><td><b>2165.529</b></td><td><b>-73</b></td><td>...</td><td>...</td><td><b>-7.62</b></td><td><b>373</b></td></tr><tr><td>Pf 25-5</td><td>2373.729</td><td>-61</td><td>9.96e<sup>-2</sup></td><td>407</td><td>-0.479</td><td>...</td></tr><tr><td>Pf 24-5</td><td>2382.173</td><td>-41</td><td>1.27e<sup>-1</sup></td><td>380</td><td>-0.577</td><td>...</td></tr><tr><td>Pf 23-5</td><td>2391.815</td><td>-65</td><td>1.77e<sup>-1</sup></td><td>308</td><td>-0.643</td><td>...</td></tr><tr><td>Pf 21-5</td><td>2415.726</td><td>-57</td><td>2.02e<sup>-1</sup></td><td>339</td><td>-0.737</td><td>...</td></tr><tr><td>Pf 20-5</td><td>2430.699</td><td>-54</td><td>2.37e<sup>-1</sup></td><td>324</td><td>-0.920</td><td>...</td></tr><tr><td>Pf 19-5</td><td>2448.332</td><td>-49</td><td>2.82e<sup>-1</sup></td><td>333</td><td>-1.133</td><td>...</td></tr></tbody></table>

<table><caption>**Table A2**
Identified He I Lines</caption><thead><tr><th>λ<sub>0</sub></th><th>V<sub>h</sub></th><th>Flux</th><th>FWHM</th><th>EQW</th><th>V<sub>PC</sub></th></tr></thead><tbody><tr><td>706.518</td><td>-123</td><td>7.75e<sup>-5</sup></td><td>342</td><td>-1.820</td><td>115</td></tr><tr><td>728.135</td><td>-127</td><td>5.30e<sup>-5</sup></td><td>327</td><td>-0.661</td><td>133</td></tr><tr><td>1031.122</td><td>-118</td><td>6.86e<sup>-3</sup></td><td>450</td><td>-0.633</td><td>147</td></tr><tr><td>1196.904</td><td>-116</td><td>2.34e<sup>-2</sup></td><td>437</td><td>-0.608</td><td>179</td></tr><tr><td>1252.751</td><td>-73</td><td>3.28e<sup>-2</sup></td><td>223</td><td>-0.315</td><td>...</td></tr><tr><td>1279.050</td><td>-118</td><td>6.36e<sup>-2</sup></td><td>373</td><td>-0.326</td><td>316</td></tr><tr><td>1278.491</td><td>14</td><td>6.51e<sup>-2</sup></td><td>386</td><td>-0.923</td><td>312</td></tr><tr><td>1296.843</td><td>-77</td><td>2.82e<sup>-2</sup></td><td>203</td><td>-0.211</td><td>193</td></tr><tr><td>1700.234</td><td>-111</td><td>2.49e<sup>-1</sup></td><td>278</td><td>-1.100</td><td>164</td></tr><tr><td>2058.129</td><td>-133</td><td>1.81</td><td>299</td><td>-6.120</td><td>135</td></tr><tr><td>2112.002</td><td>-119</td><td>9.89e<sup>-2</sup></td><td>296</td><td>-0.471</td><td>360</td></tr><tr><td>2161.701</td><td>-111</td><td>1.01e<sup>-1</sup></td><td>426</td><td>-0.164</td><td>280</td></tr></tbody></table>

Bold values in Table A1 indicate lines that were modelled more accurately; their fitted parameters can be found in Section 4.2. Hydrogen lines in bold benefited from more accurate modeling; their set of parameters is available in Section 4.2.
<table><caption>**Table A3**
Identified Iron Transitions</caption><thead><tr><th></th><th>Term</th><th>λ</th><th>V<sub>helio</sub></th><th>Flux (mJy)</th><th>EQW</th><th>FW</th><th>GW</th></tr></thead><tbody><tr><td>Fe II</td><td>z4F<sub>0</sub>-b4G</td><td>995.631</td><td>-79 ± 3</td><td>3.4 ± 0.05</td><td>-0.254</td><td>220 ± 4</td><td>94 ± 8</td></tr><tr><td>Fe II</td><td>z4F<sub>0</sub>-b4G</td><td>999.758</td><td>-63 ± 2</td><td>434.0 ± 0.17</td><td>-2.914</td><td>217 ± 2</td><td>117 ± 3</td></tr><tr><td>Fe II</td><td>z4D<sub>0</sub>-d2F</td><td>1017.392</td><td>-83 ± 7</td><td>6.1 ± 0.22</td><td>-0.647</td><td>252 ± 7</td><td>104 ± 18</td></tr><tr><td>Fe II</td><td>z4F<sub>0</sub>-d2F</td><td>1043.476</td><td>-75 ± 3</td><td>3.7 ± 0.08</td><td>-0.332</td><td>265 ± 3</td><td>69 ± 8</td></tr><tr><td>Fe II</td><td>z4F<sub>0</sub>-b4G</td><td>1050.150</td><td>-82 ± 4</td><td>26.0 ± 0.42</td><td>-0.047</td><td>236 ± 4</td><td>97 ± 9</td></tr><tr><td>Fe II</td><td>z4F<sub>0</sub>-b4G</td><td>1112.558</td><td>-72 ± 2</td><td>38.1 ± 0.37</td><td>-1.811</td><td>233 ± 2</td><td>93 ± 5</td></tr><tr><td>[Fe II]</td><td>a6D-a4D</td><td>1256.680</td><td>-47 ± 3</td><td>11.7 ± 0.2</td><td>-0.243</td><td>295 ± 4</td><td>65 ± 8</td></tr><tr><td>Fe II</td><td>z4F<sub>0</sub>-c3F</td><td>1687.320</td><td>-73 ± 2</td><td>495.0 ± 1.7</td><td>-3.612</td><td>257 ± 2</td><td>66 ± 4</td></tr><tr><td>Fe II</td><td>z4D<sub>0</sub>-c4F</td><td>1741.401</td><td>-88 ± 1</td><td>152.9 ± 2.0</td><td>-1.026</td><td>282 ± 6</td><td>138 ± 9</td></tr><tr><td>Fe II</td><td>z4F<sub>0</sub>-c3F</td><td>1974.611</td><td>-74 ± 3</td><td>282.1 ± 4.7</td><td>-1.598</td><td>266 ± 3</td><td>80 ± 7</td></tr><tr><td>Fe II</td><td>z4D<sub>0</sub>-c4F</td><td>1986.841</td><td>-76 ± 3</td><td>81.1 ± 1.1</td><td>-0.484</td><td>265 ± 3</td><td>80 ± 6</td></tr><tr><td>[Fe II]</td><td>a4P-a2P</td><td>2046.007</td><td>-55 ± 7</td><td>24.0 ± 0.64</td><td>-0.139</td><td>280 ± 7</td><td>55 ± 11</td></tr><tr><td>Fe II</td><td>z4F<sub>0</sub>-c3F</td><td>2088.810</td><td>-76 ± 2</td><td>299.0 ± 1.6</td><td>-1.690</td><td>257 ± 2</td><td>81 ± 4</td></tr><tr><td>[Fe II]</td><td>a2G-a2H</td><td>2223.760</td><td>-40 ± 6</td><td>17.9 ± 0.6</td><td>-0.096</td><td>286 ± 6</td><td>52 ± 15</td></tr><tr><td>Fe II</td><td>z4D<sub>3/2</sub>-c4P<sub>3/2</sub></td><td>2240.152</td><td>-62 ± 3</td><td>28.4 ± 0.5</td><td>-0.140</td><td>274 ± 3</td><td>54 ± 7</td></tr></tbody></table>

Note. (FW: HWHM of the flat profile; GW: HWHM of the Gaussian broadening. Both are in km s⁻¹.)

<table><caption>**Table A4**
Identified Transitions of Mg II</caption><thead><tr><th>Term</th><th>λ<sub>0</sub></th><th>V<sub>h</sub></th><th>Flux</th><th>FWHM</th><th>EQW</th></tr></thead><tbody><tr><td>2P<sub>0</sub>-2D</td><td>787.705</td><td>-94.5</td><td>2.78e<sup>-4</sup></td><td>413</td><td>-0.514</td></tr><tr><td>2P<sub>0</sub>-2D</td><td>789.637</td><td>-59.4</td><td>4.62e<sup>-4</sup></td><td>329</td><td>-0.708</td></tr><tr><td>2P<sub>0</sub>-2S</td><td>821.398</td><td>-10.3</td><td>4.39e<sup>-4</sup></td><td>236</td><td>-0.206</td></tr><tr><td>2P<sub>0</sub>-2S</td><td>823.464</td><td>-117.9</td><td>5.84e<sup>-4</sup></td><td>355</td><td>-0.411</td></tr><tr><td>2S-2P<sub>0</sub></td><td>921.825</td><td>-125.3</td><td>1.02e<sup>-2</sup></td><td>270</td><td>-0.988</td></tr><tr><td>2S-2P<sub>0</sub></td><td>924.426</td><td>-52.3</td><td>8.3e<sup>-3</sup></td><td>325</td><td>-0.980</td></tr><tr><td>2D-2P<sub>0</sub></td><td>1091.424</td><td>-58.0</td><td>4.5e<sup>-2</sup></td><td>283</td><td>-1.484</td></tr><tr><td>2D-2P<sub>0</sub></td><td>1095.177</td><td>17.9</td><td>2.0e<sup>-2</sup></td><td>273</td><td>-0.589</td></tr><tr><td>2S-2P<sub>0</sub></td><td>2136.90</td><td>-97.4</td><td>2.37e<sup>-1</sup></td><td>351</td><td>-0.948</td></tr><tr><td>2S-2P<sub>0</sub></td><td>2143.22</td><td>-82.4</td><td>1.27e<sup>-1</sup></td><td>350</td><td>-0.509</td></tr><tr><td>2D-2P<sub>0</sub></td><td>2412.46</td><td>-110.1</td><td>1.31e<sup>-1</sup></td><td>344</td><td>-0.534</td></tr><tr><td>2D-2P<sub>0</sub></td><td>2404.15</td><td>-154.9</td><td>3.82e<sup>-1</sup></td><td>353</td><td>-1.600</td></tr></tbody></table>

<table><caption>**Table A5**
Optical Forbidden Transitions</caption><thead><tr><th></th><th>Term</th><th>λ<sub>0</sub></th><th>V<sub>h</sub></th><th>Flux</th><th>FWHM</th><th>EQW</th></tr></thead><tbody><tr><td>[O I]</td><td>3P-1D</td><td>630.030</td><td>-32</td><td>1.0e<sup>-4</sup></td><td>11</td><td>-2.21</td></tr><tr><td>[N II]</td><td>3P-1D</td><td>654.804</td><td>-30</td><td>1.6e<sup>-4</sup></td><td>17</td><td>-0.72</td></tr><tr><td>[N II]</td><td>3P-1D</td><td>658.346</td><td>-32</td><td>5.1e<sup>-4</sup></td><td>6</td><td>-2.09</td></tr><tr><td>[S II]</td><td>4S<sub>0</sub>-2D<sub>0</sub></td><td>671.644</td><td>-29</td><td>3.4e<sup>-4</sup></td><td>15</td><td>-1.46</td></tr><tr><td>[S II]</td><td>4S<sub>0</sub>-2D<sub>0</sub></td><td>673.081</td><td>-31</td><td>1.9e<sup>-4</sup></td><td>8</td><td>-0.75</td></tr></tbody></table>

ORCID iDs

Francis Fortin https://orcid.org/0000-0003-3642-2267

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