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A&A
Volume 691, November 2024
Article Number A214
Number of page(s) 7
Section Stellar structure and evolution
DOI https://doi.org/10.1051/0004-6361/202451444
Published online 15 November 2024

© The Authors 2024

Licence Creative CommonsOpen Access article, published by EDP Sciences, under the terms of the Creative Commons Attribution License (https://creativecommons.org/licenses/by/4.0), which permits unrestricted use, distribution, and reproduction in any medium, provided the original work is properly cited.

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1. Introduction

On August 17th, 2017, the LIGO-Virgo Collaboration made the first observation of a gravitational-wave (GW) signal from a binary neutron star (BNS) merger (GW170817, Abbott et al. 2017). The source has component masses between 1.17 and 1.60 M⊙ with a total mass of 2 . 74 − 0.01 + 0.04 M ⊙ $ 2.74^{+0.04}_{-0.01}\,M_\odot $ considering the low-spin priors, in agreement with the masses of the known Galactic BNS systems. The second GW BNS merger (GW190425), which was detected in the third observing run (O3) of the LIGO-Virgo-KAGRA network (Abbott et al. 2020) was found to have a total mass of 3 . 4 − 0.1 + 0.3 M ⊙ $ 3.4^{+0.3}_{-0.1}\, M_\odot $, lying 5σ deviations away from the mean mass of Galactic BNS systems (Farrow et al. 2019; Zhang et al. 2019). To date, there is no confirmed detection of electromagnetic counterparts partially due to its large distance and poor sky localization (Coughlin et al. 2019; Hosseinzadeh et al. 2019). Furthermore, detailed investigations in Zhao et al. (2023) showed that heavier BNS systems are expected to generate kilonovae with lower peak luminosities and more rapid decay when compared to BNS systems with smaller total mass. Given the incredibly heavy inferred mass, our current understanding of the origin of GW190425 is still controversial.

The dominant formation channel for BNS systems is considered to be isolated binary evolution (for example, Tauris et al. 2017). In this scenario, the primary star, which is initially more massive, evolves faster to become a NS after the first supernova (SN) explosion. The secondary star expands significantly after the main-sequence phase and then initiates mass transfer (MT) via the first Lagrangian point (L1) onto the companion NS star. During this phase, the system is more likely to undergo a dynamically unstable MT phase, which probably leads to the formation of a common envelope (so-called CE phase, Paczynski et al. 1976). If the system survives the CE phase, it becomes the immediate progenitor of a BNS, consisting of a He-rich star and a NS in a close orbit. Subsequently, the close binary system could undergo different MT phases, for example, Case BB/BC MT1 (Habets 1986; Tauris et al. 2015; Jiang et al. 2021), depending on the evolutionary phase of the He-rich star. The survival of the second SN explosion can lead to the formation of the BNS system, which could become prime search targets for ground-based GW detectors, such as LIGO (LIGO Scientific Collaboration 2015), Virgo (Acernese et al. 2015), and KAGRA (Aso et al. 2013).

In the standard isolated binary evolution channel, a fast-merging channel (namely unstable Case BB MT channel) is found to be inefficient in generating massive BNS systems (Safarzadeh et al. 2020). In this scenario, Galaudage et al. (2021) further estimated that the fast-merging binaries constitute 8%–79% of BNS at birth and have a delay time (from the birth to the death) of ∼5 − 401 Myr (90% credibility). Romero-Shaw et al. (2020) pointed out that the constraint on the eccentricity (e ≤ 0.07 at 10 Hz with 90% confidence) of GW190425 cannot provide evidence for or against the unstable MT channel. Early on, it was found in Vigna-Gómez et al. (2018) that Case BB MT from He-rich stars onto NSs is most likely dynamically stable, which is in agreement with earlier findings in Tauris et al. (2015). Recent investigations in Giacobbo & Mapelli (2018) found that, in order to produce BNSs with total masses of 3.2 − 3.5 M⊙ from isolated binaries, the metallicity should be low, namely ∼5%–10% solar metallicity (Z⊙). Studies with population synthesis analysis from Kruckow (2020) show the formation of GW190425 with Milky Way-like metallicity is possible. Alternatively, assuming the first-born NS with a typical mass of ∼1.4 M⊙ and allowing for super-Eddington accretion onto NSs, Zhang et al. (2023) employed detailed binary evolution to find that the parameter space required to reproduce GW190425-like events is very limited. Additionally, fallback can contribute significantly to the mass growth of the newly formed NS, which may help explain the formation of heavy BNS systems like GW190425 (Vigna-Gómez et al. 2021).

It is worth noting that the inferred component masses are dependent on the choice of spin priors (see their Figure 3 in Abbott et al. 2020). Zhu & Ashton (2020) proposed a spin prior, extrapolated from radio pulsar observations of Galactic binary NSs, finding positive support for a spinning recycled NS in GW190425. The inferred NS masses are 1 . 64 − 0.11 + 0.13 M ⊙ $ 1.64_{-0.11}^{+0.13}\, M_\odot $ for recycled NS and 1 . 66 − 0.12 + 0.12 M ⊙ $ 1.66_{-0.12}^{+0.12}\, M_\odot $ for slow NS, respectively. There is a slight discrepancy for inferred masses due to different spin priors, while the two masses are still consistent with the individual binary components being NSs. Additionally, the further constraint on the spin tilt angle (≲60°, Zhu & Ashton 2020) shows that this system is consistent with the standard formation channel of isolated binary evolution. Mandel et al. (2021) proposed a probabilistic prescription for compact remnant masses, consistent with the formation of GW190425.

In the context of a massive, binary evolution scenario for the formation of GW190425, the immediate progenitor is a close binary system of a He-rich star and a NS. In order to form a NS rather than a black hole, the mass of its progenitor He-rich star at solar metallicity needs to be less massive (for example, ∼3.0 − 7.5 M⊙, Zhang et al. 2023). After the core-helium burning phase, the star expands and then initiates MT onto the first-born NS through Case BB/BC MT. The late MT can further stripe He-rich star, producing Type Ib/c SNe (for example, Tauris et al. 2015). If the system remains bound, the progenitor of the second-born NS could be significantly spun up via tidal interaction, leading to a fast-spinning magnetized NS (namely magnetar, see details in Hu et al. 2023).

In this work, with the newly inferred component masses of NSs for GW190425, we perform detailed binary evolution with Modules for Experiments in Stellar Astrophysics (MESA) to explore its possible formation history. Therefore, the remainder of this paper is organized as follows. The main methods implemented in detailed binary evolution calculations are introduced in Section 2. In Section 3, we show the possible formation path of GW190425 and also estimate the rotational energy of the newly formed magnetar. Furthermore, we present in Section 4 the helium envelope mass before the SN explosion and the associated ejecta mass. Finally, our main conclusions and some discussion are summarized in Section 5.

2. Methods

We performed detailed binary modeling by using the release version mesa-r15140 of the Modules for Experiments in Stellar Astrophysics (MESA) stellar evolution code (Paxton et al. 2011, 2013, 2015, 2018, 2019; Jermyn et al. 2023). We followed the same method as recent studies (for example, Hu et al. 2023; Fragos et al. 2023; Lyu et al. 2023; Zhang et al. 2023; Zhu et al. 2024) to create zero-age helium main sequence (namely ZamsHe) with different masses and adopt Z⊙ = 0.0142 as the solar metallicity (Asplund et al. 2009).

We modeled convection using the mixing-length theory (Böhm-Vitense 1958) with a parameter of αmlt = 1.93. We adopted the Ledoux criterion to treat the boundaries of the convective zone and the step overshooting as an extension given by αp = 0.1Hp, where Hp is the pressure scale height at the Ledoux boundary limit. Semiconvection (Langer et al. 1983) with an efficiency parameter α = 1.0 was adopted in our modeling. The network of approx12.net was chosen for nucleosynthesis. We treated rotational mixing and angular momentum transport as diffusive processes (Heger & Langer 2000), including the effects of the Goldreich–Schubert–Fricke instability, Eddington–Sweet circulations, as well as secular and dynamical shear mixing. We adopted diffusive element mixing from these processes with an efficiency parameter of fc = 1/30 (Chaboyer & Zahn 1992; Heger & Langer 2000). The μ-gradient is sensitive to the rotationally induced mixing and here we reduced μ-gradient by multiplying fμ (fμ = 0.05 suggested in Heger & Langer 2000).

Stellar winds of He-rich stars were modeled following the same method as in Hu et al. (2022). We modeled the evolution of He-rich stars to reach central carbon depletion, from which the baryonic remnant mass is calculated following the “delayed” supernova prescription (Fryer et al. 2012). For NSs, we converted baryonic to gravitational mass following the approach in Lattimer & Yahil (1989), Timmes et al. (1996). We also considered the neutrino loss as in Zevin et al. (2020). The maximum mass of a NS is assumed to be 2.5 M⊙. We calculated the timescale for orbital synchronization following Hurley et al. (2002) for massive stars with radiative envelopes and adopted the updated fitting formula for the tidal coefficient E2 provided in Qin et al. (2018). We assumed the Eddington-limited accretion onto NSs using the standard formulae (see Eq. (4) in Tauris et al. 2017). MT was modeled following the Kolb scheme (Kolb & Ritter 1990) and the implicit MT method (Paxton et al. 2015) is adopted.

3. Formation of GW190425-like events and newborn magnetars

thumbnail Fig. 1.

Posterior distributions of the primary and secondary masses of GW190425 by considering low-spin prior (χ < 0.05; pink) and high-spin prior (χ < 0.89; blue). The probability density functions of m1 and m2 normalized to have equal maxima are displayed in the top and right panels, respectively.

Hereafter, we directly use the newest release version2 of GW190425 within GWTC-2 to explore its formation pathway. We first employ the posterior data obtained from the PhenomPNRT-LS (PhenomPNRT-HS) template for the low-spin (high-spin) prior3. Considering 90% confidence intervals, the inferred primary and secondary NS mass are m 1 = 2 . 02 − 0.34 + 0.58 M ⊙ $ m_1 = 2.02^{+0.58}_{-0.34}\, M_\odot $ and m 2 = 1 . 35 − 0.26 + 0.26 M ⊙ $ m_2 = 1.35^{+0.26}_{-0.26}\, M_\odot $ ( m 1 = 1 . 74 − 0.09 + 0.17 M ⊙ $ m_1 = 1.74^{+0.17}_{-0.09}\, M_\odot $ and m 2 = 1 . 56 − 0.14 + 0.08 M ⊙ $ m_2 = 1.56^{+0.08}_{-0.14}\, M_\odot $) using the low-spin (high-spin) prior. Figure 1 illustrates the posterior distributions of the primary and secondary masses for GW190425, which are nearly consistent with the results in the discovered report of this BNS merger (Abbott et al. 2020).

With the newly inferred component masses of NSs, we perform detailed binary evolution of He-rich stars and NS as a point mass in close orbits. We model He-rich stars with initial mass MZamHe linearly from 2.5 M⊙ to 8.0 M⊙ at a step of 0.5 M⊙. We cover the initial orbital period, ranging from 0.04 days to 40 days with a logarithmic spacing of Δlog(P/days)≈0.16 dex. We use Z = Z⊙ as the initial metallicity of He-rich stars.

thumbnail Fig. 2.

Accreted mass (the color bar) as a function of the initial orbital period and initial mass of He-rich stars. Cross: Initial Overflow; square: Case BA; triangle: Case BB; diamond: Case BC); circle: No mass transfer (MT). Left panel: high-spin prior (NS mass m1 = 2.02 M⊙), right panel: low-spin prior (NS mass m1 = 1.74 M⊙). We mark the parameter space of Case BB and Case BC with light and dark grey backgrounds, respectively.

3.1. Mass accretion onto NS

He-rich stars, especially for initially less massive, expand by around two orders of magnitude after leaving their core-helium burning phase (see their Figure 1 in Zhang et al. 2023). Therefore, MT between a NS and a He-rich star via Roche-lobe overflow is expected to occur in a close orbit. In Figure 2, we present the mass accreted onto NSs under different initial conditions of NS–He-rich star binary systems.

In the left panel of Figure 2, the NS is assumed to have a mass of m1 = 2.02 M⊙ (namely high-spin prior). First, MT in all binary systems is found not to occur for initial orbital periods Porb, init ≳ 40 days. Additionally, a longer initial orbital period is required for binary systems to experience mass interaction via L1 as a He-rich star is less massive. This finding has been demonstrated in earlier investigations (for example, Ivanova et al. 2003; Tauris et al. 2015; Zhang et al. 2023), which is because less massive He-rich stars expand significantly when evolving to become giant stars (for example, see the findings in Fragos et al. 2023; Zhang et al. 2023). Third, He-rich stars with an initial mass from 2.5 M⊙ to 7.5 M⊙ initiate MT onto NS companions when the central carbon is ignited (namely so-called Case BC MT, see the dark grey region). Compared with Case BC MT, He-rich stars can have MT during the shell-helium burning phase (namely Case BB MT phase, see light grey region) when the stars are less massive and the orbital periods are shorter (namely Porb ≲1.0 days). We note that several systems with initially shorter periods are found to initiate MT even earlier, namely during the core-helium burning phase (Case BA MT, see the square symbols). For initial orbital periods Porb ≲ 0.06 days, systems are more likely to merge with their companions due to the initial overflow of He-rich stars at ZamsHe. Additionally, the parameter space of the binary systems to experience MT is the same when a massive NS is assumed (see the right panel of Figure 2).

Figure 3 shows the accumulated mass accreted onto NSs through Case BB and Case BC MT, respectively. Given the NS companion mass of 2.02 M⊙ (see the left panel), the accreted mass via Case BB MT is found to be a Gaussian-like distribution, in the range of ∼2.5 × 10−4 − 6.3 × 10−3 M⊙ (log(Maccreted/M⊙)∈[∼ − 3.6, ∼ − 2.2]) and with a peak at ∼1.6 × 10−3 M⊙ (log(Maccreted/M⊙)∼ − 2.8). In contrast, the accreted mass through Case BC MT is lower, approximately from ∼3.2 × 10−5 M⊙ to ∼1.8 × 10−3 M⊙ (log (Maccreted/M⊙)∈[∼ − 4.5, ∼ − 2.8]). The difference in the accreted mass between Case BB and Case BC MT is ascribed to two factors. First, the MT rate of Case BB through the whole evolution is slightly higher when compared to Case BC MT. More importantly, the duration of the accretion for Case BB is much longer. These results are consistent with earlier findings in Zhang et al. (2023). When considering a lower-mass NS companion (see the right panel), the corresponding mass accreted through Case BB and Case BC MT has a similar distribution, with the whole range slightly shifted to a lower end.

thumbnail Fig. 3.

Histogram of accreted mass for Case BB (blue) and Case BC (green) MT. Left panel: high-spin prior (m1 = 2.02 M⊙), right panel: low-spin prior (m1 = 1.74 M⊙). All the He-rich stars have a metallicity of Z = Z⊙.

3.2. GW190425 formed through stable mass transfer

After the carbon is depleted in the center of He-rich stars, we adopt the “delayed” supernova prescription (Fryer et al. 2012) to distinguish the genres of various compact objects (COs) and then calculate the baryonic remnant mass of NSs. In Figure 4, we present the NS mass with different initial conditions of binary systems. It is shown in the left panel that a white dwarf (WD) is formed when the initial He-rich star has a mass below ∼3.0 M⊙, above which a NS is produced by the iron core-collapse supernova (CCSN). Notably, the mass of the newborn NS, determined by the inner structure of its progenitor before the SN stage, is independent of the initial orbital periods considered in this work and the companion star mass (see the right panel). He-rich stars at solar metallicity are expected to lose more mass due to metallicity-dependent winds (Vink et al. 2001), producing NSs with mass ranging from ∼1.2 M⊙ to ∼2.4 M⊙ (see the color bar).

thumbnail Fig. 4.

As in Figure 2, but the color bar refers to the NS mass. Square: white dwarf; triangle up: NS formed through iron core-collapse SN (CCSN).

After the formation of double NSs, gravitational wave emission shrinks the orbits by removing their orbital angular momentum and leads to the merger of the COs. We adopt the following expression given in Peters (1964) to calculate the merger time of BNSs,

T merger = 5 256 c 5 a f 4 G 3 ( m 1 + m 2 ) 2 m r T ( e ) , $$ \begin{aligned} T_{\rm merger}=\frac{5}{256} \frac{c^{5} a_{f}^{4}}{G^{3} (m_1 + m_2)^{2} m_{r}} T(e) , \end{aligned} $$(1)

where c is the speed of light and mr is the binary’s reduced mass, m1 and m2 are the first-born and second-born NS, as well as af the separation between the two components. For simplicity, we assume the orbit of BNS at its formation is circular (namely T(e) = 1). As we are mainly focused on BNS systems formed through stable MT after the post-CE phase, the SN kicks imparted onto the second-born NS are not taken into account. This has been pointed out in Tauris et al. (2015) that the SN kicks of NSs formed from ultra-stripped progenitors are small (see more detailed discussions in Section 6.2). Similarly, electron-capture SN is generally expected to have small kicks (Vigna-Gómez et al. 2018; Giacobbo & Mapelli 2019, 2020). For short orbital periods (≲1 day), the system remains bound when a newly-born NS has a kick velocity of less than 100 km/s (see their Figure 18 in Tauris et al. 2017).

thumbnail Fig. 5.

As in Figure 2, but the color bar refers to the merger time of BNS due to the emission of GWs. We use the red plus symbol to mark the systems that resemble GW190425-like events.

In Figure 5, we show the merger time (Tmerger, see the color bar) of BNSs with their Tmerger ≲ 14 Gyr (the Hubble age). It is shown that GW190425-like events (see the red plus symbols) can be formed through either Case BC or Case BB MT depending on the specific initial conditions. In the left panel, in order to form observable BNSs (the colored symbols), the initial orbital periods need to be shorter than ∼0.5 days (log(Pinit/day)∼ − 0.3). Assuming the first-born NS mass of m2 = 2.02 M⊙, we adopt the red plus symbols to mark the targets representing the GW190425-like events, with log(Tmerger/Myr) in the range of ∼1.37 − 4.14. We note that the parameter space to form GW190425-like events is rather limited, namely 3.0 ≲ MzamsHe/M⊙ ≲ 5.5 and 0.08 ≲ Pinit/days ≲ 0.5. When considering the first-born NS mass of m1 = 1.74 M⊙, the clear difference is that the immediate progenitor mass shifts to higher values, but with a rather narrow range, namely 5.5 ≲ MzamsHe/M⊙ ≲ 6.0. The merger time, however, is slightly shorter, namely 1.10 ≲ log(Tmerger/Myr)≲3.85.

3.3. Rotational energy of the newly formed magnetar

A NS accreting mass from its companion can be significantly spun up to become a mildly recycled pulsar (see details in Tauris et al. 2012, 2015, 2017). The progenitor of the second-born NS if in close orbits, is expected to be synchronized by its companion through tides (Detmers et al. 2008; Qin et al. 2018), potentially resulting in fast-spinning highly-magnetized NS (Hu et al. 2023).

thumbnail Fig. 6.

As in Figure 2, but the color bar refers to the rotational energy of the newly formed magnetar in binary systems whose merger time is no longer than the Hubble age.

We calculate the magnetar’s initial rotational energy, Erot, i = 0.5Imag(2π/Prot, i)2. In order to estimate the magnetar’s moment of inertia, we adopt the empirical relation from Worley et al. (2008),

I mag ≈ 0.237 M R 2 [ 1 + 4.2 M M ⊙ km R + 90 ( M M ⊙ km R ) 4 ] , $$ \begin{aligned} I_{\rm mag} \approx 0.237 M R^2 \left[1 + 4.2\frac{M}{M_\odot }\frac{\mathrm{km}}{R} + 90\left(\frac{M}{M_\odot }\frac{\mathrm{km}}{R}\right)^4 \right], \end{aligned} $$(2)

where M and R are the magnetar’s mass and radius, respectively. For the radius of a NS, we use a constant R of 12.5 km (Abbott et al. 2020; Landry et al. 2020).

In Figure 6, we present the magnetar’s initial rotational energy for the systems whose merger time is no longer than the Hubble age. It is worth noting that the magnetar’s initial rotational energy varies in the range of ∼1050 − ∼1052 erg. Therefore, the newly formed magnetar from tidal spin-up is a promising progenitor to produce stripped-envelope supernovae, including Type Ic superluminous supernovae, broad-line Type Ic SNe, and fast blue optical transients. More relevant studies can be found in Hu et al. (2023) who systematically investigated the origin of magnetar-driven stripped-envelope supernovae.

4. Helium envelope mass before the SN explosion and its associated ejecta mass

He-rich stars that experience Case BB/BC MT leading to further envelope loss are generally considered to produce ultra-stripped supernovae (namely ultra-stripped SNe, Tauris et al. 2015). Figure 7 shows the remaining envelope mass of He-rich stars before the SN explosion under different initial conditions. It is found that more envelope mass is retained in He-rich stars when increasing their initial orbital periods. This is because He-rich stars tend to be more significantly stripped in closer orbits. In the left panel (the companion NS mass M1 = 2.02 M⊙), Case BC MT allows He-rich stars to retain helium envelope mass in the range of ∼0.3 − 1.7 M⊙, while Case BB MT further strips He-rich stars, resulting in less helium envelope mass from ∼0.1 M⊙ to ∼1.2 M⊙. The envelope mass could be used as an upper limit of the ejecta mass during the SN explosion. As inferred in Hachinger et al. (2012), the upper limit of helium mass required for type SNe Ic is ∼0.06 − 0.14 M⊙. Therefore, our models indicate that stable Case BB MT in post-CE binaries can be considered a potential channel to contribute a certain proportion of all type Ic events (≲1%; Tauris et al. 2013). In contrast, more helium envelope mass could be retained for He-rich stars when experiencing Case BC MT, which is more likely to produce SNe Ib events. Therefore, GW190425-like sources are likely associated with type Ib/c SN events.

thumbnail Fig. 7.

As in Figure 2, but the color bar refers to the helium envelope mass (material outside the carbon-oxygen core) retained in He-rich stars before the SN explosion.

Recently, Hu et al. (2023) employed the population synthesis study to systematically investigate the evolution of He-rich stars with various companion stars, finding that He-rich stars in close orbits evolve to form fast-spinning magnetars. With the same method adopted in Hu et al. (2023), we combine the carbon/oxygen core mass and the remnant mass to estimate He-rich stars/ ejecta mass (carbon/oxygen core mass - remnant mass). In Figure 8, the ejecta mass is found to be in the range of ∼0.25 − 2.2 M⊙. Therefore, the formation of GW190425-like events could be accompanied by some transients, for example, type Ic superluminous supernovae, broad-line Type Ic SNe, and fast blue optical transients (FBOTs). We refer readers of interest to Hu et al. (2023) for more detailed calculations.

thumbnail Fig. 8.

As in Figure 2, but the color bar refers to ejecta mass.

5. Conclusions and discussion

We first obtain the component masses of GW190425, m 1 = 2 . 02 − 0.34 + 0.58 M ⊙ $ m_1 = 2.02^{+0.58}_{-0.34}\,M_\odot $ and m 2 = 1 . 35 − 0.26 + 0.26 M ⊙ $ m_2 = 1.35^{+0.26}_{-0.26}\,M_\odot $ ( m 1 = 1 . 74 − 0.09 + 0.17 M ⊙ $ m_1 = 1.74^{+0.17}_{-0.09}\,M_\odot $ and m 2 = 1 . 56 − 0.14 + 0.08 M ⊙ $ m_2 = 1.56^{+0.08}_{-0.14}\,M_\odot $) at 90% confidence intervals using the high-spin (low-spin) prior. Assuming the immediate progenitor of GW190425 is a He-rich star and a NS, we perf form a detailed binary evolution modeling to investigate its possible formation history and explore the properties of the relevant transients associated with the second SN explosion. In order to reproduce GW190425-like events with high-spin (low-spin) prior, the immediate progenitors should have an initial mass of a He-rich star in the range of ∼3.0 − 5.5 M⊙ (∼5.5 − 6.0 M⊙) and an initial period from ∼0.08 − 0.5 days (∼0.08 − 0.4 days), respectively. The mass accreted onto NS with mass inferred using the high-spin prior is found to vary from ∼2.5 × 10−4 to ∼6.3 × 10−3 M⊙ for Case BB MT and ∼3.2 × 10−5 to ∼1.8 × 10−3 M⊙ for Case BC MT, respectively. When considering a lower-mass (namely 1.74 M⊙) NS companion, the corresponding mass accreted through Case BB and Case BC MT has a similar distribution, with the whole range slightly shifted to a lower end. Additionally, the corresponding merger time log(Tmerger/Myr) varies from ∼1.37 to ∼4.14 (from ∼1.10 to ∼3.85 for low-spin prior).

Due to subsequent mass transfer onto their companions, He-rich stars are further stripped and thus are more likely to produce ultra-stripped SNe (Tauris et al. 2015). We find that the remaining helium envelope mass before SN explosion for He-rich stars is in the range of ∼0.3 ∼ 1.7 M⊙ for Case BC MT and ∼0.1 − 1.2 M⊙ for Case BB MT. Therefore, He-rich stars that experience extra mass transfer could potentially produce type Ib/c SN. We then estimate the ejecta mass in the range of ∼0.25 − 2.2 M⊙. Recently, Wu & Fuller (2022) found that less massive He-rich stars with mass of ≈2.5 − 3 M⊙ expand by a factor of a few during O/Ne-burning and thus can be further stripped due to later MT phase. Therefore, the ejecta mass estimated above is considered the upper limit.

We note that the SN 2023zaw, a sub-luminous and rapidly evolving SN with the lowest nickel mass, was recently reported to have a helium envelope mass of ∼0.2 M⊙ (Das et al. 2024). Their findings suggest an ultra-stripped SN, originating from a low-mass He-rich star in a close binary system. Furthermore, Moore et al. (2024) adopted Bayesian analysis to find that, in addition to the radioactive decay of 56Ni, an extra energy source (for example, a magnetar or interaction with circumstellar material) is required. As demonstrated in Section 3.3, He-rich stars in close orbits can be efficiently spun up (for example, Detmers et al. 2008; Qin et al. 2018; Sciarini et al. 2024), forming fast-spinning magnetars. For the newly formed magnetar, we further estimate its rotational energy in the range of ∼1050 − 1052 erg, which is considered a promising progenitor to produce stripped-envelope supernovae, such as Type Ic superluminous supernovae, broad-line Type Ic SNe, and fast blue optical transients (see detailed modeling in Hu et al. 2023). Recently, Siebert et al. (2024) reported a discovery of SN Ic, 2023adta identified in the deep James Webb Space Telescope (JWST)/NIRCam imaging. Follow-up observations with JWST/NIRSpec suggest the classification as abroad-line Type Ic SN. The estimated explosion parameters (ejecta mass and kinetic energy in their Table 5) for SN 2023adta are roughly consistent with the findings explored in this work. It should be noted that the spin period of the resultant NS is dependent on the magnetic braking (for example, Deng et al. 2021), although it is unclear whether the He-rich star has a strong magnetic field on the surface.

It is possible that the formation scenario explored in this work could be applied to the formation of GW170817. On top of that, two high-confidence black hole-neutron star mergers (GW200105 and GW200115, Abbott et al. 2021) were found to have low misaligned spin, indicating that their immediate progenitors just after the CE phase consisted of a first-born black hole and low-mass He-rich star. The subsequent evolution very likely involves the stable Case BB/BC MT phase (Jiang et al. 2023). More recently, the LVK collaboration (Abac et al. 2024) reported the observation of a coalescing compact binary (GW230529) with two component masses 2.5 − 4.5 M⊙ and 1.2 − 2.0 M⊙ (at the 90% credible level). This binary system was inferred to have low effective inspiral spin, namely χ eff = − 0 . 1 − 0.17 + 0.12 $ \chi_{\mathrm{eff}} = -0.1_{-0.17}^{+0.12} $. We also note that GW190425 could be the merger of a light NS and a low-mass BH (Han et al. 2020). If the massive component is a black hole, their progenitor systems (GW230529 and GW190425) could share the same origin as GW200105 and GW200115.

It is important to highlight that our findings could be dependent on the model adoption, for example, remnant mass prescriptions (for example, Fryer et al. 2012). Predictions of Galactic BNS mass distribution in Vigna-Gómez et al. (2018) show that more massive components are expected (see their Figure 7) using the “delayed” supernova prescription when compared with “rapid” and “ m u ¨ ller $ \rm m \ddot{u} ller $ ” (Müller et al. 2016) prescription. Therefore, the population of heavy BNS systems is expected to be small. However, predictions of the event rate for such systems are beyond the scope of this study. In the case of the high-spin prior, the primary NS mass of GW190425 was inferred to be ∼2.0 M⊙. This requires a massive NS to be born in the first SN explosion, for example, PSR J1640+2224 (Deng et al. 2020), PSR J1614−2230 (Tauris et al. 2011), and 2A 1822-371 (Wei et al. 2023).

Notably, Moroianu et al. (2023) reported a possible association (2.8σ level) between GW190425 (Abbott et al. 2020) and a fast radio burst (FRB 20190425A) (CHIME/FRB Collaboration 2021), which could be consistent with the FRB model invoking the collapse of a supermassive neutron star (Falcke & Rezzolla 2014) following a BNS merger (Zhang 2014). Further studies of Magaña Hernandez et al. (2024), however, suggested that current state-of-the-art GW analyses disfavor the association between GW190425 and FRB 20190425A. Numerical simulations in Yamasaki et al. (2018) showed that a fraction of BNS mergers may form fast-rotating stable NSs, reproducing repeating FRBs (CHIME/FRB Collaboration 2019) like FRB 121102. CHIME/FRB Collaboration (2019) proposed that a BNS merger channel could also form a millisecond magnetar (also see Wang et al. 2020). For massive BNS systems, the primary component can be spun up through accretion but remains strong magnetic fields (Chu et al., in prep.), leading to a limited radio-pulsar lifetime and a low probability of being detected by radio emission (also see claims in Safarzadeh et al. 2020). Their findings show that this could be the main reason that none of the observed Galactic BNSs is as massive as GW190425.

1

Case BB MT initiates from a He-rich star onto its companion when the donor star is burning helium in its shell, whereas Case BC MT occurs when carbon is ignited in the stellar core.

3

We refer the dimensionless component spin magnitudes < 0.05 (< 0.89) to low-spin (high-spin) prior.

Acknowledgments

We thank the referee for constructive comments that helped improve the manuscript. We also thank Tassos Fragos and Wen-Cong Chen for their helpful comments on the manuscript. Y.Q. acknowledges support from the Anhui Provincial Natural Science Foundation (grant No. 2308085MA29), the National Natural Science Foundation of China (grant No. 12473036), and funding from the Key Laboratory for Relativistic Astrophysics at Guangxi University. J.P.Z. thanks the COMPAS group at Monash University. G.M. has received funding from the European Research Council (ERC) under the European Union’s Horizon 2020 research and innovation program (grant agreement No 833925, project STAREX. Q.W.T acknowledges support from Jiangxi Provincial Natural Science Foundation (grant Nos. 20242BAB26012 and 20224ACB211001). This work was partially supported by the National Natural Science Foundation of China (grant Nos. 12065017, 12192220, 12192221, 12133003, 12203101, U2038106, 12103003). All figures are made with the free Python module Matplotlib (Hunter 2007).

References

  1. Abac, A. G., Abbott, R., Abouelfettouh, I., et al. 2024, ApJ, 970, L34 [CrossRef] [Google Scholar]
  2. Abbott, B. P., Abbott, R., Abbott, T. D., et al. 2017, Phys. Rev. Lett., 119, 161101 [Google Scholar]
  3. Abbott, B. P., Abbott, R., Abbott, T. D., et al. 2020, ApJ, 892, L3 [NASA ADS] [CrossRef] [Google Scholar]
  4. Abbott, R., Abbott, T. D., Abraham, S., et al. 2021, ApJ, 915, L5 [NASA ADS] [CrossRef] [Google Scholar]
  5. Acernese, F., Agathos, M., Agatsuma, K., et al. 2015, Class. Quant. Grav., 32, 024001 [Google Scholar]
  6. Aso, Y., Michimura, Y., Somiya, K., et al. 2013, Phys. Rev. D, 88, 043007 [NASA ADS] [CrossRef] [Google Scholar]
  7. Asplund, M., Grevesse, N., Sauval, A. J., & Scott, P. 2009, ARA&A, 47, 481 [NASA ADS] [CrossRef] [Google Scholar]
  8. Böhm-Vitense, E. 1958, Z. Astrophys., 46, 108 [Google Scholar]
  9. Chaboyer, B., & Zahn, J. P. 1992, A&A, 253, 173 [NASA ADS] [Google Scholar]
  10. CHIME/FRB Collaboration (Andersen, B. C., et al.) 2019, ApJ, 885, L24 [Google Scholar]
  11. CHIME/FRB Collaboration (Amiri, M., et al.) 2021, ApJS, 257, 59 [NASA ADS] [CrossRef] [Google Scholar]
  12. Coughlin, M. W., Ahumada, T., Anand, S., et al. 2019, ApJ, 885, L19 [CrossRef] [Google Scholar]
  13. Das, K. K., Fremling, C., Kasliwal, M. M., et al. 2024, ApJ, 969, L11 [NASA ADS] [CrossRef] [Google Scholar]
  14. Deng, Z.-L., Gao, Z.-F., Li, X.-D., & Shao, Y. 2020, ApJ, 892, 4 [NASA ADS] [CrossRef] [Google Scholar]
  15. Deng, Z.-L., Li, X.-D., Gao, Z.-F., & Shao, Y. 2021, ApJ, 909, 174 [NASA ADS] [CrossRef] [Google Scholar]
  16. Detmers, R. G., Langer, N., Podsiadlowski, P., & Izzard, R. G. 2008, A&A, 484, 831 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  17. Falcke, H., & Rezzolla, L. 2014, A&A, 562, A137 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  18. Farrow, N., Zhu, X.-J., & Thrane, E. 2019, ApJ, 876, 18 [NASA ADS] [CrossRef] [Google Scholar]
  19. Fragos, T., Andrews, J. J., Bavera, S. S., et al. 2023, ApJS, 264, 45 [NASA ADS] [CrossRef] [Google Scholar]
  20. Fryer, C. L., Belczynski, K., Wiktorowicz, G., et al. 2012, ApJ, 749, 91 [Google Scholar]
  21. Galaudage, S., Adamcewicz, C., Zhu, X.-J., Stevenson, S., & Thrane, E. 2021, ApJ, 909, L19 [NASA ADS] [CrossRef] [Google Scholar]
  22. Giacobbo, N., & Mapelli, M. 2018, MNRAS, 480, 2011 [Google Scholar]
  23. Giacobbo, N., & Mapelli, M. 2019, MNRAS, 482, 2234 [Google Scholar]
  24. Giacobbo, N., & Mapelli, M. 2020, ApJ, 891, 141 [NASA ADS] [CrossRef] [Google Scholar]
  25. Habets, G. M. H. J. 1986, A&A, 167, 61 [NASA ADS] [Google Scholar]
  26. Hachinger, S., Mazzali, P. A., Taubenberger, S., et al. 2012, MNRAS, 422, 70 [NASA ADS] [CrossRef] [Google Scholar]
  27. Han, M.-Z., Tang, S.-P., Hu, Y.-M., et al. 2020, ApJ, 891, L5 [CrossRef] [Google Scholar]
  28. Heger, A., & Langer, N. 2000, ApJ, 544, 1016 [CrossRef] [Google Scholar]
  29. Hosseinzadeh, G., Cowperthwaite, P. S., Gomez, S., et al. 2019, ApJ, 880, L4 [NASA ADS] [CrossRef] [Google Scholar]
  30. Hu, R.-C., Zhu, J.-P., Qin, Y., et al. 2022, ApJ, 928, 163 [NASA ADS] [CrossRef] [Google Scholar]
  31. Hu, R. C., Zhu, J. P., Qin, Y., et al. 2023, ArXiv e-prints [arXiv:2301.06402] [Google Scholar]
  32. Hunter, J. D. 2007, Comput. Sci. Eng., 9, 90 [NASA ADS] [CrossRef] [Google Scholar]
  33. Hurley, J. R., Tout, C. A., & Pols, O. R. 2002, MNRAS, 329, 897 [Google Scholar]
  34. Ivanova, N., Belczynski, K., Kalogera, V., Rasio, F. A., & Taam, R. E. 2003, ApJ, 592, 475 [NASA ADS] [CrossRef] [Google Scholar]
  35. Jermyn, A. S., Bauer, E. B., Schwab, J., et al. 2023, ApJS, 265, 15 [NASA ADS] [CrossRef] [Google Scholar]
  36. Jiang, L., Tauris, T. M., Chen, W.-C., & Fuller, J. 2021, ApJ, 920, L36 [NASA ADS] [CrossRef] [Google Scholar]
  37. Jiang, L., Chen, W.-C., Tauris, T. M., Müller, B., & Li, X.-D. 2023, ApJ, 945, 90 [NASA ADS] [CrossRef] [Google Scholar]
  38. Kolb, U., & Ritter, H. 1990, A&A, 236, 385 [NASA ADS] [Google Scholar]
  39. Kruckow, M. U. 2020, A&A, 639, A123 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  40. Landry, P., Essick, R., & Chatziioannou, K. 2020, Phys. Rev. D, 101, 123007 [NASA ADS] [CrossRef] [Google Scholar]
  41. Langer, N., Fricke, K. J., & Sugimoto, D. 1983, A&A, 126, 207 [NASA ADS] [Google Scholar]
  42. Lattimer, J. M., & Yahil, A. 1989, ApJ, 340, 426 [NASA ADS] [CrossRef] [Google Scholar]
  43. LIGO Scientific Collaboration (Aasi, J., et al.) 2015, Class. Quant. Grav., 32, 074001 [Google Scholar]
  44. Lyu, F., Yuan, L., Wu, D. H., et al. 2023, MNRAS, 525, 4321 [NASA ADS] [CrossRef] [Google Scholar]
  45. Magaña Hernandez, I., d’Emilio, V., Morisaki, S., Bhardwaj, M., & Palmese, A. 2024, ApJ, 971, L5 [CrossRef] [Google Scholar]
  46. Mandel, I., Müller, B., Riley, J., et al. 2021, MNRAS, 500, 1380 [Google Scholar]
  47. Moore, T., Gillanders, J., Nicholl, M., et al. 2024, ArXiv e-prints [arXiv:2405.13596] [Google Scholar]
  48. Moroianu, A., Wen, L., James, C. W., et al. 2023, Nat. Astron., 7, 579 [NASA ADS] [CrossRef] [Google Scholar]
  49. Müller, B., Heger, A., Liptai, D., & Cameron, J. B. 2016, MNRAS, 460, 742 [Google Scholar]
  50. Paczynski, B. 1976, in Structure and Evolution of Close Binary Systems, eds. P. Eggleton, S. Mitton, & J. Whelan, 73, 75 [NASA ADS] [CrossRef] [Google Scholar]
  51. Paxton, B., Bildsten, L., Dotter, A., et al. 2011, ApJS, 192, 3 [Google Scholar]
  52. Paxton, B., Cantiello, M., Arras, P., et al. 2013, ApJS, 208, 4 [Google Scholar]
  53. Paxton, B., Marchant, P., Schwab, J., et al. 2015, ApJS, 220, 15 [Google Scholar]
  54. Paxton, B., Schwab, J., Bauer, E. B., et al. 2018, ApJS, 234, 34 [NASA ADS] [CrossRef] [Google Scholar]
  55. Paxton, B., Smolec, R., Schwab, J., et al. 2019, ApJS, 243, 10 [Google Scholar]
  56. Peters, P. C. 1964, Phys. Rev., 136, 1224 [Google Scholar]
  57. Qin, Y., Fragos, T., Meynet, G., et al. 2018, A&A, 616, A28 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  58. Romero-Shaw, I. M., Farrow, N., Stevenson, S., Thrane, E., & Zhu, X.-J. 2020, MNRAS, 496, L64 [NASA ADS] [CrossRef] [Google Scholar]
  59. Safarzadeh, M., Ramirez-Ruiz, E., & Berger, E. 2020, ApJ, 900, 13 [NASA ADS] [CrossRef] [Google Scholar]
  60. Sciarini, L., Ekström, S., Eggenberger, P., et al. 2024, A&A, 681, L1 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  61. Siebert, M. R., DeCoursey, C., Coulter, D. A., et al. 2024, ApJ, 972, L13 [NASA ADS] [CrossRef] [Google Scholar]
  62. Tauris, T. M., Langer, N., & Kramer, M. 2011, MNRAS, 416, 2130 [NASA ADS] [CrossRef] [Google Scholar]
  63. Tauris, T. M., Langer, N., & Kramer, M. 2012, MNRAS, 425, 1601 [NASA ADS] [CrossRef] [Google Scholar]
  64. Tauris, T. M., Langer, N., Moriya, T. J., et al. 2013, ApJ, 778, L23 [NASA ADS] [CrossRef] [Google Scholar]
  65. Tauris, T. M., Langer, N., & Podsiadlowski, P. 2015, MNRAS, 451, 2123 [NASA ADS] [CrossRef] [Google Scholar]
  66. Tauris, T. M., Kramer, M., Freire, P. C. C., et al. 2017, ApJ, 846, 170 [Google Scholar]
  67. Timmes, F. X., Woosley, S. E., & Weaver, T. A. 1996, ApJ, 457, 834 [NASA ADS] [CrossRef] [Google Scholar]
  68. Vigna-Gómez, A., Neijssel, C. J., Stevenson, S., et al. 2018, MNRAS, 481, 4009 [Google Scholar]
  69. Vigna-Gómez, A., Schrøder, S. L., Ramirez-Ruiz, E., et al. 2021, ApJ, 920, L17 [CrossRef] [Google Scholar]
  70. Vink, J. S., de Koter, A., & Lamers, H. J. G. L. M. 2001, A&A, 369, 574 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  71. Wang, F. Y., Wang, Y. Y., Yang, Y.-P., et al. 2020, ApJ, 891, 72 [NASA ADS] [CrossRef] [Google Scholar]
  72. Wei, N., Jiang, L., & Chen, W.-C. 2023, A&A, 679, A74 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  73. Worley, A., Krastev, P. G., & Li, B.-A. 2008, ApJ, 685, 390 [NASA ADS] [CrossRef] [Google Scholar]
  74. Wu, S. C., & Fuller, J. 2022, ApJ, 940, L27 [NASA ADS] [CrossRef] [Google Scholar]
  75. Yamasaki, S., Totani, T., & Kiuchi, K. 2018, PASJ, 70, 39 [NASA ADS] [CrossRef] [Google Scholar]
  76. Zevin, M., Spera, M., Berry, C. P. L., & Kalogera, V. 2020, ApJ, 899, L1 [Google Scholar]
  77. Zhang, B. 2014, ApJ, 780, L21 [Google Scholar]
  78. Zhang, J., Yang, Y., Zhang, C., et al. 2019, MNRAS, 488, 5020 [NASA ADS] [CrossRef] [Google Scholar]
  79. Zhang, W. T., Wang, Z. H. T., Zhu, J. P., et al. 2023, MNRAS, 526, 854 [NASA ADS] [CrossRef] [Google Scholar]
  80. Zhao, C., Lu, Y., Chu, Q., & Zhao, W. 2023, MNRAS, 522, 912 [NASA ADS] [CrossRef] [Google Scholar]
  81. Zhu, X.-J., & Ashton, G. 2020, ApJ, 902, L12 [NASA ADS] [CrossRef] [Google Scholar]
  82. Zhu, J.-P., Qin, Y., Wang, Z.-H.-T., et al. 2024, MNRAS, 529, 4554 [NASA ADS] [CrossRef] [Google Scholar]

All Figures

thumbnail Fig. 1.

Posterior distributions of the primary and secondary masses of GW190425 by considering low-spin prior (χ < 0.05; pink) and high-spin prior (χ < 0.89; blue). The probability density functions of m1 and m2 normalized to have equal maxima are displayed in the top and right panels, respectively.
In the text
thumbnail Fig. 2.

Accreted mass (the color bar) as a function of the initial orbital period and initial mass of He-rich stars. Cross: Initial Overflow; square: Case BA; triangle: Case BB; diamond: Case BC); circle: No mass transfer (MT). Left panel: high-spin prior (NS mass m1 = 2.02 M⊙), right panel: low-spin prior (NS mass m1 = 1.74 M⊙). We mark the parameter space of Case BB and Case BC with light and dark grey backgrounds, respectively.
In the text
thumbnail Fig. 3.

Histogram of accreted mass for Case BB (blue) and Case BC (green) MT. Left panel: high-spin prior (m1 = 2.02 M⊙), right panel: low-spin prior (m1 = 1.74 M⊙). All the He-rich stars have a metallicity of Z = Z⊙.
In the text
thumbnail Fig. 4.

As in Figure 2, but the color bar refers to the NS mass. Square: white dwarf; triangle up: NS formed through iron core-collapse SN (CCSN).
In the text
thumbnail Fig. 5.

As in Figure 2, but the color bar refers to the merger time of BNS due to the emission of GWs. We use the red plus symbol to mark the systems that resemble GW190425-like events.
In the text
thumbnail Fig. 6.

As in Figure 2, but the color bar refers to the rotational energy of the newly formed magnetar in binary systems whose merger time is no longer than the Hubble age.
In the text
thumbnail Fig. 7.

As in Figure 2, but the color bar refers to the helium envelope mass (material outside the carbon-oxygen core) retained in He-rich stars before the SN explosion.
In the text
thumbnail Fig. 8.

As in Figure 2, but the color bar refers to ejecta mass.
In the text

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