<figure><img src="image_1.png" /></figure>

# The Secondary Component of the Extremely Low Mass Ratio Massive Binary HD 165246 Revealed by GHOST

C. N. Rodriguez<fnref n="1" />, T. Ansin<fnref n="2" />,<fnref n="3" />, G. A. Ferrero<fnref n="2" />,<fnref n="3" />, O. G. Benvenuto<fnref n="2" />,<fnref n="3" />,<fnref n="19" />, R. Gamen<fnref n="2" />,<fnref n="3" />, N. I. Morrell<fnref n="4" />, J. I. Arias<fnref n="5" />, R. E. Higa<fnref n="2" />,<fnref n="3" />, C. Putkuri<fnref n="2" />,<fnref n="3" />, C. Johnston<fnref n="6" />,<fnref n="7" />,<fnref n="8" />, K. Chiboucas<fnref n="9" />, E. Deibert<fnref n="1" />, C. R. Hayes<fnref n="10" />, J.-E. Heo<fnref n="1" />, M. Jeong<fnref n="11" />, V. Kalari<fnref n="1" />, E. Martioli<fnref n="12" />,<fnref n="13" />, V. M. Placco<fnref n="14" />, K. A. Venn<fnref n="15" />, S. Xu<fnref n="19" />, R. Diaz<fnref n="1" />, M. Gomez-Jimenez<fnref n="1" />, R. Ruiz-Carmona<fnref n="1" />, C. Simpson<fnref n="9" />, A. W. McConnachie<fnref n="10" />, J. Pazder<fnref n="10" />,<fnref n="15" />, G. Burley<fnref n="10" />, T. A. M. Berg<fnref n="10" />,<fnref n="16" />, J. G. Robertson<fnref n="17" />,<fnref n="18" />, K. Labrie<fnref n="9" />, and J. Thomas-Osip<fnref n="1" />

## Abstract

In the pursuit of understanding the multiplicity of massive stars, the OWN Survey has undertaken spectroscopic monitoring of hundreds of targets over the past 20 yr. This effort has led to the discovery of new single-lined spectroscopic binaries. Characterizing the unseen companions of these systems is crucial to deriving precise stellar parameters, including the mass of each component, and constraining formation models for massive binary systems. In this paper, we aim to physically and evolutionarily characterize the stellar components of the massive binary HD 165246. We analyzed spectra obtained by the OWN Survey and used Kepler's second mission (K2) photometry, from which we determined the physical parameters of the system. Thanks to the high-resolution capabilities of the new echelle spectrograph Gemini/GHOST, we could directly detect spectral features of the secondary component for the first time. A very low mass ratio for the system is confirmed ($q \sim 0.16$), making this binary one of the few extremely low mass ratio systems known. The primary star has a mass of $M_a = 22 \pm 1 M_\odot$ and a radius of $R_a = 7.0 \pm 0.1 R_\odot$, while the secondary star has a mass of $M_b = 3.4 \pm 0.1 M_\odot$ and a radius of $R_b = 2.22 \pm 0.02 R_\odot$. Stellar evolution models indicate that the system has undergone approximately two million years of evolution since its formation. Moreover, the secondary star is identified as a pre-main-sequence object, progressing toward the zero-age main sequence. This paper illustrates the power of high-resolution spectrographs such as GHOST in finding much lower mass, previously unseen stellar components of binary star systems.

*Unified Astronomy Thesaurus concepts: Massive stars (732); Early-type stars (430); Fundamental parameters of stars (555); Close binary stars (254); Eclipsing binary stars (444)*

## 1. Introduction

Massive stars end their evolution in a gravitational core collapse, giving rise to the formation of a neutron star or a black hole, in most cases through a supernova explosion. Stars that reach this end, start with an initial mass of at least $8–10 M_\odot$ at the

zero-age main sequence (ZAMS; A. J. T. Poelarends et al. 2008). This also means they have an O or earlier than B3 spectral type in their first evolutionary stage. These are stars of great astrophysical importance for several reasons, including their key role in the chemical enrichment of galaxies, the dynamics of the interstellar (IS) medium, their contribution to the dynamo effect responsible for galactic magnetic fields, the ionization of gas in H II regions, the location of galactic spiral arms, and the origin of binary systems capable of emitting gravitational waves.

Observational evidence shows that most massive stars are part of multiple systems that are close enough to interact during their lifetime (H. Sana et al. 2012; H. A. Kobulnicky et al. 2014; R. H. Barbá et al. 2017; G. Banyard et al. 2022). For O-type stars, the fraction of binarity is in the range of 50%–70%

<figure><img src="image_2.png" /></figure>

<footnote n="1">¹ International Gemini Observatory/NSF NOIRLab, Casilla 603, La Serena, Chile; cinthya.rodriguez@noirlab.edu</footnote>

<footnote n="2">² Instituto de Astrofísica de La Plata, CONICET-UNLP, Paseo del Bosque S/N, B1900FWA La Plata, Argentina</footnote>

<footnote n="3">³ Facultad de Ciencias Astronómicas y Geofísicas, Universidad Nacional de La Plata, Paseo del Bosque S/N, B1900FWA La Plata, Argentina</footnote>

<footnote n="4">⁴ Las Campanas Observatory, Carnegie Observatories, Casilla 601, La Serena, Chile</footnote>

<footnote n="5">⁵ Departamento de Astronomía, Universidad de La Serena, Cisternas 1200 Norte, La Serena, Chile</footnote>

<footnote n="6">⁶ Department of Astrophysics, IMAPP, Radboud University Nijmegen, PO Box 9010, 6500 GL Nijmegen, The Netherlands</footnote>

<footnote n="7">⁷ Institute of Astronomy, KU Leuven, Celestijnenlaan 200D, 3001 Leuven, Belgium</footnote>

<footnote n="8">⁸ Max-Planck-Institut für Astrophysik, Karl-Schwarzschild-Straße 1, 85741 Garching bei München, Germany</footnote>

<footnote n="9">⁹ International Gemini Observatory/NSF NOIRLab, 670 N. A'ohoku Place, Hilo, HI 96720, USA</footnote>

<footnote n="10">¹⁰ NRC Herzberg Astronomy and Astrophysics Research Centre, 5071 West Saanich Road, Victoria, BC V9E 2E7, Canada</footnote>

<footnote n="11">¹¹ Department of Astronomy, Space Science, and Geology, Chungnam National University, Daejeon 34134, Republic of Korea</footnote>

<footnote n="12">¹² Laboratório Nacional de Astrofísica, Rua Estados Unidos 154, 37504-364 Itajubá, MG, Brazil</footnote>

<footnote n="13">¹³ Institut d'Astrophysique de Paris, CNRS, UMR 7095, Sorbonne Université, 98 bis bd Arago, 75014 Paris, France</footnote>

<footnote n="14">¹⁴ NSF NOIRLab, 950 N. Cherry Ave., Tucson, AZ 85719, USA</footnote>

<footnote n="15">¹⁵ Department of Physics and Astronomy, University of Victoria, PO Box 3055, STN CSC, Victoria, BC V8W 3P6, Canada</footnote>

<footnote n="16">¹⁶ European Southern Observatory, Alonso de Cordova, 3107, Casilla, 19001 Santiago, Chile</footnote>

<footnote n="17">¹⁷ Australian Astronomical Optics, Macquarie University, 105 Delhi Rd, North Ryde, NSW 2113, Australia</footnote>

<footnote n="18">¹⁸ Sydney Institute for Astronomy, School of Physics, The University of Sydney, NSW 2006, Australia</footnote>

<footnote>Received 2024 August 9; revised 2025 January 8; accepted 2025 January 8; published 2025 February 12</footnote>

<footnote n="19">¹⁹ Member of the Carrera del Investigador Científico of the Comisión de Investigaciones Científicas of the Provincia de Buenos Aires (CICPBA), Argentina</footnote>

<footnote>Original content from this work may be used under the terms of the [Creative Commons Attribution 4.0 licence](https://creativecommons.org/licenses/by/4.0/). Any further distribution of this work must maintain attribution to the author(s) and the title of the work, journal citation and DOI.</footnote>
(H. Sana 2017; R. H. Barbá et al. 2017). Establishing the multiplicity properties of large, statistically significant samples of newly born massive stars is crucial to understanding the initial conditions and properly computing the evolution of these objects (N. Langer et al. 2008; S. E. de Mink et al. 2009).

Studies of binary systems are important because they help to understand the properties of massive stars in general, since by characterizing their components we can obtain their fundamental parameters, increasing the number of massive stars with known parameters. Moreover, determining their masses is crucial because the evolution of stars depends mainly on them. Therefore, it is necessary to obtain accurate dynamical masses, and a robust method to do it is through Kepler's laws.

With the aim of addressing the characterization of binary systems a survey started in 2006, called OWN Survey, which is a high-resolution spectroscopic survey that monitors a large sample of the O and Wolf Rayets with nitrogen lines (WN) stars of the Southern Hemisphere. To date there are data for more than 300 stars, of which 34 are double-lined spectroscopic binary (SB2) systems and 31 are single-lined spectroscopic binary (SB1) systems (R. Gamen et al. 2008; R. H. Barbá et al. 2017; R. Gamen et al. 2025, in preparation). Thus, there are a significant number of binary systems that are SB1s. Recently, there have been several works trying to characterize the unseen companion in SB1s (G. Banyard et al. 2022; L. Mahy et al. 2022; T. Shenar et al. 2022), but more systematic studies of these systems are needed to increase the sample of SB2s from which stellar parameters can be determined with precision.

In SB1 systems, where the luminosity of binary components significantly differs, extracting the spectral signatures of companions fainter than a certain fraction of the primary's brightness becomes challenging, particularly, in O-type binary systems, in which the luminosity of the O star dominates the whole flux. The rapid rotation and the low light contribution can also lead to a nondetection of the secondary component (T. Shenar et al. 2020; C. Putkuri et al. 2021). Consequently, there is a limit on the mass of (normal)<fnref n="20" /> companions that can be extracted through this disentangling process.

M. Moe & R. Di Stefano (2017) compiled observations of 44 SB2 systems in which at least the primary component is a massive star and showed that each observational technique (spectroscopy, photometry, etc.) is only able to detect companions within a certain range of orbital periods and mass ratios. In particular, spectroscopy can only detect close companions when the mass ratio is $q \gtrsim 0.25$.

The focus of this paper is the massive spectroscopic binary HD 165246 (CD−24°13880; $V=7.6$). This object was classified as O8 V(n) by N. R. Walborn (1972). S. A. Otero (2007) discovered its eclipsing binary nature using ASAS3 data and derived a period $P = 4.5928$ days. P. Mayer et al. (2013) added spectroscopy to their analysis and modeled the system with the PHOEBE code (A. Prša & T. Zwitter 2005). They found a slightly eccentric orbit with a period $P = 4.592706$ days and a very low mass ratio of $q = 0.17$. They determined, for the more luminous component, a mass $M \approx 22 M_\odot$, radius $R = 7.4 \pm 0.2 R_\odot$, and a high rotational velocity $v \sin i = 244 \pm 3$ km s⁻¹, making this component a rapidly rotating high-mass star. Although these authors found no direct indication of the secondary lines in the spectra of HD 165246, they proposed it to be an SB2 based on the results of a spectral disentangling

method. H. Sana et al. (2014) acquired interferometric observations of this target with Sparse Aperture interferometric Masks (SAM) in NACO instrument, and they found three companions in addition to the inner eclipsing binary. The closest of them is at $\rho = 30 \pm 16$ mas, the next at $\rho = 1.93 \pm 0.04$, and the third is at more than 5". This makes HD 165246 a sextuple source, where the Aa component is the binary system we are studying here. Due to these observations resolution, any companion closer than 30 mas cannot be detected. More recently, the system was studied by C. Johnston et al. (2017, 2021). They used photometry from the Kepler K2 mission to characterize the binary model and the intrinsic variability. They also obtained 160 spectra with the HERMES spectrograph (Mercator telescope, Santa Cruz de La Palma), to investigate the presence of line profile variability. They detected variability in both the K2 photometry and HERMES spectroscopy and found it consistent with rotational modulation as well as intrinsically excited stellar pulsations. This also explained (partially) the variable and high macroturbulence, which is linked to time-dependent pulsational line broadening.

In this work we revisit HD 165246 by leveraging new optical spectra obtained with the Gemini High-resolution Optical SpecTrograph (GHOST; J. Pazder et al. 2020), a new Gemini Observatory facility instrument with a high resolving power of $R \sim 80,000$. This instrument can provide top-notch, high-quality data, which will make possible unexpected discoveries.

The remainder of this paper is organized as follows: Observations are described in Sections 2 and 3, and Section 4 present the results of the spectroscopic and photometric analyses. We discuss the evolutionary status of the system in Section 5 and we summarize our conclusions in Section 6.

## 2. Observations

### 2.1. K2 Photometry

The Kepler space telescope was a mission launched by NASA in 2009 and its principal goal was to discover exoplanets. In 2012, Kepler's primary mission was completed and a new mission called K2 "second light" was presented, using its remaining capabilities. HD 165246 was observed by the K2 mission in both long-cadence (29.42 minute integration time) and short-cadence (58.85 s integration time) modes in 2016, from April 21 to July 1. As this star is a bright target ($V=7.6$; J. R. Ducati et al. 2001), it saturates the Kepler detectors, spreading the signal along several pixels. Moreover, as this target lies in a crowded field, multiple contaminating sources were included in the pipeline mask. To deal with this, C. Johnston et al. (2017) used a custom pixel mask avoiding the contaminating sources but including the saturated columns, and finally they retrieved a ~30 day lightcurve. The authors kindly shared their lightcurve to be used in this work.

### 2.2. Spectroscopic Data

We obtained spectroscopic data between 2006 and 2019 under the OWN Survey program. We used the FEROS echelle spectrograph (A. Kaufer et al. 1999), MIKE (R. Bernstein et al. 2003), the echelle spectrograph mounted on the 2.5 m Irénée du Pont telescope,<fnref n="21" /> and the REOSC spectrograph.<fnref n="22" />

<footnote n="20">²⁰ It is important to note that the rationale outlined above may not be directly applicable to compact objects such as neutron stars or white dwarfs.</footnote>

<footnote n="21">²¹ https://www.lco.cl/?epkb_post_type_1=echelle</footnote>

<footnote n="22">²² https://casleo.conicet.gov.ar/reosc-ds-dc/</footnote>
<table><caption>**Table 1**

Technical Characteristics of the Spectrographs Used for This Work + Max Planck Gesellschaft</caption><thead><tr><th>OBSERVATORY</th><th>TELESCOPE</th><th>SPECTROGRAPH</th><th>SPECTRAL RANGE (Å)</th><th>RESOLVING POWER</th><th>NUMBER OF SPECTRA</th></tr></thead><tbody><tr><td>GEMINI</td><td>GEMINI SOUTH (8.1 M)</td><td>GHOST</td><td>~3600–10500</td><td>80,000</td><td>9</td></tr><tr><td>CASLEO</td><td>JORGE SAHADE (2.15 M)</td><td>REOSC</td><td>~3500–6100</td><td>12,600</td><td>25</td></tr><tr><td>LAS CAMPANAS</td><td>IRÉNÉE DU PONT (2.5 M)</td><td>ECHELLE</td><td>~3600–10100</td><td>40,000</td><td>12</td></tr><tr><td>LA SILLA</td><td>MPG+/ESO (2.2 M)</td><td>FEROS</td><td>~3600–9200</td><td>48,000</td><td>12</td></tr><tr><td>LAS CAMPANAS</td><td>CLAY-MAGELLAN (6.5 M)</td><td>MIKE</td><td>~3200–5000 (blue)<br/>~4900–10000 (red)</td><td>28,000<br/>22,000</td><td>3<br/>3</td></tr></tbody></table>

We also acquired high-resolution spectra using the newly commissioned GHOST at Gemini South (A. W. McConnachie et al. 2024). The data were collected as part of the System Verification Program ID: GS-2023A-SV-104 during selected nights in 2023 May. The instrument was used in its single-object, high-resolution mode (binning spectral × spatial = 1 × 2). We obtained three exposures of 200 s for three epochs, which, after being combined, resulted in a signal-to-noise ratio (SNR) of ~400 at 4400 Å and ~500 at 5850 Å.

OWN Survey spectra were reduced using standard IRAF<fnref n="23" /> routines. GHOST data were processed using version 1.0.0 of the GHOSTDR (M. J. Ireland et al. 2018; C. R. Hayes et al. 2022) software within the DRAGONS 3.0 environment (K. Labrie et al. 2019, 2022). Table 1 summarizes the main properties of these instruments.

Since the companion named Ab by H. Sana et al. (2014) is located at 30 mas from the inner eclipsing binary, our spectra are expected to be contaminated by the light from this star. In fact, the slit widths used were ~3" (REOSC, CASLEO), 1" (du Pont's echelle, Las Campanas Observatory (LCO)), 1" (MIKE, LCO), and ~0.5 and 2.7 for the fiber integral field unit entrances of GHOST (Gemini) and FEROS (ESO La Silla), respectively. Nevertheless, it is important to note that there is no evidence that this component is gravitationally bound with the inner binary. In either way, we do not expect that this companion interferes with the radial velocity (RV) measurements or the orbital parameters of the system (we do not expect Kozai-Lidov oscillations).

### 3. Spectroscopic Analysis

#### 3.1. Radial Velocities and Spectral Classification

The composite spectrum of HD 165246 exhibits features consistent with an O8 V(n) type star, including intense broad lines of H, He I, and He II. Initially, our investigation focused on the OWN Survey spectra, aiming to uncover any subtle indications of a hidden companion, such as line asymmetries during quadrature phases, or some metallic spectral features originating in a cooler star. Our efforts yielded promising results, as we were able to discern the faint He I λ5876 line, both blue- and redshifted, in four spectra obtained during the quadrature phases. In addition, after obtaining the GHOST spectra, we were also able to identify He I λ4471 and Mg II λ4481. In Figure 1, we compare the aspect of the He I λ5876 line in a FEROS spectrum and a GHOST spectrum obtained at nearly identical orbital phases.

<figure><img src="image_3.png" /><figcaption>**Figure 1.** Observed composite spectra of GHOST (left panel) and FEROS (right panel) around the He I λ5876 line at phases close to quadrature. In the GHOST spectrum, the detection of the secondary component is evident, which was suspected on the FEROS one. In both panels, the green line shows the secondary component.</figcaption></figure>

Given that these lines arise in a very faint component, and the secondary's contribution to the composite spectrum is minimal, we chose to determine the RVs of the primary star by fitting Gaussian profiles to the He I λ4471, He I λ4922, He I λ5016, He II λ4542, He II λ5412, O III λ5592, and C IV λ5801 spectral lines. Measurements were performed using the IRAF task SPLOT. Since these lines did not exhibit significant RV differences among them, we calculated their average to represent the motion of the primary component.

Regarding the secondary star, only the He I λ5876 line could be measured in the OWN Survey spectra. In the GHOST data, as just explained, we also detected the He I λ4471 and Mg II λ4481 lines, and, since their RVs closely matched those of He I λ5876, we decided to use the average of the three lines in our analysis.

The RV measurements, along with the corresponding instruments used, are detailed in Table 5 in the Appendix. As we used five different instruments, we also measured the IS line Na I λ5890 (D₂) to assess the need for a zero-point correction. The mean velocities of the IS D₂ line for the different instruments: FEROS, du Pont echelle, REOSC, and GHOST are -51.6, -50.2, -51.0, and -50.9 km s⁻¹, with rms of 0.1, 0.7, 1.0, and 0.5 km s⁻¹, respectively. The rms serve as an estimate of the instrumental error in the RV measurements. We concluded that a zero-point correction was not necessary.

In Figure 2 we show two GHOST spectra obtained at opposite quadratures, focusing on the wavelength ranges encompassing the three most prominent lines of the secondary star. It can be seen here that the intensities of Mg II λ4481 and He I λ4471 are similar. Since the comparison between these lines is the key criterion for classifying late B stars (C. J. Evans et al. 2015), the detection of Mg II λ4481 and He I λ4471

<footnote n="23">²³ NOIRLab's IRAF is distributed by the Community Science and Data Center at NSF NOIRLab, which is managed by the Association of Universities for Research in Astronomy (AURA) under a cooperative agreement with the U.S. National Science Foundation.</footnote>
<figure><img src="image_4.png" /><figcaption>**Figure 2.** Observed composite GHOST spectrum of HD 165246 for phases $Φ = 0.86$ (upper panel) and $Φ = 0.24$ (bottom panel), near quadrature. In both spectra, the He I λ4471, Mg II λ4481, and He I λ5876 lines are visible for the primary (orange) and the secondary (green) components.</figcaption></figure>

suggests that this component has a spectral type of B8. Because no other spectral features are detected, it is impossible to determine its luminosity class.

## 3.2. Quantitative Analysis

The determination of line broadening requires careful consideration. C. Johnston et al. (2021) discuss the challenges in determining these parameters, particularly in the presence of pulsations. Using a methodology similar to ours, they employed the IACOB-BROAD tool to separate macroturbulent from rotation broadening. They analyzed 10 spectra taken at different times, finding a mean $v \sin i$ of 230 km s$^{-1}$ with a standard deviation of 13 km s$^{-1}$ using the goodness-of-fit (GOF) method, while their estimates of macroturbulence exceeded 100 km s$^{-1}$ in all cases, with a mean of 148 km s$^{-1}$ and a standard deviation of 41 km s$^{-1}$. Furthermore, when assuming zero macroturbulence, they determined $v \sin i = 268 \pm 25$ km s$^{-1}$. They concluded that the $v \sin i$ estimated using GOF with zero macroturbulence was the most reliable and conducted an atmospheric analysis using a genetic algorithm wrapped around the non-LTE radiative transfer code FASTWIND to optimize various atmospheric parameters, including macroturbulence, estimating it to be $20 \pm 6$ km s$^{-1}$. On the other hand, G. Holgado et al. (2022) estimated a $v \sin i$ of 251 km s$^{-1}$ without explicitly providing a macroturbulence value, assuming broadening parameters below 40 km s$^{-1}$ as zero. Our determination of $v \sin i$ is consistent with these previous works; however, the discrepancy lies in the macroturbulence velocity.

In this work, we estimated the broadening components using our higher-SNR composite GHOST spectrum, taken near quadrature when the secondary star's contribution to the primary spectral profiles is marginal, thereby allowing for

more accurate measurements. We analyzed the He I λλλ4471, 5876, and 5016 lines (as did C. Johnston et al. 2021) and the O III λ5592 metallic line (as G. Holgado et al. 2022). The results averaged across these lines yielded $v \sin i = 247 \pm 13$ km s$^{-1}$ and $v_{mac} = 74 \pm 13$ km s$^{-1}$. To further investigate potential variations in $v_{mac}$, we analyzed three spectra at different orbital phases. We did not detect significant changes, with $v_{mac} = 77 \pm 13$ km s$^{-1}$ ($v \sin i = 251 \pm 13$ km s$^{-1}$) at phase = 0.63, $v_{mac} = 73 \pm 14$ km s$^{-1}$ ($v \sin i = 263 \pm 14$ km s$^{-1}$) at phase = 0.22, and $v_{mac} = 85 \pm 12$ km s$^{-1}$ ($v \sin i = 244 \pm 12$ km s$^{-1}$) at phase = 0.05. These determinations were made using ESO and LCO spectra with an SNR greater than 250. The discrepancies reported in the literature, where $v_{mac}$ values exceed 100 km s$^{-1}$ or fall below 40 km s$^{-1}$, may be related to pulsations, but a reliable analysis is out of scope for this work.

We adopted the broadening parameters found with IACOB-BROAD tool: $v \sin i = 247 \pm 13$ km s$^{-1}$ and $v_{mac} = 74 \pm 13$ km s$^{-1}$, and fixed for IACOB-GBAT the associated helium abundances ($Y_{He} = 0.09$), the microturbulence velocities ($\xi_{1} = 5$ km s$^{-1}$), wind strength$^{24}$ ($Q = -13$), and the wind parameter $\beta = 0.8$; also, solar metallicity was assumed. We fitted the following lines: H$\alpha$, H$\beta$, H$\gamma$, H$\delta$, H$\epsilon$, He I λ4026, He I λ4471, He I λ4713, He I λ5015, He I λ5876, He II λ4200, He II λ4542, He II λ4686, and He II λ5412. It should be noted that the surface gravity derived from the wings of the hydrogen Balmer lines needs to be corrected for centrifugal effects due to rotation (A. Herrero et al. 1992; T. Repolst et al. 2004), and then the actual surface gravity becomes $log g = 3.95 \pm 0.09$. The derived physical parameters are described in Table 2 and the best model is presented in Figure 3.

<footnote>$^{24}$ The strength wind is defined by $Q = \frac{M}{(R v_{\infty})^{1.5}}$.</footnote>
<table><caption>**Table 2**
Parameters Obtained from the Quantitative Spectroscopic Analysis of the Primary Component of HD 165246</caption><thead><tr><th>Parameters</th><th>This Work</th></tr></thead><tbody><tr><td>v sin i [km s<sup>-1</sup>]</td><td>247 ± 13</td></tr><tr><td>v<sub>mac</sub> [km s<sup>-1</sup>]</td><td>74 ± 13</td></tr><tr><td>T<sub>eff</sub><sup>a</sup> [K]</td><td>34,000 ± 500</td></tr><tr><td>log g [dex]</td><td>3.95 ± 0.09</td></tr><tr><td>log Q [dex]</td><td>-13.0 (fixed)</td></tr><tr><td colspan="2">Fundamental parameters adopting M<sub>V</sub> = -4.41.<sup>b</sup></td></tr><tr><td>R [R<sub>☉</sub>]</td><td>8.2 ± 0.1</td></tr><tr><td>log (L/L<sub>☉</sub>) [dex]</td><td>4.96 ± 0.02</td></tr><tr><td>M<sub>sp</sub> [M<sub>☉</sub>]</td><td>18.3 ± 2.7</td></tr></tbody></table>

**Notes.**

<sup>a</sup> This value represents the hemisphere-averaged temperature.

<sup>b</sup> Adopted from F. Martins & B. Plez (2006).

# 4. Binary Model

## 4.1. Orbital Solution

We derived a spectroscopic orbit for the system using the RVs measured for both components (as detailed in Section 3.1). Initial parameters were adopted from previous studies (C. Johnston et al. 2017, 2021), and the analysis was conducted using the Fotometric Elements (FOTEL) program (P. Hadrava 2004), which resulted in a robust RV orbital solution. The code efficiently converged to the final solution, utilizing the RV data presented in the Appendix, with the resulting parameters summarized in Table 3 with the values found by C. Johnston et al. (2021) for comparison (see also Figure 4). It can be seen that all values agree within the errors with those obtained by C. Johnston et al. (2021), except for the eccentricity, which we will comment on below.

## 4.2. Combined Spectroscopic and Photometric Analysis

Following the spectroscopic results, we modeled the eclipsing binary system through the PHOEBE2 code (A. Prša et al. 2016). We assumed the gravity-darkening coefficients and the bolometric albedo to be $\alpha = 1$ and $\beta = 1$ for both components, respectively, since they have radiative envelopes (A. Claret 2001, 2003). The limb-darkening coefficients were manually set using W. van Hamme (1993) and A. Claret & S. Bloemen (2011) for bolometric and passband, respectively, in the square root prescription. Since we have contamination of the 30 mas component in the inner binary, we also considered a third source of light, whose value (18%) was adopted from C. Johnston et al. (2017).

The initial model was built using the results obtained with FOTEL (including the calculated RVs) and the results of IACOB-GBAT (see Sections 3.2 and 4.1). We use the Nelder-Mead optimizer method to find the best-fitting solutions for the RV and lightcurves individually. From this fit, we obtained an eccentricity of $e = 0.04 \pm 0.004$, which is more consistent with the one obtained by C. Johnston et al. (2017), while still showing an ~2σ difference. We also tried a fit with null eccentricity, which did not provide a good binary model when compared to the RVs.

In order to estimate the uncertainties, we run a Markov Chain Monte Carlo (MCMC) algorithm with the EMCEE implementation (D. Foreman-Mackey et al. 2013) in the

<figure><img src="image_5.png" /><figcaption>**Figure 3.** Comparison between the template found for the primary component (black) and the best FASTWIND model according to the parameters fitted with IACOB-GBAT. Note that IACOB-GBAT only uses H, He I, and He II lines for the fit.</figcaption></figure>

<table><caption>**Table 3**
Spectroscopic Solution Given by FOTEL</caption><thead><tr><th>Element</th><th>This Work</th><th>C. Johnston et al. (2021)</th></tr></thead><tbody><tr><td>P [days]</td><td>4.59276 ± 0.00006</td><td>4.59270 ± 0.00001</td></tr><tr><td>T<sub>0</sub><sup>a</sup> [HJD – 2400000]</td><td>57215.43 ± 0.02</td><td>57215.4 ± 0.06</td></tr><tr><td>T<sub>c</sub><sup>b</sup> [HJD – 2400000]</td><td>57215.46 ± 0.03</td><td>...</td></tr><tr><td>e</td><td>0.077 ± 0.002</td><td>0.030 ± 0.003</td></tr><tr><td>V<sub>γ</sub> [km s<sup>-1</sup>]</td><td>-7.9 ± 1</td><td>-7.9 ± 0.1</td></tr><tr><td>ω [°]</td><td>87.6 ± 1.2</td><td>93.4 ± 5.8</td></tr><tr><td>K<sub>a</sub> [km s<sup>-1</sup>]</td><td>52.7 ± 3.3</td><td>53.0 ± 0.2</td></tr><tr><td>K<sub>b</sub> [km s<sup>-1</sup>]</td><td>329.4 ± 36</td><td>...</td></tr><tr><td>q = M<sub>b</sub>/M<sub>a</sub></td><td>0.17 ± 0.01</td><td>0.16 ± 0.02</td></tr></tbody></table>

**Notes.** We consider the initial phase ($\Phi = 0$) at the primary eclipse. The given errors represent the inner errors of the numerical method used to fit.

<sup>a</sup> Time of periastron passage.

<sup>b</sup> Time of primary conjunction.

PHOEBE environment. To reduce the parameter space, we fixed the period and the epoch of the superior conjunction. We also consider the rotational velocities for the components, but, as we do not have an estimation for the secondary, we adopted the primary value. We did several trials with different values for the secondary and noted that this parameter has not observable effects on the binary model.
<figure><img src="image_6.png" /><figcaption>**Figure 4.** Left panel: RV curve showing RV measurements (blue for the primary and green for the secondary), along with the model fitted with PHOEBE2 (solid line for the primary, dashed lines for the secondary). Orange and red points represent the values of RV obtained from GHOST data for the primary and secondary, respectively. Green points are measurements of the secondary from FEROS data. Note that there is a value with residuals ~100 km s⁻¹, which was assigned a lower weight in the fitting. Blue points represent RV of the primary component obtained from FEROS, MIKE, REOSC, and the echelle of Irénée du Pont. The residuals of the fit are illustrated following the same color code, in units of m s⁻¹. Right panel: the lightcurve of the system where black dots represent the K2 data and the red line represents the fitted model calculated with PHOEBE2. The residuals to the fit are also shown. Note that phase Φ = 0 corresponds with the primary eclipse.</figcaption></figure>

<table><caption>**Table 4**
Binary Model Parameters Calculated with PHOEBE2</caption><thead><tr><th>Element</th><th>This Work</th><th>C. Johnston et al. (2021)</th></tr></thead><tbody><tr><td>q<sub>binary</sub></td><td>0.157<sup>+0.005</sup><sub>-0.005</sub></td><td>0.16 ± 0.02</td></tr><tr><td>a<sub>binary</sub> [R<sub>☉</sub>]</td><td>34.2<sup>+0.5</sup><sub>-0.5</sub></td><td>35.3<sup>+0.6</sup><sub>-0.7</sub></td></tr><tr><td>i<sub>binary</sub> [°]</td><td>84.5<sup>+0.2</sup><sub>-0.2</sub></td><td>84 ± 1</td></tr><tr><td>T<sub>eff,b</sub>/T<sub>eff,a</sub></td><td>0.4077<sup>+0.0028</sup><sub>-0.0028</sub></td><td>0.3559 ± 0.1393</td></tr><tr><td>R<sub>equiv,b</sub> [R<sub>☉</sub>]</td><td>2.22<sup>+0.04</sup><sub>-0.04</sub></td><td>2.4<sup>+0.3</sup><sub>-0.4</sub></td></tr><tr><td>R<sub>equiv,a</sub> [R<sub>☉</sub>]</td><td>7.0<sup>+0.2</sup><sub>-0.2</sub></td><td>7.3<sup>+0.3</sup><sub>-0.4</sub></td></tr><tr><td>T<sub>eff,a</sub> [K]</td><td>34,000.0 (fixed)</td><td>36,100 ± 200</td></tr><tr><td>T<sub>eff,b</sub> [K]</td><td>13, 861<sup>+102</sup><sub>-102</sub></td><td>12,100 ± 600</td></tr><tr><td>logg<sub>a</sub></td><td>4.09<sup>+0.01</sup><sub>-0.01</sub></td><td>4.08<sup>+0.02</sup><sub>-0.04</sub></td></tr><tr><td>logg<sub>b</sub></td><td>4.28<sup>+0.02</sup><sub>-0.02</sub></td><td>4.26<sup>+0.06</sup><sub>-0.14</sub></td></tr><tr><td>M<sub>a</sub> [M<sub>☉</sub>]</td><td>21.9<sup>+1.1</sup><sub>-1.1</sub></td><td>23.7<sup>+1.1</sup><sub>-1.4</sub></td></tr><tr><td>M<sub>b</sub> [M<sub>☉</sub>]</td><td>3.5<sup>+0.2</sup><sub>-0.2</sub></td><td>3.8<sup>+0.4</sup><sub>-0.5</sub></td></tr></tbody></table>

Note. Errors represent the MCMC fit and are estimated at 2σ.

We initialized the EMCEE sampler with a distribution of tiny Gaussian balls around $q, a, i, (R_a + R_b)/a, R_b/R_a, T_{eff,b}/T_{eff,a}$, and $L_{pb}$. As we found problems with convergence, we added a noise nuisance parameter $\sigma_{inf}$ associated with an underestimated error in the K2 fluxes. We used 48 walkers and had them take 2000 steps, after which it was evident that all the parameters had converged. We used the built-in PHOEBE plot of the autocorrelation function to quantify convergence. Then, we removed the first 1000 iterations as a burn-in phase and let them take another 5000 steps to ensure convergence. The final parameters with their uncertainties, taken as $2\sigma$ from the distributions, are listed in Table 4 with the results from C. Johnston et al. (2021) for comparison. We noticed a difference between the two results in the effective temperature of the primary. Because we used only single-color K2 observations and PHOEBE fits the ratio $T_{eff,b}/T_{eff,a}$, we fixed $T_{eff,a}$ to the one found in the quantitative analysis. This also leads to a difference in $T_{eff,b}$.

<figure><img src="image_7.png" /><figcaption>**Figure 5.** Evolutionary tracks for solar-composition, rotating stellar models during the main sequence in the surface gravity vs. effective temperature plane. The initial equatorial velocity of all objects is 250 km s⁻¹. Solid gray lines represent the tracks, labeled at the bottom with their respective masses on the ZAMS given in solar units. Tracks that best represent the components of the pair are in red, others are in black. Isochrones corresponding to ages with steps of 1 Myr (with zero for the ZAMS) are also shown. Isochrones for 0, 5, 10, 15, and 20 Myr are represented with solid lines, while intermediate ones are represented with short dashed lines. The one that corresponds to 2 Myr is in green, while the rest are in black. The data corresponding to HD 165246 are represented with black boxes. Note that the primary is compatible with an age of ≈2 Myr while the secondary falls on higher age isochrones.</figcaption></figure>

The final RV and lightcurves are shown in Figure 4. The dispersion in RV measurements with the fitted model observed in this figure is most likely related to the line profile variability characterized by C. Johnston et al. (2021).

## 5. Evolutionary Status of the System

To investigate the evolutionary status of HD 165246 we have created a set of models with our code, described in O. G. Benvenuto & M. A. De Vito (2003) and
<figure><img src="image_8.png" /><figcaption>**Figure 6.** Same as Figure 5 but also considering ages and rotation during pre-main-sequence evolution. We have set initial rotation velocities so that, after initial contraction, stars reach a maximum velocity of 250 km s⁻¹ near the ZAMS. Pre-main-sequence portion of each track is denoted with a blue line, while the main-sequence one is shown with a black one. For clarity, here we have included only the isochrone corresponding to 2 Myr. It corresponds to the main sequence for M ≥ 5.50 M☉, and to the pre-main sequence for less massive objects (denoted with solid green and red lines, respectively). Note that both components of the pair are better fitted in this case.</figcaption></figure>

C. Putkuri et al. (2022). In this code, the equation of state is solved by considering a Saha law for the main elements present in the mixture. Nuclear reaction rates have been taken from G. R. Caughlan & W. A. Fowler (1988), and radiative Rosseland opacities from C. A. Iglesias & F. J. Rogers (1996). Rotation has been included following the shell-based approximation (J. P. Zahn 1992). Basically, we solve for the rotation equations employing another Henyey scheme, which is coupled to the structure of the star considering a modified hydrostatic structure (G. Meynet & A. Maeder 1997). Thus, structure and rotation are found self-consistently. The changes in the chemical composition are computed by solving Equation (3.11) of A. Maeder & J.-P. Zahn (1998) assuming noninstantaneous mixing by convection, semiconvection, and meridional circulation. The mass-loss rate is from J. S. Vink et al. (2001), allowing its enhancement due to rotation (N. Langer 1998). As a first step in our analysis, we constructed a set of rotating, solar-composition (X = 0.7289, Z = 0.014) stellar models on the main sequence for masses between 2.92 M☉ and 26.25 M☉ logarithmically evenly spaced. We imposed a rotation velocity of 250 km s⁻¹ (adopting the value found for the primary) for all models on the ZAMS. The mixing length parameter αMLT has been set to 1.76 by fitting to present Sun conditions. Also, moderate overshooting has been considered (αOV = 0.2). Let us remark

that the value of α<sub>OV</sub> is rather uncertain. For example, S. Martinet et al. (2021) have constructed rapidly rotating stellar models considering different values of α<sub>OV</sub> ranging from 0.1 to 0.6. A solid determination of α<sub>OV</sub> remains an open problem. Since the pair is well detached, we have considered the stars isolated in computing these models, neglecting the effects due to the companion's presence.

To apply the models in analyzing the observational data, we employ a log g versus log T<sub>eff</sub> diagram. These data come directly from spectroscopic observations and are not affected by uncertainties in the distance to the pair or IS absorption, which are hard to measure accurately. The results are presented in Figures 5–6.

In Figure 5 we present the theoretical log g versus log T<sub>eff</sub> plane together with observations, setting zero ages on the ZAMS. From the analysis of this figure, we deduce that the primary star is compatible with an age of ≈2 Myr while the rather wide error box representing the secondary covers isochrones of higher age values. At the ZAMS, the primary had ≈24.0 M☉ and suffered a moderate mass loss, while the companion evolved with a constant mass of ≈3.7 M☉.

Due to the disparity of the masses of the components in HD 165246, it is interesting to generalize our previous analysis employing tracks starting from the pre-main sequence. For this purpose, we started the sequences on the corresponding
<figure><img src="image_9.png" /><figcaption>**Figure 7.** Main characteristics of the stars affecting the orbital evolution while the primary is on the main sequence. Upper panel: the masses of the components. Middle panel: radii of the components. Bottom panel: radius of gyration. The green dashed region represents the observed value with its uncertainty.</figcaption></figure>

Hayashi line including a given initial equatorial rotational velocity. After contracting toward the ZAMS, we have chosen this velocity, so that the stars reach a maximum rotational velocity of 250 km s⁻¹. For this case,<fnref n="25" /> the results are given in Figure 6. We depicted the isochrone of 2 Myr, and it can be seen that in the region of higher masses, the shape of the isochrone is very similar to that shown in Figure 5. However, for lower masses, there are qualitative differences. Now, both components of the pair are nicely fit by the isochrone of 2 Myr. Therefore, in this case, it is important to consider the pre-main-sequence phase to obtain consistent results. Remarkably, while the primary is on the main sequence, the companion seems to be on the pre-main sequence, still evolving toward the ZAMS.

It is important to note that the estimations of age of the cluster NGC 6530 (to which this system belongs) have a wide range, 1–5 Myr as seen in N. J. Mayne et al. (2007) and R. D. Jeffries et al. (2023), and references therein. It is possible that this cluster has undergone sequential formation, where the central stars have an age of ~2 Myr, and the ones further from the center, ~5 Myr.

<figure><img src="image_10.png" /><figcaption>**Figure 8.** Upper panel: evolution of the orbital period and the rotational period of each component. The magenta line represents the orbital period, whereas the blue and red lines correspond to the rotational periods of the primary and secondary stars, respectively. The green line depicts the observed value of the orbital period. Bottom panel: evolution of the eccentricity. The green dashed region represents the observed value with its uncertainty. In both panels, long dashed vertical lines indicate the estimated age of the pair.</figcaption></figure>

Now we can study the evolution of the system's orbit. For this purpose, we apply the formalism presented in S. Repetto & G. Nelemans (2014). To do so, for compatibility with the hypotheses on which the formalism is based, we constructed nonrotating models with masses of 24.0 $M_\odot$ and 3.75 $M_\odot$, and then we added rigid body rotation.<fnref n="26" /> The main quantities that affect the tidal evolution of the orbit, i.e., masses, radii, and the radius of gyration *k* (defined as $k^2 = I/\text{MR}^2$ where *I* is the moment of inertia of the star) are presented in Figure 7 as a function of time during $6 \times 10^6$ yr, which represents almost the entire main-sequence evolution of the primary. We have fitted these quantities with polynomial expressions and solved for the evolution of the main parameters of the orbit (i.e., orbital period, eccentricity, the velocity of rotation, and the inclination of the rotational axis to the plane of the orbit of each component) assuming a set of initial conditions.

After exploring the parameter space of the problem, we found that the initial inclination of the axes and the rotation rate of the secondary star do not affect noticeably the orbital evolution. Therefore, we set the rotation rate of the secondary to be the same fraction of its critical rate as that assumed for the primary and vary the initial values of the orbital period, eccentricity, and rotation rate of the primary. The best model solution (considered this way since it accounts for the

<footnote n="25">²⁵ Here we have considered models evolving on the pre-main sequence with constant mass. We preferred to do it this way for simplicity, without considering the treatment of the birth line, as it has been done in L. Haemmerlé et al. (2019). This does not affect the results presented here since, as shown in Figure 7 of that paper (see also their discussion in Section 4.1), evolutionary tracks resulting from both treatments are indistinguishable from each other near the ZAMS. Moreover, initial evolution is so fast that the isochrone of 2 Myr (of key relevance for our analysis) is unaffected by considering constant mass models or more realistic ones starting from a birth line.</footnote>

<footnote n="26">²⁶ See this section's final paragraph for a discussion of the consistency of this procedure.</footnote>
<figure><img src="image_11.png" /><figcaption>**Figure 9.** Upper panel: evolution of the equatorial velocity of the components. The blue and red lines correspond to the primary and secondary stars, respectively. The green dashed region represents the observed value with its uncertainty. Bottom panel: fractional change of the orbital period as a function of time. As in Figure 8, long dashed vertical lines indicate the estimated age of the pair.</figcaption></figure>

observations at the present age of the system) is presented in Figures 8 and 9. Its main characteristic is that for the primary to have a present rotational speed so high, it must have been even faster on the ZAMS. The orbit should have been almost circular and the period slightly shorter than now. As it can be seen in the upper panel of Figure 8, the primary needs almost its entire life on the main sequence to get approximately synchronized ($P_{\text{orbit}} \approx P_{\text{prim}}$). On the other hand, because of its weak tidal coupling due to its small radius, the secondary remains rotating at an almost constant rate. At these ages, the system tends to complete circularization but it does not reach such a condition as it can be seen in the lower panel of Figure 8.

For the sake of completeness, in the upper panel of Figure 9 we show the evolution of the components' rotational velocities. The primary transfers angular momentum to the orbit and slows down appreciably, while the secondary remains almost unaffected as quoted above. Finally, in the lower panel of Figure 9 we show the expected fractional change of the orbital period, far from the present observational capabilities.

In closing this section, a general comment is in order. The treatment we have performed first for studying the evolution of the stars of the system and then to analyze the orbital evolution is not fully consistent. In the formalism presented in S. Repetto & G. Nelemans (2014), it is assumed that stars rotate as rigid bodies. However, for our stellar interior models, following the treatment presented by J. P. Zahn (1992), we allowed a

nonconstant profile for the internal angular velocity. Constructing completely self-consistent models for this kind of binary system represents a very difficult task, which is far beyond the scope of this paper. In any case, we remark that the lack of consistency should not affect the main conclusions of this section.

## 6. Summary and Outlook

In this paper, we have studied the massive binary system HD 165246. We successfully detected spectral lines of the secondary component for the first time, initially, in a subset of OWN Survey data and, subsequently, in GHOST spectra collected during a System Verification observing run. The key lines identified include He I λ5876, Mg II λ4481, and He I λ4471 (only in GHOST data for the last two). The relative intensities of the latter lines serve as the primary classification criteria for late B-type stars, indicating a spectral type of B8 for the secondary star. Their RVs were measured and combined with the primary's velocities to calculate an SB2 orbital solution. An extremely low mass ratio of $q = 0.157 \pm 0.002$ was derived for the system.

Taking into account the eclipsing binary nature of HD 165246, as revealed in the K2 photometry, we employed a PHOEBE binary model to derive astrophysical parameters for each star. Our analysis yielded the following values: masses of $M_a = 22 M_\odot$ and $M_b = 3.4 M_\odot$, radii of $R_a = 7.0 R_\odot$ and $R_b = 2.2 R_\odot$, and effective temperatures of $T_{\text{eff,a}} = 34,000$ K and $T_{\text{eff,b}} = 13,900$ K for the primary and secondary components, respectively.

We modeled the evolution of each component, considering rotation. The tracks are presented in a log g versus log T<sub>eff</sub> diagram (see Figures 5–6). The positions of the components are consistent with tracks of ZAMS masses of ≈24 $M_\odot$ and ≈3.7 $M_\odot$ for the primary and secondary, respectively. Considering the time elapsed since their formation, both stars closely match an isochrone of 2 Myr (see Figure 6). Notably, while the primary is on the main sequence, the secondary appears to be evolving toward the ZAMS. Thus, the secondary should be considered as a pre-main-sequence object.

We also performed modeling of the tidal evolution of the system (see Figures 7–9), from which we infer that the primary component is not yet synchronized with the orbit and the secondary will not be able to reach synchronization in the lifetime of the detached system. The modeled parameters coincide with the observation at the considered age of the system (2 Myr). Furthermore, the equatorial velocity of the primary is expected to fall considerably during its lifetime on the main sequence, while the secondary will not show any changes.

The analysis presented in this paper indicates that HD 165246 has a very low mass ratio, making it the lowest known to date. In the log($P_{\text{orb}}$) versus q plane it falls below the region of detection corresponding to SB2s (see Figure 1 in M. Moe & R. Di Stefano 2017). In Figure 10, we show this plane for massive systems with data available to date and mass ratio q calculated from spectroscopy. It can be seen that HD 165246 lies in a region in which it was expected to be very hard to detect pairs, as proposed by M. Moe & R. Di Stefano (2017). So, this paper represents a contribution to populate this relevant region of the log($P_{\text{orb}}$) versus q plane.
<figure><img src="image_12.png" /><figcaption>**Figure 10.** log(*P*) vs. *q* plane for SB2 systems known to date where the primary masses are greater than 9 *M*⊙ and *q* is obtained from a spectroscopic analysis. In different colors are the data taken from H. Levato et al. (1987), H. A. Abt et al. (1990), H. Sana et al. (2012), H. A. Kobulnicky et al. (2014), G. Banyard et al. (2022), and J. Li et al. (2024). The black star represents the case of HD 165246.</figcaption></figure>

The primary star of this system will fill its Roche lobe and start a mass-transfer episode in a few million years. A short time after, it is expected to undergo a common envelope episode (B. Paczynski 1976). The extreme mass ratio indicates that the thermal timescales of the components are very different. Therefore, the most massive star of the pair will begin to transfer mass to the companion at a high rate. The companion will be unable to accommodate quickly enough the received matter. Thus, the companion will swell considerably up to fill its own Roche lobe. Then, the pair enters a common envelope stage. It is difficult to envisage its final fate since the common envelope phenomenon is yet to be completely understood (see P. Marchant & J. Bodensteiner 2023, and references therein).

With the arrival of new high-resolution spectrographs (such as GHOST) some systems now known as SB1 will become SB2, as has been the case with HD 165246. This will allow us to populate zones of the log(*P*<sub>orb</sub>) versus *q* plane with little or no data yet. This will be very helpful to constrain theories of massive binary star formation.

## Acknowledgments

C.N.R. acknowledges support from Consejo Nacional de Investigaciones Científicas y Técnicas (CONICET) through the Beca Interna Doctoral grant. C.N.R., G.A.F., T.A., R.G., R.E. H., and C.P. acknowledge support from grant PICT 2019-0344. The work of C.N.R., K.C., E.D., J.E.H., V.K., V.M.P., S.X., R.D., M.G.J., R.R.C., C.S., K.L., and J.T.O. is supported by NOIRLab, which is managed by the Association of Universities for Research in Astronomy (AURA) under a cooperative agreement with the U.S. National Science Foundation. Based on observations obtained under Program ID GS-2023A-SV-104, at the International Gemini Observatory, a program of NSF NOIRLab, which is managed by the Association of Universities for Research in Astronomy (AURA) under a cooperative agreement with the U.S. National Science Foundation on behalf of the Gemini Observatory partnership: the U.S. National Science Foundation (United States), National Research Council (Canada), Agencia Nacional de Investigación y Desarrollo (Chile), Ministerio de

Ciencia, Tecnología e Innovación (Argentina), Ministério da Ciência, Tecnologia, Inovações e Comunicações (Brazil), and Korea Astronomy and Space Science Institute (Republic of Korea).

GHOST was built by a collaboration between Australian Astronomical Optics at Macquarie University, National Research Council Herzberg of Canada, and the Australian National University, and funded by the International Gemini partnership. The instrument scientist is Dr. Alan McConnachie at NRC, and the instrument team is also led by Dr. Gordon Robertson (at AAO) and Dr. Michael Ireland (at ANU). The authors would like to acknowledge the contributions of the GHOST instrument build team, the Gemini GHOST instrument team, the full SV team, and the rest of the Gemini operations team that were involved in making the SV observations a success. Data processed using Data Reduction for Astronomy from Gemini Observatory North and South (DRAGONS). C.J. gratefully acknowledges support from the Netherlands Research School of Astronomy (NOVA). E.M. acknowledges funding from FAPEMIG under project number APQ-02493-22 and a research productivity grant number 309829/2022-4 awarded by the CNPq, Brazil. E.M. acknowledges funding from FAPEMIG under project number APQ-02493-22 and a research productivity grant number 309829/2022-4 awarded by the CNPq, Brazil. K.A.V. thanks the National Sciences and Engineering Research Council of Canada for funding through the Discovery Grants and CREATE programs.

**Facilities:** Gemini:South, CASLEO:JST, Du Pont, Magellan:Clay, and Max Planck:2.2m.

**Software:** NOIRLab's IRAF (D. Tody 1986, 1993; M. Fitzpatrick et al. 2024) and MATPLOTLIB (J. D. Hunter 2007).

## Appendix

### Radial Velocities Measured

In Table 5 we show the RVs measured for both components of HD 165246, which are indicated with the subscripts “a” and “b.” In successive columns we tabulate the heliocentric Julian day (HJD) and the RVs measured for each component (see main text). To estimate the error of an individual RV measurement, we repeat many times the Gaussian fitting of a
<table><caption>**Table 5**

Radial Velocity Measurements for the Primary Component (RV<sub>a</sub>) and Secondary (RV<sub>b</sub>)</caption><thead><tr><th>HJD - 2400000 days</th><th>RV<sub>a</sub> (km s<sup>-1</sup>)</th><th>RV<sub>b</sub> (km s<sup>-1</sup>)</th><th>Instrument</th></tr></thead><tbody><tr><td>53868.7834</td><td>-16.8</td><td></td><td>REOSC</td></tr><tr><td>53938.7759</td><td>-3.9</td><td></td><td>du Pont</td></tr><tr><td>53965.5722</td><td>-35.7</td><td></td><td>FEROS</td></tr><tr><td>53966.5548</td><td>21.0</td><td></td><td>FEROS</td></tr><tr><td>53967.5295</td><td>43.5</td><td>-202.0</td><td>FEROS</td></tr><tr><td>53968.5284</td><td>-14.0</td><td></td><td>FEROS</td></tr><tr><td>53969.5321</td><td>-57.6</td><td>278.5</td><td>FEROS</td></tr><tr><td>53970.5551</td><td>-1.6</td><td></td><td>FEROS</td></tr><tr><td>53954.6740</td><td>-31.9</td><td>136.1</td><td>MIKE</td></tr><tr><td>53955.6908</td><td>-48.0</td><td></td><td>MIKE</td></tr><tr><td>53986.5697</td><td>10.0</td><td></td><td>FEROS</td></tr><tr><td>53987.5474</td><td>-32.5</td><td></td><td>FEROS</td></tr><tr><td>53988.5404</td><td>-33.4</td><td></td><td>FEROS</td></tr><tr><td>53989.5429</td><td>25.5</td><td></td><td>FEROS</td></tr><tr><td>53991.5540</td><td>-18.6</td><td></td><td>FEROS</td></tr><tr><td>54198.9193</td><td>-55.0</td><td></td><td>du Pont</td></tr><tr><td>54199.9107</td><td>-38.4</td><td></td><td>du Pont</td></tr><tr><td>54209.9196</td><td>33.37</td><td></td><td>FEROS</td></tr><tr><td>54211.9021</td><td>-18.4</td><td></td><td>FEROS</td></tr><tr><td>54219.8519</td><td>49.1</td><td>-328.9</td><td>REOSC</td></tr><tr><td>54220.9207</td><td>18.3</td><td></td><td>REOSC</td></tr><tr><td>54246.8379</td><td>42.0</td><td></td><td>FEROS</td></tr><tr><td>54251.8714</td><td>60.3</td><td></td><td>REOSC</td></tr><tr><td>54252.8665</td><td>26.5</td><td></td><td>REOSC</td></tr><tr><td>54310.7503</td><td>10.0</td><td></td><td>REOSC</td></tr><tr><td>54316.6298</td><td>45.0</td><td></td><td>REOSC</td></tr><tr><td>54338.6585</td><td>5.4</td><td></td><td>REOSC</td></tr><tr><td>54339.6447</td><td>48.4</td><td></td><td>REOSC</td></tr><tr><td>54340.6446</td><td>-15.4</td><td></td><td>REOSC</td></tr><tr><td>54582.8803</td><td>82.8</td><td></td><td>REOSC</td></tr><tr><td>54608.9178</td><td>-29.7</td><td></td><td>REOSC</td></tr><tr><td>54644.8198</td><td>-55.6</td><td></td><td>REOSC</td></tr><tr><td>54626.7012</td><td>-36.9</td><td></td><td>REOSC</td></tr><tr><td>54660.7353</td><td>40.0</td><td></td><td>REOSC</td></tr><tr><td>54693.6769</td><td>3.7</td><td></td><td>REOSC</td></tr><tr><td>54955.9314</td><td>-15.0</td><td></td><td>FEROS</td></tr><tr><td>54956.9286</td><td>-59.5</td><td></td><td>FEROS</td></tr><tr><td>55761.7646</td><td>10.1</td><td></td><td>MIKE</td></tr><tr><td>55846.5352</td><td>5.8</td><td></td><td>du Pont</td></tr><tr><td>56438.8904</td><td>18.2</td><td></td><td>du Pont</td></tr><tr><td>58408.5269</td><td>79.1</td><td></td><td>du Pont</td></tr><tr><td>58409.5687</td><td>-14.3</td><td></td><td>du Pont</td></tr><tr><td>58596.8129</td><td>60.5</td><td></td><td>REOSC</td></tr><tr><td>58596.8511</td><td>45.0</td><td></td><td>REOSC</td></tr><tr><td>58597.7615</td><td>34.5</td><td></td><td>REOSC</td></tr><tr><td>58597.7860</td><td>38.6</td><td></td><td>REOSC</td></tr><tr><td>58598.8536</td><td>-38.3</td><td></td><td>REOSC</td></tr><tr><td>58598.8777</td><td>-38.8</td><td></td><td>REOSC</td></tr><tr><td>58599.7602</td><td>-38.5</td><td></td><td>REOSC</td></tr><tr><td>58599.7862</td><td>-17.3</td><td></td><td>REOSC</td></tr><tr><td>58600.7708</td><td>40.3</td><td></td><td>REOSC</td></tr><tr><td>58600.7955</td><td>42.1</td><td></td><td>REOSC</td></tr><tr><td>60077.8113</td><td>-25.9</td><td>316.1</td><td>CHOST</td></tr><tr><td>60079.7428</td><td>55.9</td><td>-237.1</td><td>CHOST</td></tr><tr><td>60080.7370</td><td>45.4</td><td>-212.1</td><td>CHOST</td></tr></tbody></table>

Note. The typical error of a measurement is 2.5 km s<sup>-1</sup> for the primary and 5 km s<sup>-1</sup> for the secondary.

line in a particular spectrum and calculate the standard deviation of the central wavelengths obtained. Then we repeat this procedure in several lines from different spectra. In this

way, we estimate a typical error of a particular measurement as 2.5 km s<sup>-1</sup> for the primary and 5 km s<sup>-1</sup> for the secondary.

## ORCID iDs

C. N. Rodriguez https://orcid.org/0000-0002-8744-803X

T. Ansin https://orcid.org/0009-0004-3544-1751

G. A. Ferrero https://orcid.org/0000-0001-9623-6218

O. G. Benvenuto https://orcid.org/0009-0005-0973-470X

R. Gamen https://orcid.org/0000-0002-5227-9627

N. I. Morrell https://orcid.org/0000-0003-2535-3091

J. I. Arias https://orcid.org/0000-0001-7500-7352

C. Johnston https://orcid.org/0000-0002-3054-4135

K. Chiboucas https://orcid.org/0000-0002-9020-5004

E. Deibert https://orcid.org/0000-0001-9796-2158

J.-E. Heo https://orcid.org/0000-0003-2530-3000

M. Jeong https://orcid.org/0009-0009-7838-7771

V. Kalari https://orcid.org/0000-0002-4641-2532

E. Martioli https://orcid.org/0000-0002-5084-168X

V. M. Placco https://orcid.org/0000-0003-4479-1265

K. A. Venn https://orcid.org/0000-0003-4134-2042

S. Xu https://orcid.org/0000-0002-8808-4282

R. Diaz https://orcid.org/0000-0001-9716-5335

R. Ruiz-Carmona https://orcid.org/0000-0001-7518-1393

C. Simpson https://orcid.org/0000-0001-8589-4055

A. W. McConnachie https://orcid.org/0000-0003-4666-6564

G. Burley https://orcid.org/0000-0001-7548-5354

T. A. M. Berg https://orcid.org/0000-0002-2606-5078

J. G. Robertson https://orcid.org/0000-0001-5528-7801

K. Labrie https://orcid.org/0000-0002-6633-7891

J. Thomas-Osip https://orcid.org/0000-0003-1033-4402

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