# SN2023fyq: A Type Ibn Supernova With Long-standing Precursor Activity Due to
Binary Interaction

Yize Dong (董一泽) Department of Physics and Astronomy, University of California,
1 Shields Avenue, Davis, CA 95616-5270, USA  Daichi Tsuna TAPIR, Mailcode
350-17, California Institute of Technology, Pasadena, CA 91125, USA  Research
Center for the Early Universe (RESCEU), School of Science, The University of
Tokyo, 7-3-1 Hongo, Bunkyo-ku, Tokyo 113-0033, Japan  Stefano Valenti
Department of Physics and Astronomy, University of California, 1 Shields
Avenue, Davis, CA 95616-5270, USA  David J. Sand Steward Observatory,
University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA
Jennifer E. Andrews Gemini Observatory, 670 North A‘ohoku Place, Hilo, HI
96720-2700, USA  K. Azalee Bostroem LSSTC Catalyst Fellow  Steward
Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ
85721-0065, USA  Griffin Hosseinzadeh Steward Observatory, University of
Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA  Emily Hoang
Department of Physics and Astronomy, University of California, 1 Shields
Avenue, Davis, CA 95616-5270, USA  Saurabh W. Jha Department of Physics and
Astronomy, Rutgers, the State University of New Jersey,
136 Frelinghuysen Road, Piscataway, NJ 08854-8019, USA  Daryl Janzen Department of Physics & Engineering Physics, University of Saskatchewan, 116 Science Place, Saskatoon, SK S7N 5E2, Canada  Jacob E. Jencson Department of Physics and Astronomy, Johns Hopkins University, 3400 North Charles Street, Baltimore, MD 21218, USA  Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA  Michael Lundquist W. M. Keck Observatory, 65-1120 Māmalahoa Highway, Kamuela, HI 96743-8431, USA  Darshana Mehta  Department of Physics and Astronomy, University of California, 1 Shields Avenue, Davis, CA 95616-5270, USA  Aravind P. Ravi Department of Physics and Astronomy, University of California, 1 Shields Avenue, Davis, CA 95616-5270, USA  Nicolas E. Meza Retamal Department of Physics and Astronomy, University of California, 1 Shields Avenue, Davis, CA 95616-5270, USA  Jeniveve Pearson Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA  Manisha Shrestha Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA  Alceste Z. Bonanos IAASARS, National Observatory of Athens, Metaxa & Vas. Pavlou St., 15236, Penteli, Athens, Greece  D. Andrew Howell Las Cumbres Observatory, 6740 Cortona Drive, Suite 102, Goleta, CA 93117-5575, USA  Department of Physics, University of California, Santa Barbara, CA 93106-9530, USA  Nathan Smith Steward Observatory, University of Arizona, 933 North Cherry Avenue, Rm. N204, Tucson, AZ 85721-0065, USA  Joseph Farah  Las Cumbres Observatory, 6740 Cortona Drive, Suite 102, Goleta, CA 93117-5575, USA  Department of Physics, University of California, Santa Barbara, CA 93106-9530, USA  Daichi Hiramatsu Center for Astrophysics | Harvard & Smithsonian, 60 Garden Street, Cambridge, MA 02138-1516, USA  The NSF AI Institute for Artificial Intelligence and Fundamental Interactions, USA  Koichi Itagaki (板垣公一)  Itagaki Astronomical Observatory, Yamagata 990-2492, Japan  Curtis McCully Las Cumbres Observatory, 6740 Cortona Drive, Suite 102, Goleta, CA 93117-5575, USA  Department of Physics, University of California, Santa Barbara, CA 93106-9530, USA  Megan Newsome  Las Cumbres Observatory, 6740 Cortona Drive, Suite 102, Goleta, CA 93117-5575, USA  Department of Physics, University of California, Santa Barbara, CA 93106-9530, USA  Estefania Padilla Gonzalez Las Cumbres Observatory, 6740 Cortona Drive, Suite 102, Goleta, CA 93117-5575, USA  Department of Physics, University of California, Santa Barbara, CA 93106-9530, USA  Emmanouela Paraskeva IAASARS, National Observatory of Athens, Metaxa & Vas. Pavlou St., 15236, Penteli, Athens, Greece  Craig Pellegrino Department of Astronomy, University of Virginia, Charlottesville, VA 22904, USA  Giacomo Terreran Las Cumbres Observatory, 6740 Cortona Drive, Suite 102, Goleta, CA 93117-5575, USA  Department of Physics, University of California, Santa Barbara, CA 93106-9530, USA  Joshua Haislip Department of Physics and Astronomy, University of North Carolina, 120 East Cameron Avenue, Chapel Hill, NC 27599, USA  Vladimir Kouprianov Department of Physics and Astronomy, University of North Carolina, 120 East Cameron Avenue, Chapel Hill, NC 27599, USA  Daniel E. Reichart Department of Physics and Astronomy, University of North Carolina, 120 East Cameron Avenue, Chapel Hill, NC 27599, USA

###### Abstract

We present photometric and spectroscopic observations of SN 2023fyq, a type
Ibn supernova in the nearby galaxy NGC 4388 (D$\simeq$18 Mpc). In addition, we
trace long-standing precursor emission at the position of SN 2023fyq using
data from DLT40, ATLAS, ZTF, ASAS-SN, Swift, and amateur astronomer Koichi
Itagaki. Precursor activity is observed up to nearly three years before the
supernova explosion, with a relatively rapid rise in the final 100 days. The
double-peaked post-explosion light curve reaches a luminosity of $\sim
10^{43}~{}\rm erg\,s^{-1}$. The strong intermediate-width He lines observed in
the nebular spectrum of SN 2023fyq imply the interaction is still active at
late phases. We found that the precursor activity in SN 2023fyq is best
explained by the mass transfer in a binary system involving a low-mass He star
and a compact companion. An equatorial disk is likely formed in this process
($\sim$0.6$\rm M_{\odot}$), and the interaction of SN ejecta with this disk
powers the main peak of the supernova. The early SN light curve reveals the
presence of dense extended material ($\sim$0.3$\rm M_{\odot}$) at
$\sim$3000$\rm R_{\odot}$ ejected weeks before the SN explosion, likely due to
final-stage core silicon burning or runaway mass transfer resulting from
binary orbital shrinking, leading to rapid rising precursor emission within
$\sim$30 days prior to explosion. The final explosion could be triggered
either by the core-collapse of the He star or by the merger of the He star
with a compact object. SN 2023fyq, along with SN 2018gjx and SN 2015G, forms a
unique class of Type Ibn SNe which originate in binary systems and are likely
to exhibit detectable long-lasting pre-explosion outbursts with magnitudes
ranging from $-$10 to $-$13.

Core-collapse supernovae (304), Circumstellar matter (241), Stellar mass loss
(1613)

††facilities: ADS, DLT40 (Prompt5, Prompt-MO), ATLAS, LCOGT (SBIG, Sinistro,
FLOYDS), Gemini:North (GMOS), Keck:I (LRIS, DEIMOS), NED, SOAR (Goodman),
Swift (UVOT), LBT (MODS) ††software: Astropy (Astropy Collaboration et al.,
2013, 2018), emcee (Foreman-Mackey et al., 2013) HOTPANTS (Becker, 2015),
Matplotlib (Hunter, 2007), NumPy (Harris et al., 2020), PYRAF (Science
Software Branch at STScI, 2012), Pandas (Wes McKinney, 2010), SciPy (Virtanen
et al., 2020), SWarp (Bertin et al., 2002), HOTPANTS (Becker, 2015),
LCOGTSNpipe (Valenti et al., 2016), Light Curve Fitting (Hosseinzadeh & Gomez,
2020), LPipe (Perley, 2019)

##  1 Introduction

Type Ibn supernovae (SNe) are a subclass of interaction-powered SNe that show
narrow helium (He) lines but not hydrogen (H) lines in their spectra (e.g.,
Smith, 2017; Modjaz et al., 2019). Although it has been more than two decades
since the discovery of the first Type Ibn SN (SN 1999cp, Matheson et al.
2000), our understanding of Type Ibn progenitors remains limited. The light
curves of Type Ibn SNe tend to be short-lived and some of them even resemble
the evolution of fast-evolving transients (Ho et al., 2023; Fox & Smith,
2019). A general interpretation is that SNe Ibn are Wolf-Rayet/He stars that
experience enhanced mass loss right before the SN explosion. The interaction
of SN ejecta with the surrounding dense He-rich circumstellar material (CSM)
powers some of the SN light curve and ionizes the outer CSM, producing the
narrow lines we observe (Pastorello et al., 2007; Hosseinzadeh et al., 2017).

Light curve modeling of Type Ibn SNe has supported the presence of dense CSM
close to the progenitors (Gangopadhyay et al., 2020; Pellegrino et al., 2022;
Ben-Ami et al., 2023). Both SNe Ibn and their H-rich counterparts, SNe IIn,
have CSM interaction signatures that point to pre-SN mass loss that is much
stronger than normal massive-star winds (Smith, 2014, 2017). However, the
mechanisms driving the enhanced mass loss near the time of explosion remain a
subject of active debate. This enhanced mass loss could be attributed to the
final-stage stellar activities of massive stars, where the dense CSM could be
produced by eruptive outbursts through pulsational pair instability (Yoshida
et al., 2016; Woosley, 2017) or wave-driven outbursts excited by late-stage
nuclear burning (Quataert & Shiode, 2012; Shiode & Quataert, 2014; Fuller,
2017; Fuller & Ro, 2018; Morozova et al., 2020). Alternatively, the dense CSM
might be generated through binary interactions (Smith, 2014; Smith & Arnett,
2014; Metzger, 2022; Wu & Fuller, 2022; Dessart et al., 2022; Tsuna et al.,
2024). In this scenario the progenitor does not necessarily have to be a very
massive star, as the mass loss would be significantly enhanced by the presence
of a binary companion.

One way to constrain the progenitor of Type Ibn SNe is by searching for
evidence of a massive star or a binary companion in deep images once the SN
fades. The absence of evidence for massive star progenitors and the possible
detection of binary companions have been reported for some Type Ibn SNe (Maund
et al., 2016; Shivvers et al., 2017; Hosseinzadeh et al., 2019).

Alternatively, a direct way to constrain the mass loss history of SN
progenitors is by searching for signs of pre-explosion activity or precursor
emission prior to the SN explosion. Precursor emission is commonly observed in
Type IIn SNe (e.g., Mauerhan et al., 2013; Smith et al., 2010; Ofek et al.,
2013; Tartaglia et al., 2016; Pastorello et al., 2013, 2018; Hiramatsu et al.,
2024). The bright precursor outbursts in Type IIn SNe may be due to eruptive
mass loss from LBV-like progenitors (e.g., Smith, 2017) or pulsational pair
instability outbursts (Smith & McCray, 2007; Woosley et al., 2007; Smith,
2014). Alternatively, these outbursts could be caused by red supergiants with
a compact object companion (Fryer & Woosley, 1998; Schrøder et al., 2020;
Smith et al., 2024; Tsuna et al., 2024), or other late-stage binary
interaction (Smith & Arnett, 2014). To date, precursor emission has been
identified in two Type Ibn SNe, SN 2006jc (Pastorello et al., 2007) and SN
2019uo (Strotjohann et al., 2021). The precursor outbursts in these events are
shorter and fainter compared to those observed in Type IIn SNe, and have been
interpreted as resulting from single massive star activities or binary
interactions (Pastorello et al., 2007; Foley et al., 2007; Smith et al., 2008;
Tsuna et al., 2024).

In this paper we present the optical observations of SN 2023fyq, one of the
closest SNe Ibn. The light curves and spectra of this object closely resemble
those of Type Ibn SNe. Notably, relatively steady precursor activity is
observed up to approximately three years prior to the SN explosion. The
detection of precursor emission in SN 2023fyq allows us to investigate the
final-stage stellar activity and the nature of its progenitor system. The pre-
explosion observations of SN 2023fyq are also presented in Brennan et al.
(2024), where they identify an asymmetric CSM structure, likely related to
unstable stellar activities of the progenitor.

The paper is organized as follows: the photometric and spectroscopic
observations are described in Section 2. We constrain the reddening and
distance of SN 2023fyq in Section 3. We describe the photometric and
spectroscopic evolution of SN 2023fyq in Sections 4 and 5. The progenitor
scenario and the physical mechanism of precursor activities are discussed in
Section 6. We summarize the main results in Section 7.

Figure 1: Composite $gri$ image of SN 2023fyq in NGC 4388 obtained with the
Las Cumbres Observatory on 2023 August 11. The position of SN 2023fyq is
indicated by white tick markers.  Figure 2: Photometric limits and detections
of SN 2023fyq prior to and after explosion. Detections with S/N$>$4 are
indicated by large solid symbols, while detections with 3$<$S/N$\leq$4 are
indicated by hollow symbols. The smaller symbols are nondetection limits with
S/N$\leq$3\. The precursor activities detected in Type Ibn SN 2006jc ($R$
band) and SN 2019uo ($r$ band) are indicated in the red and green rectangles,
respectively. The limits on the precursor activities on Type Ibn SN 2015G are
shown with the purple dashed line. All of the bands are in the AB magnitude
system.

##  2 Observations

SN 2023fyq was discovered on 2023 April 17 by the Zwicky Transient Facility
(ZTF) survey at RA(2000) $=$ 12h25m45$\fs$847, Dec(2000) $=+12\arcdeg
39\arcmin 48\farcs 87$ in NGC 4388 (De, 2023) (see Figure 1). On 2023 June 14
a rapid rebrightening of SN 2023fyq was observed and reported by amateur
astronomer Koichi Itagaki. On 2023 June 25 SN 2023fyq was classified as a
peculiar Type Ib due the presence of helium lines and the lack of hydrogen
lines in the optical spectrum (Valerin et al., 2023).

In this section we present the photometric data of SN 2023fyq taken by Las
Cumbres Observatory (Brown et al., 2013) via the Global Supernova Project, the
Distance Less Than 40 Mpc (DLT40, Tartaglia et al., 2018) survey, ZTF (Bellm
et al., 2019; Graham et al., 2019), the Asteroid Terrestrial-Impact Last Alert
System (ATLAS, Tonry 2011; Tonry et al. 2018; Smith et al. 2020), the All-Sky
Automated Survey for Supernovae (ASAS-SN, Shappee et al. 2014; Kochanek et al.
2017), the Neil Gehrels Swift Observatory (Gehrels et al., 2004), and amateur
astronomer Itagaki. We also report the spectroscopic followup of SN 2023fyq
taken after the SN explosion. All spectroscopic observations from this paper
can be found at https://github.com/yizedong/SN2023fyq_data and will be
available on WISeREP (Yaron & Gal-Yam, 2012)111http://www.weizmann.ac.il.

Figure 3: The light curve evolution of SN 2023fyq. The $Clear$ filter is
calibrated to the $r$ band. The hollow symbol indicates the data with
3$<$S/N$\leq$4, while the solid symbol indicates the data with S/N$>$4\. Light
curves in the bottom panel have been shifted by the indicated amounts to
enhance clarity. All of the bands are in the AB magnitude system. The black
dashed line marks the epoch of the first light of the SN ($-$11 d), as adopted
in the paper.

###  2.1 Photometric Observations

For the photometry we adopt a signal-to-noise threshold of 3 for source
detections and a signal-to-noise threshold of 5 for computing the upper limit,
following the suggestions of Masci (2011). The light curves are shown in
Figure 2 and 3.

####  2.1.1 Las Cumbres Observatory Observations

Our multiband photometric followup campaign with Las Cumbres Observatory was
initiated on 2023 July 26. The images were reduced using the PyRAF-based
photometric reduction pipeline lcogtsnpipe (Valenti et al., 2016). Apparent
magnitudes were calibrated using the APASS ($g,r,i$) and Landolt ($U,B,V$)
catalogs.

####  2.1.2 DLT40 Observations

The DLT40 survey is a targeted one-day cadence SN search for very young
transients within 40 Mpc (Tartaglia et al., 2018; Yang et al., 2019).

DLT40 has been monitoring the field of SN 2023fyq since 2014 in the $Clear$
filer. All of the images have been visually inspected to remove those with bad
qualities. A deep template was made with the images taken between 2014 June 20
and 2015 February 01 using Swarp (Bertin et al., 2002). The rest of the images
were stacked in windows of 15 days and were then subtracted against the
template using HOTPANTS (Becker, 2015). We used aperture photometry at the
position of SN 2023fyq through a pipeline based on Photutils (Bradley et al.,
2022). The photometry was calibrated to the $r$ band.

####  2.1.3 ZTF Observations

ZTF is a time-domain survey using a wide-field camera mounted on the Palomar
48-inch Schmidt telescope (Bellm et al., 2019; Graham et al., 2019). The ZTF
public survey searches for transients and variables in the northern sky with a
three-day cadence in $g$ and $r$ filters.

The position of SN 2023fyq has been monitored by ZTF since 2018. We obtained
the reference image subtracted forced photometry from the ZTF Forced
Photometry Service (Masci et al., 2023). We removed bad-quality data following
the instructions in Masci et al. (2023). For images taken after $-$300d, the
transient was bright enough to be detected in single images, and so the
observations were stacked in 1-day time bins. For images taken prior to
$-$300d, the observations were stacked in 15-day time bins to improve the
signal to noise ratio (S/N).

####  2.1.4 ATLAS Observations

The ATLAS survey is an all-sky daily cadence survey (Smith et al., 2020)
carried out in two filters, cyan ($c$) and orange ($o$), roughly equivalent to
Pan-STARRS filters $g+r$ and $r+i$, respectively.

The position of SN 2023fyq has been monitored by ATLAS since 2015. Forced
photometry at the supernova position was obtained from the ATLAS forced
photometry server (Shingles et al., 2021). Using the method presented in Young
(2022), we stacked the measurements to improve the signal-to-noise ratio and
obtain deeper upper limits. For images taken after $-$300d, the observations
were stacked in 1-day time bins. For images taken before $-$300d, the
observations were stacked in 15-day time bins.

####  2.1.5 ASAS-SN Observations

ASAS-SN is an untargeted all-sky survey to a depth of g$\sim$18.5 mag.
(Shappee et al., 2014; Kochanek et al., 2017). We obtained the ASAS-SN
reference image subtracted forced photometry from the ASAS-SN sky
portal222https://asas-sn.osu.edu/.

####  2.1.6 Swift Observations

The position of SN 2023fyq has been observed by the UVOT instrument on the
Neil Gehrels Swift Observatory (Gehrels et al., 2004) since 2015. We performed
aperture photometry at the position of SN 2023fyq on Swift UVOT images using
the High-Energy Astrophysics software (HEA-Soft). Background variations in
individual images were removed using an aperture placed on a blank section of
the sky. To remove the underlying galaxy background contamination, we
subtracted the flux extracted from Swift UVOT images taken on 2016 November
08. Zero-points were chosen from Breeveld et al. (2011) with time-dependent
sensitivity corrections updated in 2020.

####  2.1.7 Koichi Itagaki’s Observations

We also incorporated observations taken with Koichi Itagaki’s Bitran BN-83MCCD
imager mounted on a 0.5m telescope in Okayama Prefecture, Japan. We solved the
astrometry of the images using Astrometry.net (Lang et al., 2010). The
aperture photometry was performed using a pipeline based on Photutils (Bradley
et al., 2022) and was calibrated to r-band magnitudes in the Sloan system
(Fukugita et al., 1996).

###  2.2 Spectroscopic Observations

We collected four optical spectra from the FLOYDS spectrograph (Brown et al.,
2013) on the 2m Faulkes Telescope South in Australia at the Las Cumbres
Observatory via the Global Supernova Project. The FLOYDS spectra were reduced
following standard procedures using the FLOYDS pipeline (Valenti et al.,
2014). We triggered Gemini-North Target of Opportunity (ToO) observations with
the Gemini Multi-Object Spectrograph (GMOS; Hook et al., 2004) and the B600
grating on 2023 July 27 and 2023 August 01 through proposal GN-2023A-Q-136.
The Gemini spectra were reduced by using the IRAF Gemini package. We triggered
further ToO observations with the Andalucia Faint Object Spectrograph and
Camera (ALFOSC) on the Nordic Optical Telescope (NOT) at the Spanish “Roque de
los Muchachos” Observatory (ORM) on 2023 August 04 through proposal 67-112.
The NOT ALFOSC spectrum was observed using Grism #4 and a 1.$\arcsec$0 slit
and was reduced using the PypeIt pipeline (Prochaska et al., 2020, 2020). We
obtained spectra on 2023 December 12 and 2024 May 1 from the Low-Resolution
Imaging Spectrometer (LRIS; Oke et al., 1995) on the Keck I telescope. The
LRIS spectra were reduced in a standard way using the LPipe pipeline (Perley,
2019). A low-resolution spectrum was taken on 2024 January 23 with the Goodman
High Throughput Spectrograph (GHTS) on the Southern Astrophysical Research
Telescope (SOAR; Clemens et al., 2004), and was reduced with the Goodman
pipeline (Torres et al., 2017). One spectrum was obtained with the Multi-
Object Double Spectrographs (MODS, Pogge et al., 2010) on the twin 8.4 m Large
Binocular Telescope (LBT) at Mount Graham International Observatory. The
spectrum was reduced using standard techniques, including bias subtraction and
flat-fielding using the MODSCCDred package (Pogge, 2019) and further reduced
with IRAF including cosmic ray rejection, local sky subtraction, and
extraction of one-dimensional spectra. A log of the spectroscopic observations
is presented in Table A1. We also present an unpublished nebular spectrum of
Type Ibn SN 2019kbj taken at 80 d after the peak. The spectrum was taken on
2019 September 23 with the DEep Imaging Multi-Object Spectrograph (DEIMOS,
Faber et al., 2003) on the Keck II telescope (Table A2). The DEIMOS spectrum
was reduced using the PypeIt pipeline (Prochaska et al., 2020, 2020). A
detailed analysis of SN 2019kbj has been presented in Ben-Ami et al. (2023).

Figure 4: $r/R$ Light curve comparison between SN 2023fyq, a sample of Type
Ibn SNe, and well-studied normal SESNe. The Vega magnitudes have been
converted to the AB magnitude system. The evolution of SN 2023fyq is similar
to those of Type Ibn SNe. The SNe used in this plot includes Type IIb SN 1993J
(Filippenko et al., 1993), Type Ib SN 2008D (Modjaz et al., 2009), Type Ic SN
2007gr (Hunter et al., 2009)), and Type Ibn SNe: SN 2015U (Tsvetkov et al.,
2015; Pastorello et al., 2015a; Hosseinzadeh et al., 2017), iPTF15ul
(Hosseinzadeh et al., 2017), iPTF14aki (Hosseinzadeh et al., 2017), iPTF15akq
(Hosseinzadeh et al., 2017), SN 2019deh (Pellegrino et al., 2022), SN 2021jpk
(Pellegrino et al., 2022), SN 2005la (Pastorello et al., 2008), SN 2020nxt
(Wangq et al., 2024), SN 2018gjx (Prentice et al., 2020), ASASSN-15ed
(Pastorello et al., 2015b), SN 2010al (Pastorello et al., 2015c), SN 2015G
(Shivvers et al., 2017; Hosseinzadeh et al., 2017), SN 2006jc (Pastorello et
al., 2007), SN 2019uo (Gangopadhyay et al., 2020), and SN 2019kbj (Ben-Ami et
al., 2023). SN 2018gjx, ASASSN-15ed, SN 2010al, SN 2015G, SN 2006jc, SN
2019uo, and SN 2019kbj will be used for further comparison in the paper, while
a broader sample of SNe Ibn are shown in tan. SESNe are shown in grey.

Figure 5: The pre- and post-explosion bolometric light curve (upper two
panels) and the blackbody temperature and radius evolution (bottom panel) of
SN 2023fyq at the precursor phases and the early SN phases. The uncertainties
are indicated by the shaded area.

##  3 Observational Properties

###  3.1 Reddening

The empirical correlation between the equivalent width (EW) of the Na I D line
and the amount of gas and dust along the line of sight has often been used in
extinction estimations (Munari & Zwitter, 1997). In order to measure the line-
of-sight reddening towards SN 2023fyq, we analyzed the medium-resolution
spectrum (R$\sim$1800) taken with Gemini North on 2023 August 1. The measured
EW of the host galaxy Na I D $\lambda$5890 ($\rm D_{2}$) and Na I D
$\lambda$5896 ($\rm D_{1}$) are $0.27\pm 0.04$ Å and $0.15\pm 0.04$ Å,
respectively. The measured EW of the Galactic Na I $\rm D_{2}$ and Na I $\rm
D_{1}$ are $0.23\pm 0.02$ Å and $0.16\pm 0.01$ Å respectively. Using Eq.9 in
Poznanski et al. (2012) and applying the renormalization factor of 0.86 from
Schlafly et al. (2010), we found a host extinction of $E(B-V)_{\rm
host}=0.037\pm 0.01$ mag. The Milky Way extinction is measured to be
$E(B-V)_{\rm MW}=0.035\pm 0.01$ mag, which is consistent with the Milky Way
extinction of $E(B-V)_{\rm MW}$ = 0.0286 mag from the extinction map by
Schlafly & Finkbeiner (2011). We adopt the latter for the Milky Way
extinction. Throughout the paper, we will adopt a total extinction of $E(B-V)$
= $0.066\pm 0.01$ mag.

We note that Brennan et al. (2024) found a larger host extinction value
($E(B-V)_{\rm host}=0.4\pm 0.1$ mag) using the Balmer ratio measured from the
host emission lines. The disagreement is probably because this method measures
the full column of gas including the background. In this case, there is likely
some dust between the SN and the underlying HII region, which is responsible
for this greater implied extinction value.

###  3.2 Distance

The distance of NGC 4388 listed on the NASA/IPAC Extragalactic Database (NED)
ranges from 13.6 Mpc to 25.7 Mpc ($\mu$ = 30.67 – 32.05 mag). We adopt the
most recent Tully-Fisher distance (based on photometry at 3.6$\mu$m with
Spitzer Space Telescope), 18.0$\pm$3.7 Mpc ($\mu$ = 31.28$\pm$0.45 mag; Tully
et al. 2016).

##  4 Photometric Evolution

In Figure 2 we present the photometric evolution of SN 2023fyq dating back to
2015, illustrating our search for precursor activities. In Figure 3 we take a
closer look at the evolution from one year before the SN explosion. All phases
mentioned in the paper are with respect to the maximum light in the $r$ band,
which is measured to be at JD = 2460154 after fitting the light curve with a
spline function. At $\sim-11$ d, a sudden rise of $\sim$1.5 mag within
$\sim$17 hrs is clearly observed (see lower panel of Figure 3). As we will
discuss below, we attribute this rapid rise to the SN first light.
Consequently, we divide the photometric evolution of SN 2023fyq into two
phases: the precursor phase ($<-11$ d) and the SN phase ($>-11$ d).

###  4.1 Precursor Detections

The precursor is detected from $\sim$$-$1000 d to $\sim$$-$11 d. There are
also single detections at around $-$2300 d and $-$1300 d. These detections
have 3$<$S/N$\leq$4, and are bracketed by nondetections of similar depth.
Therefore, they are likely not true detections of precursor emission. As
illustrated in Figure 2, the precursor activities remain relatively stable at
$-10$ to $-12$ mag between $\sim-$1300 d and $\sim-$100 d. Then, starting from
$-100$ d, the object slowly brightens to $\sim$$-15$ mag. Between
$\sim$$-2500$ and $\sim$$-100$ d, the UV observations from Swift only give
nondetection limits (See Figure 2). As the precursor gets brighter, at
$\sim$$-28$ d, a source is detected in the $UVW1$ filter at $\sim$$-$13 mag,
with similar magnitudes observed in $g$ and $o$ bands. From $-300$ to $-11$ d,
the precursor light curves seem to exhibit multiple bumps, indicative of pre-
explosion activities, such as small eruptions, from the progenitor star. As
shown in Figure 2, the precursor emission detected in SN 2023fyq appears
fainter and longer compared to that observed in Type Ibn SN 2006jc (Pastorello
et al., 2007) and SN 2019uo (Strotjohann et al., 2021), even when accounting
for uncertainties in the distance measurement of SN 2023fyq. Pre-explosion
activities were not detected for Type Ibn SN 2015G down to $-$13.3 $\pm$ 0.5
mag (Shivvers et al., 2017). It should be noted that the precursor searches
for SN 2006jc and SN 2019uo only go down to around $-$13 mag. Therefore,
fainter precursor activities like those observed in SN 2023fyq can not be
excluded for these events.

###  4.2 SN Light Curve

The bluer-band ($UVW2$, $UVM2$, $UVW1$) light curves of SN 2023fyq exhibit a
notable bump from $-11$ d to $-4$ d, before reaching the second peak and then
falling off rapidly. This initial bump in the blue bands is likely
attributable to the cooling following shock breakout. For the rest of the
bands, the SN light curves show a fast rise and also a fast decline. The peak
$r$-band magnitude is measured to be $M_{r}=-18.5$ mag. In Figure 4, we
compare the $r$-band light curve of SN 2023fyq with the $r/R$-band light
curves of a sample of Type Ibn SNe and well-studied normal stripped-envelop
SNe (SESNe). At early times SN 2023fyq appears more luminous than the typical
SESNe, and the evolution of SN 2023fyq is overall similar to those of Type Ibn
SNe. At late times SN 2023fyq declines similarly to SN 2018gjx and SN 2015G,
but slower than SN 2006jc. The steep decline of SN 2006jc in the optical is
likely due to dust formation in the SN ejecta or in the surrounding CSM (e.g.,
Smith et al., 2008). The slower decline of SN 2023fyq, SN 2018gjx, and SN
2015G at late times could be an indication of less efficient dust formation
than in SN 2006jc. However, due to the lack of late-phase observations of Type
Ibn SNe, it is not clear if SN 2006jc is really an outlier. SN 2023fyq
declines faster than normal SESNe at nebular phases. This may be due to an
inefficient trapping of $\gamma$-rays in SN 2023fyq if the light curve tail is
powered by $\rm{}^{56}Ni$ decay, a power source other than $\rm{}^{56}Ni$
decay, or dust formation in SN 2023fyq.

Figure 6: Left: The optical spectroscopic evolution of SN 2023fyq. The phase
is measured from the $r$-band maximum. Right: The evolution of the He I
$\lambda$5876 line. The pre-maximum spectra marked in grey are from Brennan et
al. (2024). The He I $\lambda$5876 line shows a high-velocity component
(marked with the blue band) and a low-velocity component (marked with the red
band), which may come from the SN ejecta and He-rich CSM, respectively. The
grey bands mark the emission lines from the galaxy.  Figure 7: Optical
spectral comparison of SN 2023fyq at $\sim$0 d to other Type Ibn SNe and
normal SESNe.

Figure 8: Upper: Optical spectral comparison of SN 2023fyq at $\sim$7 d to
other Type Ibn SNe and normal SESNe. Bottom: The optical spectrum taken at
$\sim$7 d compared to the mean spectra (the solid lines) and the standard
deviations (the shaded regions) of SN Ib and Ic at $\sim$10 d from Liu et al.
(2016). SN 2023fyq has several features in common with these normal SESNe,
suggesting SN 2023fyq is likely from an explosion of a stripped star.  Figure
9: Nebular spectral comparison of SN 2023fyq to other Type Ibn SNe with
nebular spectra and normal SESNe. The phases are relative to the time of
maximum light. A continuum spectrum of the background galaxy is subtracted
from the spectrum of SN 2023fyq. At nebular phases, SNe Ibn appear to fall
into two distinct classes: one exhibiting only narrow He lines (SN 2019kbj and
SN 2006jc), and another displaying intermediate-width He lines and oxygen
lines (SN 2023fyq, SN 2015G, and SN 2018gjx).

###  4.3 Bolometric Light Curve

We constructed the bolometric light curve of SN 2023fyq using data from ZTF,
ATLAS, ASAS-SN, Swift, and Itagaki. To build the spectral energy distribution
(SED) in the regions without complete multiband coverage, we reconstruct the
multiband light curves using a neural network based light curve fitting method
presented in Demianenko et al. (2023). This method is able to capture
correlations across different observations over time and among various
passbands, and compute an approximate light curve within the specified time
and wavelength ranges. The final bolometric light curve is calculated by
fitting the SED with a blackbody function using a Markov Chain Monte Carlo
(MCMC) routine in the Light Curve Fitting package (Hosseinzadeh & Gomez,
2020). The blackbody temperatures measured from the pre-explosion spectra of
SN 2023fyq in Brennan et al. (2024) are used as priors for the SED fitting. We
present the bolometric light curve of SN 2023fyq, and the corresponding
blackbody temperature ($T_{BB}$) and radius ($R_{BB}$), in the precursor phase
and the SN phase, in Figure 5. We note that we only focus on the long-term
evolution of the bolometric light curve, and small variations in the light
curves are not reflected in the final bolometric light curve.

Before $\sim$$-100$ d, the precursor of SN 2023fyq is in a relatively stable
state with a luminosity of $\sim 1\times 10^{40}$ erg s-1. During that time,
$T_{BB}$ and $R_{BB}$ are around 10,000 K and 600 $\rm R_{\odot}$,
respectively. After $-$100 d, SN 2023fyq shows a faster rise and, at
$\sim$$-$11 d, the luminosity suddenly increases over an order of magnitude
(i.e., from $\sim 4\times 10^{41}$ erg s-1 to $\sim 7\times 10^{42}$ erg s-1).
Later, after a brief decline, the SN reaches its main peak and declines
afterwards. The decline of luminosity shortly after $\sim$$-11$ d is likely
due to the shock cooling after the shock breakout. For $T_{BB}$, after jumping
to $\sim$22,000 K at $\sim$$-$11 d, it rapidly declines until entering a brief
plateau phase between $\sim$$-5$ and $0$ d with $T_{BB}$$\simeq$10,000K. The
initial rapid decrease of $T_{BB}$ is likely associated with the shock cooling
process, while the plateau phase is likely due to the recombination of He I
and will be further discussed in section 6.1. After around $-40$ d, $R_{BB}$
shows a gradual expansion with a velocity of $\sim$700 $\rm km\,s^{-1}$. After
$-11$ d, $R_{BB}$ continuously increase, reflecting an increase of the
photospheric radius with the expansion of SN ejecta. The expansion rate of
$R_{BB}$ is $\sim$14,000 $\rm km\,s^{-1}$ initially, which slows down to
$\sim$7000 $\rm km\,s^{-1}$ after around -2 d. After around 5 d, as will be
discussed in the next section, the spectra of SN 2023fyq are dominated by
absorption lines from the SN ejecta, so $R_{BB}$ may not accurately reflect
the position of the photosphere.

##  5 Spectroscopic Evolution

The spectroscopic evolution of SN 2023fyq is presented in Figure 6. At -1.3 d,
the spectrum shows a blue continuum with a prominent He I $\lambda$5876 line.
Other He lines, such as He I $\lambda$5015, He I $\lambda$6678, He I
$\lambda$7065, and He I $\lambda$7281, are also observed. The He I
$\lambda$5876 line shows a rather asymmetric profile (right panel of Figure
6). In the blue wing, the He I $\lambda$5876 line shows a two-component
profile, with a narrow absorption feature at $\sim$$-$1000 $\rm km\,s^{-1}$
and a broad absorption feature at $\sim$$-$7000 $\rm km\,s^{-1}$. The
detection of a two-component He I line profile in SN 2023fyq is consistent
with those observed in other Type Ibn SNe (Pastorello et al., 2016), and is
likely from different emitting regions. The broad component is from the fast
moving ejecta, while the narrow component is likely from the surrounding
unshocked He-rich CSM. In the red wing, there is an additional emission
component at around 1500 $\rm km\,s^{-1}$. This component is also observed
during the pre-explosion phase of SN 2023fyq (Brennan et al., 2024), and could
be due to an asymmetric CSM structure formed before the SN explosion. A few
days later the object quickly becomes redder, and the Ca II H&K
$\lambda\lambda 3934,3969$ and Ca II $\lambda\lambda$8498, 8542, 8662 lines
appear more prominent. No broad hydrogen features are observed in the spectra
of SN 2023fyq. However, we can not exclude the presence of narrow hydrogen
lines since the spectra are heavily contaminated by the host-galaxy emission.
At $\sim$137 d, the spectrum is dominated by strong [O I]
$\lambda\lambda$6300, 6364 and [Ca II] $\lambda\lambda$7291, 7323. He lines,
such as He I $\lambda$5876 and He I $\lambda$7065 are also strong at this
phase. Other lines, including Mg I] $\lambda$4571 and Ca II
$\lambda\lambda$8498, 8542, 8662, can be seen in the spectrum. After that, the
spectra we have are mainly dominated by the host, while weak [O I]
$\lambda\lambda$6300, 6364 lines are still present.

We compare the spectra of SN 2023fyq around 0 d and 7 d with other SNe Ibn and
normal SESNe at similar phases in Figure 7 and Figure 8. At around 0 d, other
SNe Ibn show blue continua plus narrow He I $\lambda$5876 lines in their
spectra. The velocities of those narrow He I $\lambda$5876 lines are
consistent with that of the narrow component of the He I $\lambda$5876 line in
SN 2023fyq. At around 0 d, normal SESNe are redder than SN 2023fyq and other
SNe Ibn. This is probably due to the ongoing CSM interaction in the SNe Ibn,
which is not significant in SESNe. SESNe start to show lines from iron-group
elements at this phase, whereas these features are not strong in SN 2023fyq or
other SNe Ibn at a similar phase. The He lines in Type Ib/IIb SNe are also
much broader than those shown in SN 2023fyq.

At around 7 d, SN 2023fyq is very similar to SNe Ibn SN 2018gjx, ASASSN-15ed,
SN 2010al, and SN 2015G, which start to show signatures from deeper layers of
the ejecta. The He I $\lambda$5876 lines of SN 2018gjx, ASASSN-15ed, SN
2010al, and SN 2015G grow broader, with velocities similar to that of the
broad component of He I $\lambda$5876 in SN 2023fyq. Interestingly, some
similarities between SN 2023fyq and normal SESNe are also observed at around 7
d. To better illustrate this, we flatten the spectrum of SN 2023fyq at $\sim$7
d using SNID following the procedure outlined in Blondin & Tonry (2007) and
compare the flattened spectrum with Type Ib and Ic templates at 10 d from Liu
et al. (2016) in the bottom panel of Figure 8. This comparison clearly
indicates that SN 2023fyq exhibits spectral features similar to those of Type
Ic SNe, suggesting that its progenitor is likely a stripped/He star.

When the object enters the nebular phase, the ejecta become optically thin,
providing an unique opportunity to study the core of the progenitor star.
However, it is challenging to follow up SNe Ibn at nebular phases since they
rapidly get fainter. In Figure 9, we compare the nebular spectrum of SN
2023fyq at $\sim$136.5d with a few SNe Ibn with late-time observations and
normal SESNe at similar phases. The underlying continuum of the background
galaxy, obtained from a pre-explosion spectrum taken at -504 d as presented in
Brennan et al. (2024) when the signal from the host is dominant, is subtracted
from the spectrum presented here. SN 2023fyq shows strong intermediate-width
He emission lines, similar to Type Ibn SN 2018gjx and SN 2015G, but the [O I]
$\lambda\lambda$6300, 6364 line in SN 2023fyq is significantly stronger than
those in other objects. Type Ibn SN 2006jc and SN 2019kbj only show narrow He
lines and have no signatures of oxygen. SNe Ibn at nebular phases seem to fall
into two distinct classes, with one still showing only narrow lines and
another showing intermediate-width He lines and oxygen lines. This topic will
be further discussed in Section 6.4. Compared to normal SESNe SN 1993J, SN
2008D, and SN 2007gr, SN 2023fyq shows prominent He emission lines, but
otherwise SN 2023fyq is similar to those normal SESNe at the nebular phase.

Overall, the spectroscopic evolution SN 2023fyq is similar to those of some
SNe Ibn. However, the difference between SESNe and SN 2023fyq shortly after
the light curve maximum is less evident. A transition between Type Ibn and
Type Ic is clearly observed. Similar behaviors have been reported in several
previous studies of other Type Ibn SNe (e.g., Pastorello et al., 2015b;
Prentice et al., 2020). If SN 2023fyq is indeed dominated by CSM interaction
at peak light, the transition to Type Ic could be due to the CSM-interaction
region becoming transparent over time, allowing us to see more signatures from
the SN ejecta. It is also possible that the SN ejecta has moved beyond the
dense CSM. This suggests that SN 2023fyq is likely exploded from a stripped/He
star within He-rich CSM. The He lines observed at the nebular phase indicate
that the interaction with the He-rich CSM is still ongoing. It is natural to
link the pre-existing He-rich CSM with the pre-explosion activities of the
progenitor system, which likely also produces the precursor emission observed
in SN 2023fyq. This topic will be further discussed in Section 6.3.

##  6 Discussions

The detection of sustained precursor emission in SN 2023fyq provides an
invaluable opportunity to study the progenitor system of Type Ibn SNe. Below
is a summary of the primary observed characteristics of SN 2023fyq:

  1. 1.

A long-standing and continuously rising precursor emission starting from years
before the SN explosion;

  2. 2.

The light curve following the explosion exhibits an evolution similar to Type
Ibn SNe; the bolometric light curve exhibits two peaks.

  3. 3.

The early- and late-phase spectra both show narrow/intermediate-width He
lines. The nebular spectra show prominent [O I] $\lambda\lambda$6300, 6364
emission, suggesting that SN 2023fyq is likely a stripped/He star exploded
within He-rich CSM.

Any progenitor scenario for SN 2023fyq needs to explain the above behaviors.
In this section we will discuss the progenitor system and possible powering
mechanisms of the precursor and the SN light curve.

###  6.1 What Powers The First Peak of The SN Bolometric Light Curve?

The light curve of SN 2023fyq reaches its initial peak at around $-$11 d. The
later decrease of luminosity is associated with a prompt decline of $T_{BB}$
and a rapid expansion of $R_{BB}$. This process is likely the shock cooling
phase after the shock breakout. During this phase, the expansion of the ejecta
is nearly adiabatic, converting the thermal energy into kinetic energy. The
rapid decline of the photospheric temperature can produce a decrease in
brightness in bluer bands and an increase in brightness in redder bands as the
temperature moves through the optical bands, which is consistent with what we
see in SN 2023fyq (Figure 3). It is noteworthy that, around the shock
breakout, $R_{BB}$ is about 2$\sim$3$\times 1000\rm~{}R_{\odot}$
($\sim$1-2$\times 10^{14}$ cm), so the shock breakout likely originates from
an extended envelope/CSM wind instead of from the stellar surface. A similar
conclusion is also drawn by Brennan et al. (2024) based on the pre-explosion
spectroscopic and photometric observations of SN 2023fyq.

When $T_{BB}$ drops down to $\simeq$10,000 K, it enters a brief plateau phase
(Figure 5). Interestingly this temperature is consistent with the
recombination temperature of He I ($\sim$10,000 K) (Kleiser & Kasen, 2014). In
the meantime, the expansion of $R_{BB}$ slows down. Given that the early SN
spectra are dominated by He lines, the outer envelope is likely He-rich. We
argue that this $T_{BB}$ plateau phase is due to the recombination of He I,
and the decrease of $R_{BB}$ expansion rate is due to the recession of the
photosphere into the extended envelope. After this process, the outer envelope
becomes almost transparent due to the drop of electron scattering opacity.
This is consistent with the fact that we start to see more signals, such as Ca
lines, from the deeper SN ejecta after 0 d.

In conclusion, the first peak of the SN bolometric light curve of SN 2023fyq
is likely due to shock breakout in an extended envelope/CSM wind located at
$\sim$$2000-3000\rm~{}R_{\odot}$.

###  6.2 What Powers The Second Peak of The SN Bolometric Light Curve?

At 0 d, SN 2023fyq reaches its second peak. It should be noted that all bands
(from UV to optical) show peaks at this phase, so this second peak is not an
effect of temperature evolution and is instead powered by energy processes.

####  6.2.1 radioactive decay (RAD)?

We first consider the possibility that the SN light curve around the second
peak is powered by the $\rm{}^{56}Ni$ decay. The early light curve evolution
of SNe is regulated by the photon diffusion time, which depends on the SN
ejecta mass, the ejecta velocity, and the opacity (Arnett, 1982). Assuming
that the rise time of the light curve is equal to the photon diffusion time
and Arnett’s law holds for this object, i.e., the peak luminosity is close to
the instantaneous decay power at the peak, we can estimate the $\rm{}^{56}Ni$
mass ($M_{Ni}$) and the ejecta mass ($M_{ej}$). We fix the optical opacity
$\kappa_{opt}$ to be 0.1 $\rm cm^{2}\,g^{-1}$. Given a peak luminosity of
$9.5\times 10^{42}$ erg $s^{-1}$, we get $M_{Ni}$$\simeq$0.28 $\rm M_{\odot}$
and $M_{ej}\simeq 0.54\rm M_{\odot}$($v_{ph}/\rm 7000km\,s^{-1})(t/10d)^{2}$.

Therefore, to power the light curve with only $\rm{}^{56}Ni$ decay, around
half of the ejecta is composed of $\rm{}^{56}Ni$. This ratio is much higher
than those in typical CCSNe (e.g., Lyman et al., 2016) and similar to those
found in Type Ia SNe (e.g., Könyves-Tóth et al., 2020; Graham et al., 2022).
If the ejecta is $\rm{}^{56}Ni$-rich, when the ejecta become optically thin,
the optical spectra would be dominated by forbidden lines from Fe and Co.
However, as we discussed in Section 5, the nebular spectrum of SN 2023fyq is
mainly dominated by He, O and Ca. Therefore, we disfavor the $\rm{}^{56}Ni$
decay as the dominant power source of the early light curve of SN 2023fyq.

Figure 10: Upper-Left: Fits to the bolometric light curve of SN 2023fyq using
a combination of shock breakout and CSM interaction models. Bottom: Fits to
the bolometric light curve of SN 2023fyq using a combination of shock
breakout, CSM interaction, and $\rm{}^{56}Ni$ decay models. The dip observed
around 30 days in the $\rm{}^{56}Ni$ decay model is due to the transition from
the photospheric phase to the nebular phase (see Valenti et al. (2008) for
more details). The upper-right panel is a zoom-in of the bottom panel to
better illustrate the fit close to the SN peak. The initial bump is well-
fitted by the shock breakout model. The hollow point is at the precursor
phase, so it is not included in the fit.

####  6.2.2 CSM interaction?

Since the evolution of SN 2023fyq is similar to those of Type Ibn SNe, it is
likely that the light curve around the second peak is powered by CSM
interaction. We use the model presented in Jiang et al. (2020), which
generalizes the self-similar solution to the interaction of stellar ejecta
with surrounding CSM originally presented in (Chevalier, 1982). In this model,
the density of CSM is described by a power law, $\rho\propto qr^{-s}$, while
the ejecta are divided by an inner region ($\rho_{ej}\propto r^{-\delta}$) and
an outer region ($\rho_{ej}\propto r^{-n}$). We fix the optical opacity
($\kappa$) to be 0.1 $\rm cm^{2}\,g^{-1}$, $n=10$, $s=0$, and $\delta=1$
following Pellegrino et al. (2022). The value of $\kappa\approx 0.1\ {\rm
cm^{2}\ g^{-1}}$ is motivated by the opacity of singly-ionized He at $\sim
10^{4}$ K (e.g., Kleiser & Kasen, 2014). We also attempted to fit the data
with $s=2$ (wind-like CSM), but did not achieve a reasonable fit. This result
is consistent with the findings reported by Karamehmetoglu et al. (2017),
Gangopadhyay et al. (2020), and Ben-Ami et al. (2023). The ejecta velocity
(7,000 $\rm km\,s^{-1}$) is obtained from the velocity of the P-Cygni minimum
of the He I lines near peak. The free parameters in our fit are the explosion
epoch ($t_{exp}$), the ejecta mass ($M_{ej}$), the inner radius of the CSM
($R_{0}$), the CSM mass ($M_{csm}$), the density of the CSM at $R_{0}$
($\rho_{csm,0}$), and the conversion efficiency of the shock kinetic energy to
radiation ($\epsilon$).

To account for the initial shock cooling phase we have incorporated the shock
breakout (SBO) model presented by (Margalit, 2022). This model provides an
analytic solution for the shock cooling phase following shock breakout from
extended optically thick material, which is suitable for the case of SN
2023fyq. We fix the velocity of the inner envelope at 7,000 $\rm km\,s^{-1}$.
Additionally, we introduce two free parameters into our fit: the radius of the
extended material ($R_{e}$) and the mass of the extended material ($M_{e}$).

The model fit to the observed light curve is performed using an MCMC routine.
As illustrated in the upper-left panel of Figure 10, both the initial bump and
the subsequent evolution of the light curve are well-fitted by the model. The
best-fitting parameters are detailed in Table 1 (CSM+SBO model). It is
important to note that the models presented here are likely simplified, so the
parameters derived can only be considered as estimations of the order of
magnitude.

The $M_{ej}$ and $M_{csm}$ derived for SN 2023fyq are roughly consistent with
those found in other studies (e.g., Pellegrino et al., 2022; Ben-Ami et al.,
2023). The low ejecta mass implies that the progenitor is likely a low-mass He
star. However, this model can only fit the light curve around the peak and
cannot explain the light curve flattening at late times (see Figure 5). At
later times, the light curve is likely powered by another source of energy.

####  6.2.3 RAD+CSM interaction?

Since SN 2023fyq is similar to normal SESNe shortly after peak and during
nebular phases, it is plausible that a certain amount of $\rm{}^{56}Ni$ is
produced during the explosion. Therefore, it is natural to consider
$\rm{}^{56}Ni$ decay as an additional energy source. A $\rm{}^{56}Ni$ decay
model has been employed to interpret the late-time light curves of many other
Type Ibn SNe, often revealing low $\rm{}^{56}Ni$ masses across previous
studies (Gangopadhyay et al., 2020; Pellegrino et al., 2022; Ben-Ami et al.,
2023).

We use the $\rm{}^{56}Ni$ decay model presented in (Arnett, 1982; Valenti et
al., 2008). The full SN light curve is fitted by a combination of CSM
interaction, shock breakout, and $\rm{}^{56}Ni$ decay models. We fix the
optical opacity to be $\kappa=$0.1 $\rm cm^{2}\,g^{-1}$ and the $\gamma$-ray
opacity to be 0.03 $\rm cm^{2}\,g^{-1}$. The ejecta velocity is fixed to be
7,000 $\rm km\,s^{-1}$. The best-fit model is shown in the upper-right panel
and the bottom panel of Figure 10, and the best-fit parameters are presented
in Table 1 (the CSM+SBO+RAD model). Both the amount of $\rm{}^{56}Ni$
($\sim$0.02 $\rm M_{\odot}$) and the ejecta mass ($\sim$$1.2\rm M_{\odot}$)
are lower than those of SESNe (Lyman et al., 2016). The low ejecta mass
implies that the progenitor of SN 2023fyq is less massive than those of normal
SESNe right before the SN explosion. One caveat of the model is that we did
not consider the CSM interaction at late phase, so the $\rm{}^{56}Ni$ mass we
derive here can only be treated as an upper limit.

The radius of the extended material ($R_{e}$) is around $21\times 10^{13}\rm
cm$ ($\sim$3000 $\rm R_{\odot}$). This large radius is consistent with the
blackbody radius of SN 2023fyq around the shock breakout (Figure 5). This
indicates that, at the explosion, the progenitor is surrounded by an extended
envelope with a mass of 0.3 $\rm M_{\odot}$ at a radius of
$R_{e}$$\sim$3000$\rm R_{\odot}$, consistent with what we discussed in Section
6.1. Considering the width of the narrow line component in the SN spectra
(Figure 5) and the narrow lines observed pre-explosion (Brennan et al., 2024),
the extended material likely expands with a velocity of $\sim$1000 $\rm
km\,s^{-1}$. Such a velocity suggests that the material at around
$\sim$3000$~{}\rm R_{\odot}$ was formed within around 20 days before the
explosion.

In such a scenario, the pre-explosion photophere would be located within the
extended material where the optical depth is sufficiently high. For a wind
profile $\rho\propto r^{-2}$, $R_{BB}$ is roughly proportional to
$\dot{M}/V_{wind}$, where $\dot{M}$ is the mass-loss rate and $V_{wind}$ is
the expansion velocity of the extended material. Consequently the expansion of
$R_{BB}$, starting from around $-$100 d (Figure 5), is likely due to an
increase of mass loss. The more pronounced rise between $\sim-$40 d and $-$11
d can be attributed to a more eruptive mass loss immediately preceding the
explosion. If the majority of the material characterized by $M_{e}$ is formed
during this eruptive phase, the mass loss rate can be estimated to be

$\dot{M}\approx\frac{M_{e}V_{wind}}{R_{e}}\approx
4.5~{}\rm{M_{\odot}\,yr^{-1}}\frac{M_{e}}{0.3\rm M_{\odot}}\frac{3000\rm
R_{\odot}}{R_{e}}\frac{V_{wind}}{1000\rm km\,s^{-1}}.$ (1)

Interestingly, eruptive mass ejections on the order $\sim$0.1–1 $\rm
M_{\odot}$ are anticipated for low-mass He stars with masses of 2.5–3.2 $\rm
M_{\odot}$ due to core silicon deflagration or detonation weeks prior to core
collapse (Woosley, 2019; Ertl et al., 2020). The mass and velocity of the
ejected material depends on the amount of silicon that burns (Ertl et al.,
2020). An ejection mass of $\sim$0.3 $\rm M_{\odot}$ with a velocity of
$\sim$1000$\rm\ km\,s^{-1}$ is consistent with the typical values of such
events (see figure 14 and table 4 of Woosley 2019).

The CSM characterized by $M_{CSM}$ is likely more extended and formed during
the earlier phase of the precursor activities. Detailed discussion on this
topic are provided in Section 6.3.2.

Shortly after the peak, the spectra of SN 2023fyq exhibit broad absorption
lines from the SN ejecta, indicating an optically thin CSM interaction region
between the observer and the SN ejecta. However, the model fit indicates that
the light curve is still predominantly influenced by the CSM interaction. One
possible explanation for this discrepancy is that our analytical model is
oversimplified, leading to an overestimation of the contribution from the CSM
interaction. Alternatively, the CSM may not be spherically symmetric. For
instance, if the SN were surrounded by a disk/torus-like CSM, strong CSM
interaction would mainly occur in the equatorial region. Consequently, an
observer looking along the polar direction would observe less obscured signals
from the SN ejecta while the majority of the luminosity arises from the CSM
interaction. The physical picture of this disk-like CSM scenario has been
extensively discussed in Smith (2017).

In summary, neither radioactive decay nor CSM interaction alone can be the
power source of SN 2023fyq. Approximately a few weeks before the explosion,
about 0.3 $M_{\odot}$ of material is ejected with a velocity of $\sim$1000
$\rm km\,s^{-1}$ due to an increase in mass loss from the progenitor. This
material expands to a radius of $\sim$3000 $\rm R_{\odot}$ at the time of the
explosion. After the explosion, the energy deposited by the shock breakout
from the extended material produces the initial light curve bump. Around 0 d
the light curve is at least partiall powered by the interaction between the SN
ejecta and the surrounding CSM, with the kinetic energy of the ejecta
converted into thermal energy, resulting in a bright peak. After that, as the
strength of the CSM interaction decreases over time, the light curve becomes
more influenced by radioactive decay, leading to a relatively flat light
curve.

Table 1: Best-fit parameters of the CSM+Shock Breakout model and the CSM+Shock Breakout+RAD model.  Model | $t_{exp}$ | $M_{ej}$ | $R_{0}$ | $M_{csm}$ | $\rho_{csm,0}$ | $\epsilon$ | $R_{e}$ | $M_{e}$ | $M_{Ni}$
---|---|---|---|---|---|---|---|---|---
| (Day) | ($\rm M_{\odot}$) | ($\rm 10^{13}\,cm$) | ($\rm M_{\odot}$) | ($\rm log10(g\,cm^{-3})$ |  | ($\rm 10^{13}\,cm$) | ($\rm M_{\odot}$) | ($\rm M_{\odot}$)
CSM+SBO | $-11.5^{+0.1}_{-0.1}$ | $1.3^{+0.1}_{-0.1}$ | $16.0^{+14.2}_{-9.7}$ | $0.7^{+0.1}_{-0.1}$ | $-11.9^{+0.1}_{-0.1}$ | $5^{+0.1}_{-0.1}\times 10^{-2}$ | $24.2^{+0.6}_{-1.1}$ | $0.4^{+0.1}_{-0.1}$ | $\cdots$
CSM+SBO+RAD | $-11.4^{+0.1}_{-0.1}$ | $1.2^{+0.1}_{-0.1}$ | $15.0^{+12.5}_{-10.0}$ | $0.6^{+0.1}_{-0.1}$ | $-12.2^{+0.1}_{-0.1}$ | $5^{+0.1}_{-0.1}\times 10^{-2}$ | $21.4^{+0.7}_{-0.6}$ | $0.3^{+0.1}_{-0.1}$ | $0.02^{+0.01}_{-0.01}$

###  6.3 What Powers The Precursor of SN 2023fyq?

####  6.3.1 Single Massive Star Activities?

SN Precursors have been commonly observed in Type IIn SNe (e.g., Mauerhan et
al., 2013; Smith et al., 2010; Ofek et al., 2013, 2014; Tartaglia et al.,
2016; Pastorello et al., 2013, 2018; Strotjohann et al., 2021; Hiramatsu et
al., 2024), but are rarely found in Type Ibn SNe and Type II SNe. To date, the
pre-explosion activities for Type Ibn SNe have only been detected in SN 2006jc
(Pastorello et al., 2007) and SN 2019uo (Strotjohann et al., 2021). Searches
for precursors in other SNe Ibn yielded only upper limits, ranging from around
$-15$ to $-13$ mag (e.g., Pastorello et al., 2008; Shivvers et al., 2017;
Wangq et al., 2024). This may be because those SNe Ibn had no precursors or
only fainter and shorter ones, and also because most of these events occur at
greater distances than SN 2023fyq. Compared to SN 2006jc and SN 2019uo, one
unique characteristic of SN 2023fyq is the long-standing precursor emission.
Precursor emission observed in SN 2006jc and SN 2019uo is around hundreds of
days before the SN explosions with duration of $\sim$10 days. The precursor
observed in these events are much shorter and brighter than that in SN 2023fyq
(see Figure 2).

We first consider the possibility that the precursor of SN 2023fyq is produced
by the final-stage stellar activities of a single massive star. In this case,
the precursor can be powered by mass ejection driven by wave transport during
the late-stage nuclear burning in the core (Quataert & Shiode, 2012; Shiode &
Quataert, 2014; Fuller, 2017; Fuller & Ro, 2018; Morozova et al., 2020) or
pulsational pair instability (Yoshida et al., 2016; Woosley, 2017).

Massive stars with He core masses of 30 – 64 $\rm M_{\odot}$ experience
pulsational pair instability after carbon burning, producing violent mass
ejections before their cores collapse (Woosley, 2017). Pulsational pair
instability in massive stars have been suggested to be a promising channel of
Type Ibn SNe (Yoshida et al., 2016; Woosley, 2017; Leung et al., 2019; Renzo
et al., 2020). The pulsing activities can last for hours to 10,000 years,
depending on the He core mass, before the SN explosion (Yoshida et al., 2016;
Woosley, 2017). In SN 2023fyq, precursor emission is detected for $\sim$3
years before the SN explosion. Therefore, if pulsational pair instability
powers the precursor emission of SN 2023fyq, the progenitor would be a He star
with a ZAMS mass larger than $\sim$52 $\rm M_{\odot}$ (Woosley, 2017).
However, the outbursts caused by the pulses of these more massive stars are
usually energetic and can result in sharply rising light curves, which is
inconsistent with the relatively steady precursor emission of SN 2023fyq.
Additionally, the low ejecta mass we derived in Section 6.2 does not align
with a very massive He star progenitor. Therefore, we disfavor pulsational
pair instability as the powering mechanism of precursor emission in SN
2023fyq.

Strong temperature gradients can form during late-stage nuclear burning in
massive stars, which generates convection, exciting internal gravity waves.
The gravity waves may carry their energy to the envelope of the star and
deposit it there (Quataert & Shiode, 2012; Shiode & Quataert, 2014; Fuller,
2017; Fuller & Ro, 2018), which may trigger eruptive mass ejections (Leung &
Fuller, 2020; Matzner & Ro, 2021). The mass ejection itself and the collision
between the ejecta generated from multiple outbursts can potentially produce
SN precursor emission (Leung & Fuller, 2020; Strotjohann et al., 2021; Tsuna
et al., 2023). However, it would be difficult to reproduce the time scale of
the observed precursor with a single event of dynamical envelope ejection from
a stripped star (Tsuna et al., 2024). This is because the timescale is
regulated by radiative diffusion from the precursor ejecta, which is only
weeks to months for stripped stars, thus it would work for the precursors of
SN 2006jc or SN 2019uo (Tsuna et al., 2024), but not for SN 2023fyq. In order
to produce the precursor emission seen in SN 2023fyq, multiple fine-tuned mass
ejections would be needed. Therefore, a more plausible scenario is a
continuous mass loss over the timescale of years, with some continuous
powering mechanism for the precursor.

####  6.3.2 Binary Interaction?

A low-mass He star in a binary system has been proposed to be a possible
progenitor scenario for Type Ibn SNe (Maund et al., 2016; Dessart et al.,
2022; Tsuna et al., 2024), which is supported by the lack of star formation at
the site of some members of the class (Sanders et al., 2013; Hosseinzadeh et
al., 2019). In this section we explore the possibility that the progenitor of
SN 2023fyq is an exploded stripped star, such as a He star, in a binary system
and that the binary mass transfer generated the precursor activities.

The stripped SN progenitor in a binary system expands at some point in its
evolution near core-collapse, filling its Roche lobe and initiating mass
transfer onto the companion. Such a scenario is expected for stripped stars
with He core masses in the range of 2.5–3 $M_{\odot}$, which can inflate their
envelopes to up to $\sim 100\ R_{\odot}$ at the oxygen/neon burning phase in
the final years to decades of their lives (e.g., Wu & Fuller, 2022, and
references therein). Thus for orbital separations of $\sim$(1–few) $\times
100\ R_{\odot}$ (orbital period of order 100 days for a companion of order
$\sim 1M_{\odot}$), we expect intense mass transfer to initiate during this
time period.

If the accretor is a compact object, the mass transfer rate is typically
orders of magnitude higher than its Eddington rate, $\dot{M}_{\rm Edd}\sim
2\times 10^{-8}\ {\rm M_{\odot}\ yr^{-1}}(M_{\rm
comp}/{1M_{\odot}})(\kappa_{opt}/{0.1\ {\rm cm^{2}\ g^{-1}}})^{-1}$ (where a
radiation efficiency of $10\%$ was assumed), and thus most of the transferred
mass actually escapes from the binary system without being accreted onto the
compact object. Even if the companion is not a compact object, for large mass
transfer rates of $\gtrsim 10^{-4}$–$10^{-3}\ M_{\odot}$ yr-1, most of the
mass is expected to still escape through the binary’s outer Lagrange point (Lu
et al., 2023). In either case, this escaped material becomes the CSM that
later powers the bright SN.

In Section 6.2 we found that the CSM required to power the main SN light curve
is around $0.6^{+0.1}_{-0.1}$ $\rm M_{\odot}$, which requires a time-averaged
mass loss rate of around a few $0.1M_{\odot}$ yr-1 given that the mass loss is
linked to the observed 1000-day precursor.

For binary systems exhibiting such high mass loss rates suggested by Wu &
Fuller (2022), those with orbital periods ranging from 10 to 100 days are
favored. These systems have orbital velocities of $\sim$100 – a few $100~{}\rm
km\,s^{-1}$. Assuming the velocity of the CSM that escapes the binary system
is $\sim$200 $\rm km\,s^{-1}$, the mass loss rate via mass transfer should be
at least larger than $\sim$2$\times 10^{-2}$ $M_{\odot}$ yr-1 to power the
light curve peak (the detailed derivation is shown in Appendix A), which is
consistent with what we found in Section 6.2.

Given the required $\dot{M}$, we can consider two mechanisms to power the
precursor emission. The first is a collision of the mass-transfer outflow with
external material, which may exist due to a previous mass-transfer episode
(e.g., Pejcha et al., 2016; Metzger & Pejcha, 2017). While we remain agnostic
to the origin of the pre-existing matter, the maximum available power is given
by the kinetic luminosity of the outflow as

$\displaystyle L_{\rm out}\approx\frac{1}{2}\dot{M}v_{\rm CSM}^{2}\sim
1.3\times 10^{39}\ {\rm erg\ s^{-1}}$
$\displaystyle\times\left(\frac{\dot{M}}{0.1M_{\odot}\ {\rm
yr}^{-1}}\right)\left(\frac{v_{\rm CSM}}{200\ {\rm km\ s^{-1}}}\right)^{2}.$
(2)

Thus the precursor may be explained, but only for favorably high CSM velocity
as well as high efficiencies for dissipation and radiation conversion close to
unity.

In the case for a compact object companion, an accretion disk forming around
the compact object can be a promising energy source. While most of the
transferred mass is removed from the outer L2 point, a small fraction can
still accrete onto the companion and form a disk. The disk, if its accretion
rate is super-Eddington, can launch a fast radiation-driven wind that can
collide with the rest of the mass and dissipate its kinetic energy.

The hydrodynamics of the transferred mass has been considered recently in Lu
et al. (2023). For a neutron star companion with an orbital separation of
$a\approx$ (1–few)$\times 100\ R_{\odot}$ and mass transfer rate $\gg 10^{-3}\
M_{\odot}$ yr-1, most of the mass is indeed lost from the L2 point (their
$f_{\rm L2}\sim 1$). However the accretion rate can still reach $\dot{M}_{\rm
acc}\sim(3$–$7)\times 10^{-4}\ M_{\odot}$ yr-1 (Figure A2 of Lu et al. 2023),
which is orders of magnitude larger than the Eddington rate.

For a binary mass ratio of $q=M_{\rm NS}/M_{\rm*}\approx 0.5$, the (Keplerian)
circularization radius of the disk is found from the fitting formula in Lu et
al. (2023) as

$R_{c}\approx 0.10a\sim 7\times 10^{11}\ {\rm
cm}\left(\frac{a}{100R_{\odot}}\right).$ (3)

We expect a disk wind to be launched roughly where the local luminosity
exceeds the Eddington luminosity of the NS, within a disk radius (equation 31
of Lu et al. 2023)

$\displaystyle R_{\rm sph}$ $\displaystyle\approx$
$\displaystyle\frac{\dot{M}_{\rm acc}\kappa}{4\pi c}$ (4) $\displaystyle\sim$
$\displaystyle 2\times 10^{10}\ {\rm cm}\left(\frac{\dot{M}_{\rm acc}}{5\times
10^{-4}M_{\odot}\ {\rm yr^{-1}}}\right)\left(\frac{\kappa}{0.2\ {\rm cm^{2}\
g^{-1}}}\right),$

which is typically less than $R_{\rm c}$ for an orbital separation of $a\sim
100~{}R_{\odot}$. We have taken the opacity here to be $\kappa\approx 0.2\
{\rm cm^{2}\ g^{-1}}$ as helium is expected to be fully ionized in the
interior of the disk. In line with many theoretical works that model super-
Eddington disk winds, we assume a power-law accretion rate $\dot{M}$ of
$\dot{M}\propto r^{p}$ ($R_{\rm NS}<r<R_{\rm sph}$), where we adopt $R_{\rm
NS}=10$ km. This means that a fraction of the accreted mass is expelled at
each radius, and we assume that the wind velocity is equivalent to the local
disk escape velocity. Consequently, the wind kinetic luminosity, integrated
over the range of $r$, is estimated as

$\displaystyle L_{\rm wind}$ $\displaystyle\approx\frac{p}{2(1-p)}\dot{M}_{\rm
acc}\frac{GM_{\rm NS}}{R_{\rm NS}}\left(\frac{R_{\rm NS}}{R_{\rm
sph}}\right)^{p}$ $\displaystyle\sim 2\times 10^{40}\ {\rm erg\
s^{-1}}\left(\frac{\dot{M}_{\rm acc}}{5\times 10^{-4}M_{\odot}\ {\rm
yr^{-1}}}\right)^{1/2}$ $\displaystyle\times\left(\frac{M_{\rm
NS}}{1.4M_{\odot}}\right)\left(\frac{\kappa}{0.2\ {\rm cm^{2}\
g^{-1}}}\right)^{-1/2}$ (5)

where we have adopted $p=0.5$ in the last equation while a possible range of
$0.3\leq p\leq 0.8$ is suggested (Yuan & Narayan, 2014). We thus find that the
disk wind carries the appropriate kinetic luminosity to explain the precursor
in the steady-state phase.

As the disk wind carries much smaller mass than the rest of the material
around the system, its kinetic energy will be efficiently dissipated by their
collision. We check that the dissipated energy would be successfully radiated
as the precursor. For a wind profile the diffusion timescale in the CSM is

$\displaystyle t_{\rm diff}$ $\displaystyle\approx\frac{\kappa\dot{M}}{4\pi
v_{\rm CSM}c}$ $\displaystyle\sim 8\times 10^{4}\ {\rm
sec}\left(\frac{\dot{M}}{0.1M_{\odot}\ {\rm
yr}^{-1}}\right)\left(\frac{\kappa}{0.1\ {\rm cm^{2}\
g^{-1}}}\right)\left(\frac{v_{\rm CSM}}{200\ {\rm km\ s^{-1}}}\right)^{-1}$
(6)

and the adiabatic expansion timescale from the dissipation region, whose size
is roughly comparable to the orbital separation, is

$t_{\rm exp}\approx\frac{a}{v_{\rm CSM}}\sim 3\times 10^{5}\ {\rm
sec}\left(\frac{a}{100\ R_{\odot}}\right)\left(\frac{v_{\rm CSM}}{200\ {\rm
km\ s^{-1}}}\right)^{-1}$ (7)

Thus we expect that the dissipated energy can be successfully radiated away
without adiabatic losses. The radiation will be reprocessed in the CSM, and
finally be emitted as optical radiation at $r\approx R_{\rm BB}$. The mass
loss via the L2 point can form an equatorial disk (e.g., Lu et al., 2023). The
interaction of the equatorial disk with the SN ejecta may contribute to the
main peak of the SN light curve. In this case, the parameter $M_{CSM}$
mentioned in Section 6.2 roughly characterizes the mass of the equatorial
disk. The interaction of SN ejecta with this dense CSM may still continue in
the nebular phase, producing the intermediate-width He lines we observe.

####  6.3.3 What About The Rise After $-$100 d In The Pre-explosion Light
Curve?

As we mentioned in Section 4.3, the pre-explosion light curve shows a rapid
rise after $-$100 d, with a more pronounced rise occurring between $-$40 d and
$-$11 d. This may be associated with eruptive mass loss right before the SN
explosion. For the more pronounced rise between $-$40 d and $-$11 d, we
consider two possibilities: 1) the rise is due to orbital shrinking of the
binary, leading to a runaway of mass transfer and resulting in a rapid-rising
pre-explosion light curve (i.e., MacLeod et al., 2018). 2) The rise is
influenced by the core silicon burning of the He star, which ejects a large
amount of material and powers the fast-rising light curve just before the core
collapses.

For the first case we initially consider the orbital evolution of this binary
system over the few-year timescale during which we observe the precursor. The
mass loss from the Lagrange point carries away angular momentum as well, which
can affect the orbital separation of the binary. This generally leads to
shrinking of the orbit, which may have been witnessed as the sharp rise of the
light curve as we approach the explosion epoch. From Figure 5 of Lu et al.
(2023) we find the orbital shrinking rate for mass ratio $q=0.5$ and $f_{\rm
L2}=1$ as

$\frac{\dot{a}}{a}\approx(-5)\frac{\dot{M}}{M_{*}}\sim-(6\ {\rm
yr})^{-1}\left(\frac{\dot{M}}{0.1M_{\odot}\ {\rm
yr}^{-1}}\right)\left(\frac{M_{*}}{3M_{\odot}}\right)^{-1}$ (8)

which means that for a mass loss rate of $\sim 0.1M_{\odot}$ yr-1, the orbital
separation can significantly shrink in the several years that we observe the
precursor. The orbital shrinking of the binary may cause an unstable mass
transfer and accretion onto the compact object, resulting in a runaway mass
loss. This may explain the rapid rise after around $-$40 d in the precursor
light curve. Given the anticipated significant orbital shrinking within
several years for the system under consideration, the shallower rise in the
light curve between $-$100 d and $\sim-40$ d is likely also influenced by the
orbital shrinking. This may only lead to a gently increase in the accretion
rate onto the compact companion, resulting in the rise of the light curve.

In this scenario the final SN explosion can be due to the merger of the He
star with a compact object (e.g., Chevalier, 2012; Soker, 2019; Metzger,
2022). Such merger-driven explosions have been proposed to explain some long
gamma-ray bursts (Fryer & Woosley, 1998; Zhang & Fryer, 2001; Thöne et al.,
2011; Fryer et al., 2013), which are usually associated with a subtype of Type
Ic SNe that exhibit broad spectral lines. This He-merger scenario can connect
the observed rapid increase in the light curve’s brightness at the end of the
precursor phase with the following SN-like explosion. However, the
characteristics of the final explosion post-merger remain poorly understood.
For example, the predicted explosion energies are uncertain by many orders of
magnitude (Fryer & Woosley, 1998; Zhang & Fryer, 2001; Schrøder et al., 2020).
While the merger-driven explosion might explain the spectral features
observed, detailed spectral modeling of these events is still lacking.

For the second case, a core-collapse SN explosion is anticipated after
significant mass transfer over years from low-mass stripped stars ranging from
$2.5$ to $3M_{\odot}$ (Wu & Fuller, 2022). Additionally an explosive mass
ejection weeks before the explosion due to silicon burning is indeed expected
in recent studies for low mass He stars with masses of 2.5 – 3.2 $\rm
M_{\odot}$ (Woosley, 2019). The mass ejected can range from $10^{-2}$ to 1
$\rm M_{\odot}$ with velocities from $\sim$100 $\rm km\,s^{-1}$ to a few 1000
$\rm km\,s^{-1}$. In Section 6.2 we found that there is likely an eruptive
mass loss of $\sim$0.3 $\rm M_{\odot}$ a few weeks before the SN explosion
with a velocity of $\sim$1000 $\rm km\,s^{-1}$, which is consistent with the
silicon burning phase for low-mass He stars. The eruptive mass loss may
explain the more pronounced rise of the precursor light curve between
$\sim-$40 d and $-$11 d, and the ejected material in turn produces the first
SN peak. However, we note that detailed light curve modeling is necessary to
confirm this hypothesis. In this case, the shallower rise in the light curve
between $-$100 d and $\sim-40$ d is likely still attributed to the orbital
shirking of the binary system, like discussed above.

In this scenario the final SN explosion results from the core collapse of the
He star. This explanation accounts for the observed spectral similarities
between SN 2023fyq and SESNe both post-peak and during the nebular phases.

Both the merger-driven and core-collapse scenarios can account for certain
observed features of SN 2023fyq. In either case, the progenitor system would
likely be asymmetric, which aligns with observations of SN 2023fyq. The 56Ni
yields from a merger-driven explosion are likely low (Fryer et al., 2013;
Metzger, 2022) and, similarly, low 56Ni production is expected from core-
collapse explosions in low-mass helium stars (Woosley, 2019). These
predictions are consistent with the low 56Ni mass derived from the late-time
light curves of SN 2023fyq.

An important difference between these two scenarios is that a merger-driven
explosion typically results in a single compact object in the remnant, whereas
a core-collapse explosion generally leaves behind a compact binary. In the
latter case fallback accretion post-explosion could produce observable X-ray
emissions approximately 100 to 1000 days after the explosion, which may show
time variations tied to the orbital motion of the binary (Kashiyama et al.,
2022). For SN 2023fyq, conducting X-ray follow-up years after the explosion
could be helpful in distinguishing between these two scenarios in future
studies (see Appendix B for details).

In conclusion the timescale and brightness of the precursor observed in SN
2023fyq before $-$100 d can be attributed to mass transfer in a binary system.
The companion star is likely a compact object, as the energetics of the disk
wind launched from super-Eddington accretion onto the compact object can
naturally explain the luminosity of the precursor. An equatorial circumbinary
disk, formed during the mass transfer, later interacts with the SN ejecta,
powering the main SN peak. During the nebular phases the ongoing interaction
between the equatorial disk and the SN ejecta produces the intermediate-width
He lines observed. The rise of the light curve between $-$100 d and $\sim-40$
d is likely due to orbital shrinking. The more pronounced rise of the light
curve starting around $-$40 d may be linked to 1) an eruptive mass ejection
due to final-stage silicon burning, or 2) runaway mass transfer caused by
orbital shrinking of the binary system. In the first scenario, the subsequent
explosion would result from the core-collapse of the He star. In the second
scenario, it would result from the merger of the He star with the compact
object. Both scenarios can launch materials into the polar region. The shock
breakout from this extended material and the following cooling emission power
the first bright SN peak.

###  6.4 Connections to Other Transient Phenomena and Implications on The CSM
Structure

It is noteworthy that the light curve morphology (both the pre- and post-
explosion phase) of SN 2023fyq is quite similar to those of luminous red novae
(Soker & Tylenda, 2003; Tylenda et al., 2011; Mauerhan et al., 2015; Smith et
al., 2016; Blagorodnova et al., 2017), which are generally understood to be
the product of binary mergers (e.g., Metzger & Pejcha, 2017). The pre-
explosion activities in luminous red novae are often associated with binary
mass transfer (e.g., Pejcha, 2014), and the pre-explosion brightening is due
to the increase in the mass-loss rate caused by orbital shrinking. The post-
explosion light curves of luminous red novae are double-peaked, in which the
first peak is likely from the shock cooling and the second peak is from the
interaction between the ejecta and a pre-existing equatorial disc formed
during binary mass transfer (Metzger & Pejcha, 2017).

The scenario for luminous red novae is analogous to what we proposed for SN
2023fyq, and the primary difference is just the explosion energy source. Such
an asymmetric CSM structure is consistent with the multi-component profile of
the He I $\lambda$5876 line as we discussed in Section 5 and also the
asymmetric line profiles observed during the pre-explosion phase of SN 2023fyq
(Brennan et al., 2024). Similarities between luminous red novae and
interaction-powered SNe have also been reported in previous studies (e.g.,
Hiramatsu et al., 2024).

The SN light curve evolution of SN 2023fyq is similar to those of ultra-
stripped SNe (De et al., 2018; Yao et al., 2020). The first bright SN light
curve peak in these ultra-stripped SNe is generally understood as a result of
shock breakout from the dense CSM ejected weeks before the SN explosion. The
second peak of these objects is usually around $10^{42}\rm erg\,s^{-1}$, much
fainter than that of SN 2023fyq, and is thought to be powered by
$\rm{}^{56}Ni$ decay (De et al., 2018; Yao et al., 2020). It may be that in
these objects the CSM is more confined and a more extended ($\sim 10^{15}$ cm)
dense equatorial disk is lacking, resulting in insufficient CSM at these radii
to power the second peak through interaction like that observed in SN 2023fyq.

SNe Ibn can show a wide variety of spectral features at early phases
(Hosseinzadeh et al., 2017), which is not surprising if all SNe Ibn experience
strong interaction with asymmetric CSM (e.g., Smith et al., 2015; Smith,
2017). Only a few SNe Ibn are observed until late phases since they can
decline fast. Interestingly, as we show in Figure 9, at late times, these SNe
Ibn seem to fall into two distinct classes: Class I that shows broad lines and
share many similarities with normal SESNe (SN 2023fyq, SN 2015G, SN 2018gjx)
and Class II that is still dominated by narrow emission lines (SN 2006jc, SN
2019kbj). Assuming the progenitors of all these SNe Ibn are He stars, the
objects in Class II may be surrounded by more massive CSM and/or have lower
explosion energy (Dessart et al., 2022).

For the objects in Class I, the intensity of the [O I] $\lambda\lambda$6300,
6364 line can vary significantly among different objects while the other
spectral features are quite similar. If the progenitors of all these objects
are surrounded by an equatorial disk, the difference in the intensity of the
[O I] $\lambda\lambda$6300, 6364 line can be naturally explained by different
viewing angles (See Figure 11). If the system is observed from the equatorial
direction, the central [O I] $\lambda\lambda$6300, 6364 line forming region
can be obscured by the disk. Instead, a polar observer would be able to see
the whole nebular emission from the inner ejecta. For both observers,
intermediate-width He emission lines from the ongoing interaction of the SN
ejecta with the equatorial disk can be seen.

A disk/torus-like CSM is also invoked in previous studies to explain the
spectroscopic evolution of SNe Ibn (Prentice et al., 2020) and SNe IIn (e.g.,
Smith & Arnett, 2014; Smith et al., 2015; Andrews & Smith, 2018; Smith &
Andrews, 2020). Such a disk/torus-like CSM scenario could potentially explain
the diversity we see in SNe Ibn in Class I, and is consistent with the
precursor model we discussed in Section 6.3.2. This suggests that Class I SNe
Ibn may originate from a similar progenitor channel but with variations in
viewing angles.

Long-lasting and relatively stable precursor activities due to binary
interaction are commonly seen in luminous red novae (e.g., Tylenda et al.,
2011; Mauerhan et al., 2015; Blagorodnova et al., 2017). Given the similarity
of the progenitor scenario of luminous red novae and SN 2023fyq, it is
possible that precursor activities are not rare in SNe Ibn in Class I. If this
is true, the long-lasting and slowly rising pre-explosion emission may serve
as a unique early warning for this subclass of Type Ibn SNe. The evolution of
the precursor light curves may vary depending on the viewing angle, as the
emission could be obscured by the equatorial disk for observers near the
equatorial plane. Given that the viewing angle also influences the intensity
of the [OI] lines in the nebular spectra, combining the precursor emission
with late-time spectroscopy could serve as a unique probe for the progenitor
scenario we propose.

Figure 11: A sketch of the possible progenitor system of SN 2023fyq. Upper:
around a few years before the explosion, the progenitor (a He star with a mass
of $\sim$2.5 – 3 $M_{\odot}$) expands at the oxygen/neon burning phase,
filling its Roche lobe. This triggers mass transfer onto its companion compact
object, resulting in the precursor emission we observe. Around weeks before
the explosion, an eruptive mass ejection is triggered through core silicon
burning in the low-mass He star or runaway mass transfer due to orbital
shrinking, launching dense material to the polar region. The subsequent
explosion is likely due to either by core-collapse of the He star or by the
merger of the He star with its compact object companion. Bottom: Immediately
after the explosion, the shock breaks out from the dense polar material formed
weeks before the explosion, producing the first light curve peak. The
interaction of SN ejecta with the equatorial disk formed by the pre-explosion
binary interaction contributes to the second peak.

##  7 Summary

The evolution of SN 2023fyq closely resemble that of Type Ibn SNe. The optical
spectra post-peak and the nebular spectrum of SN 2023fyq share similarities
with those of normal SESNe, implying that the progenitor is a stripped/He
star. The SN light curve can be reproduced by a CSM interaction + shock
breakout + 56Ni decay model, implying the presence of dense CSM around the
progenitor, a low progenitor mass and a low $\rm{}^{56}Ni$ production. The
precursor emission of SN 2023fyq is observed up to around three years before
the SN explosion. The long-duration precursor activities are best explained by
the mass transfer in a binary system involving a low-mass He star.

Putting all these together, we summarize a possible timeline for SN 2023fyq:

  1. 1.

$\sim$$-$1000 d to $\sim$$-$100 d (upper panel of Figure 11): A low-mass He
star (2.5 – 3 $\rm M_{\odot}$) expands substantially at the oxygen/neon
burning phase, triggering mass transfer to its companion compact object, which
produces the precursor emission we observe. The outflow via L2 point produces
the He-rich CSM around the progenitor system and forms an equatorial disk
($\sim$0.6$\rm M_{\odot}$).

  2. 2.

$\sim$$-$100 d to $\sim$$-$11 d: The shrinkage of the orbit leads to an
increase in the accretion rate onto the companion compact object, resulting in
a rise in the light curve. The more pronounced light curve rise after
$\sim-$40 d is likely due to either the core silicon burning or the runaway
mass transfer caused by orbital shrinking, which triggers an eruptive mass
ejection ($\sim$0.3$\rm M_{\odot}$) with a velocity of $\sim$1000$\rm
km\,s^{-1}$. This launches dense material to the polar region.

  3. 3.

$\sim$$-$11 d (bottom panel of Figure 11): A SN explosion is triggered either
by the core-collapse of the He star or by the merger of the He star with a
compact object, which sends a shock through the polar material ($\sim$3000
$\rm R_{\odot}$). The energy deposited during the shock breakout produces the
initial bump of the light curve.

  4. 4.

$\sim$$-$11 d to $\sim$20 d: The SN ejecta collide with the equatorial He-rich
CSM ($\sim$0.6$\rm M_{\odot}$), converting the kinetic energy of the SN ejecta
into thermal energy, contributing to the SN light curve and generating a very
blue spectrum with only prominent He lines. With the expansion of the ejecta,
the optical depth decreases so that more signals from the SN ejecta are
observed.

  5. 5.

after $\sim$20 d: The strength of the CSM interaction decreases and the SN
fades, and radioactive decay likely starts to contribute more to the light
curve. Later, the ejecta become more optically thin and the object transitions
into the nebular phase. Given our proximity to the polar direction of the
system, signals from the inner part of the ejecta are revealed, which closely
resemble those of normal SESNe at nebular phases. Additionally, the continuing
interaction between the ejecta and the He-rich equatorial CSM produces strong
intermediate-width He emission lines.

Given the similarities between SN 2023fyq and other Type Ibn SNe, precursor
activities may be common for a certain subclass of Type Ibn SNe. If an
equatorial disk is indeed formed during the precursor phase, the precursor
emission and the intensity of the [OI] lines at the nebular phases for this
class of objects would be dependent on the viewing angle. It is worth noting
that this mechanism does not apply to the very brief, singular pre-explosion
outburst observed in SN 2006jc and SN 2019uo. For the upcoming LSST survey, a
single 30-second visit will achieve a 5$\sigma$ depth of approximately 24 mag
(Bianco et al., 2022). By stacking images, even deeper limits can be achieved.
This enables LSST to effectively constrain the precursors of Type Ibn SNe,
such as SN 2023fyq, within 150 Mpc, assuming a typical precursor brightness of
$-$12 mag. A sample of Type Ibn SNe with well-constrained precursor
activities, combined with the late-time spectroscopy, will test the progenitor
scenario we propose. We also encourage detailed spectral and light curve
modeling of merger-driven explosions, as well as the silicon burning phase in
low-mass He stars just prior to core collapse. By comparing these models with
a large sample of observations, we can deepen our understanding of the final
stages of stellar evolution.

## Acknowledgements

We would like to thank Jim Fuller for the assistance with the manuscript in
its early stages. We would like to thank Kyle Davis for sharing the SOAR
spectrum from his program. Y.D. would like to thank L.Z. for redesigning and
redrawing Figure 11 in the paper.

Research by Y.D., S.V., N.M.R, E.H., and D.M. is supported by NSF grant
AST-2008108. D.T. is supported by the Sherman Fairchild Postdoctoral
Fellowship at the California Institute of Technology.

Time-domain research by the University of Arizona team and D.J.S. is supported
by NSF grants AST-1821987, 1813466, 1908972, 2108032, and 2308181, and by the
Heising-Simons Foundation under grant #2020-1864.

This work makes use of data from the Las Cumbres Observatory global telescope
network. The LCO group is supported by NSF grants AST-1911225 and AST-1911151.

A.Z.B. acknowledges support from the European Research Council (ERC) under the
European Union’s Horizon 2020 research and innovation program (grant agreement
No. 772086).

This publication was made possible through the support of an LSST-DA Catalyst
Fellowship to K.A.B, funded through Grant 62192 from the John Templeton
Foundation to LSST Discovery Alliance. The opinions expressed in this
publication are those of the authors and do not necessarily reflect the views
of LSST-DA or the John Templeton Foundation.

Based on observations obtained at the international Gemini Observatory, a
program of NSF’s NOIRLab, which is managed by the Association of Universities
for Research in Astronomy (AURA) under a cooperative agreement with the
National Science Foundation. On behalf of the Gemini Observatory partnership:
the 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).

This work was enabled by observations made from the Gemini North telescope,
located within the Maunakea Science Reserve and adjacent to the summit of
Maunakea. We are grateful for the privilege of observing the Universe from a
place that is unique in both its astronomical quality and its cultural
significance.

This work includes observations obtained at the Southern Astrophysical
Research (SOAR) telescope, which is a joint project of the Ministério da
Ciência, Tecnologia e Inovações (MCTI/LNA) do Brasil, the US National Science
Foundation’s NOIRLab, the University of North Carolina at Chapel Hill (UNC),
and Michigan State University (MSU).

Some of the data presented herein were obtained at Keck Observatory, which is
a private 501(c)3 non-profit organization operated as a scientific partnership
among the California Institute of Technology, the University of California,
and the National Aeronautics and Space Administration. The Observatory was
made possible by the generous financial support of the W. M. Keck Foundation.
The authors wish to recognize and acknowledge the very significant cultural
role and reverence that the summit of Maunakea has always had within the
indigenous Hawaiian community. We are most fortunate to have the opportunity
to conduct observations from this mountain.

The LBT is an international collaboration among institutions in the United
States, Italy and Germany. LBT Corporation Members are: The University of
Arizona on behalf of the Arizona Board of Regents; Istituto Nazionale di
Astrofisica, Italy; LBT Beteiligungsgesellschaft, Germany, representing the
Max-Planck Society, The Leibniz Institute for Astrophysics Potsdam, and
Heidelberg University; The Ohio State University, and The Research
Corporation, on behalf of The University of Notre Dame, University of
Minnesota and University of Virginia.

This research has made use of the NASA/IPAC Extragalactic Database (NED),
which is funded by the National Aeronautics and Space Administration and
operated by the California Institute of Technology.

This research made use of Photutils, an Astropy package for detection and
photometry of astronomical sources (Bradley et al., 2022).

##  Appendix A The Mass Loss Rate in binary interaction scenario

This appendix calculates the mass loss rate needed for a binary system to
explain the observations, as discussed in Section 6.3.2. We begin with
estimating the required mass loss rate $\dot{M}$ of the CSM, which in our
scenario is equivalent to the mass transfer rate if the rate is much larger
than the Eddington rate and the companion cannot accrete most of the
transferred material. The CSM must be optically thick within the observed
blackbody radius $R_{\rm BB}\approx 600~{}R_{\odot}$ at the precursor phase.
For a mass loss rate of $\dot{M}$, the optical depth at $R_{\rm BB}$ is

$\displaystyle\tau_{\rm CSM}(r=R_{\rm BB})$
$\displaystyle\approx\frac{\kappa\dot{M}}{4\pi R_{\rm BB}v_{\rm CSM}}$
$\displaystyle\sim 60\left(\frac{\dot{M}}{0.1\ M_{\odot}\ {\rm
yr}^{-1}}\right)\left(\frac{\kappa}{0.1\ {\rm cm^{2}\
g^{-1}}}\right)\left(\frac{v_{\rm CSM}}{200\ {\rm km\ s^{-1}}}\right)^{-1}$
(A1)

where $v_{\rm CSM}$ is the velocity of the CSM that escapes the binary system.
This is typically the orbital velocity for outflows from mass transfer, which
is $\sim 200$ km s-1 for the orbital separation of interest (see Section
6.3.2), but the arguments below would not depend much on the adopted value.
The value of $\kappa\approx 0.1\ {\rm cm^{2}\ g^{-1}}$ is motivated from that
of singly-ionized helium at around $10^{4}$ K (e.g., Kleiser & Kasen, 2014).
The optical depth then poses a lower limit in $\dot{M}$ of

$\dot{M}\geq\dot{M}_{\rm min}\approx 2\times 10^{-3}M_{\odot}\ {\rm
yr}^{-1}\left(\frac{\kappa}{0.1\ {\rm cm^{2}\
g^{-1}}}\right)^{-1}\left(\frac{v_{\rm CSM}}{200\ {\rm km\ s^{-1}}}\right)$
(A2)

which confirms the super-Eddington mass transfer rate333As the blackbody
temperature is $\sim 10^{4}$ K during the precursor phase, even for
$\dot{M}\gg\dot{M}_{\rm min}$, we expect that the blackbody radius would not
be too much larger than the observed value. This is because the temperature
drops as a function of radius, and the opacity at $r>R_{\rm BB}$ will rapidly
drop with radius due to helium recombination (analogous to the recombination
front of SN II-P).. As a cross check, we can also roughly infer $\dot{M}$ from
the observed SN. The collision of the SN with the CSM generates a shock that
powers the SN light curve. The kinetic energy dissipation rate is

$\displaystyle L_{\rm kin}$ $\displaystyle=2\pi r^{2}\left(\frac{\dot{M}}{4\pi
r^{2}v_{\rm CSM}}\right)v_{\rm sh}^{3}$ $\displaystyle\sim 5.5\times 10^{43}\
{\rm erg\ s^{-1}}\left(\frac{\dot{M}}{0.1M_{\odot}\ {\rm
yr}^{-1}}\right)\left(\frac{v_{\rm CSM}}{200\ {\rm km\
s^{-1}}}\right)^{-1}\left(\frac{v_{\rm sh}}{7000\ {\rm km\
s^{-1}}}\right)^{3}$ (A3)

where $v_{\rm sh}$ is the forward shock velocity. Assuming that the luminosity
at the second peak is generated by the interaction with CSM generated in the
precursor phase, we infer a mass loss rate of

$\dot{M}\sim 2\times 10^{-2}\ M_{\odot}\ {\rm
yr}^{-1}\epsilon^{-1}\left(\frac{L_{\rm rad}}{10^{43}\ {\rm erg\
s^{-1}}}\right)\left(\frac{v_{\rm CSM}}{200\ {\rm km\
s^{-1}}}\right)\left(\frac{v_{\rm sh}}{7000\ {\rm km\ s^{-1}}}\right)^{-3},$
(A4)

where $\epsilon=L_{\rm rad}/L_{\rm kin}\leq 1$ is the radiation conversion
efficiency. While this estimate is quite sensitive to the assumed $v_{\rm
sh}$, it implies that a similarly high $\dot{M}$ is also required to explain
the SN. The required mass transfer rate of $\sim 0.02$–$0.2M_{\odot}$ yr-1 for
$\epsilon\approx 0.1$–$1$ roughly overlaps with the range obtained from
simulations of binaries composed of a low-mass ($2.5$–$3\ M_{\odot}$) He star
and a neutron star, years to decades before the SN (Wu & Fuller, 2022, Figure
2).

##  Appendix B Late-time X-ray detectability of SN 2023fyq

In this appendix, we roughly estimate the X-ray detectability of SN 2023fyq
for future followup. We expect X-ray emission when it is transparent to
photoionization by oxygen and carbon in the ejecta. Our modeling favors low-
mass (a few $M_{\odot}$) helium stars for the progenitor, with carbon-oxygen
cores of mass $\approx 1.5$–$2~{}M_{\odot}$. For an explosion ejecta from such
progenitors, we infer the mass of carbon/oxygen-rich material to be roughly
$M_{\rm ej,C/O}\sim 0.1$–$1~{}M_{\odot}$. The lower limit applies if a neutron
star is left behind in the explosion (as in ultra-stripped SNe considered in
Kashiyama et al. 2022), and the upper limit is if the bulk of the CO-core is
disrupted (e.g. by a merger) and becomes part of the SN ejecta. Adopting the
ejecta velocity of $v_{\rm ej}=7000$ km s-1 and the X-ray photoionization
cross section of $\sigma_{\rm X}\sim 10^{-19}\ {\rm cm^{2}}(h\nu/{\rm
keV})^{-3}$, we expect X-rays with energy $h\nu$ to be transparent at

$\displaystyle t_{\rm trans}$ $\displaystyle\sim$
$\displaystyle\sqrt{\frac{\sigma_{\rm X}M_{\rm ej,C/O}/14m_{p}}{4\pi v_{\rm
ej}^{2}}}\sim 1~{}{\rm yr}\left(\frac{M_{\rm
ej,C/O}}{0.1~{}M_{\odot}}\right)^{1/2}\left(\frac{h\nu}{5~{}{\rm
keV}}\right)^{-3/2}.$ (B1)

Thus follow-up in hard X-rays at years after the explosion is encouraged,
although the X-ray luminosity would depend on the uncertain degree of
fallback. If the fallback is similar to the ultra-stripped SN models in
Kashiyama et al. (2022), we expect the source to be detectable by current
X-ray facilities thanks to the proximity of this event.

##  Appendix C Spectroscopic Observations

Table A1 shows a log of the spectroscopic observations of SN 2023fyq. Table A2
shows a log of the spectroscoppic observation of SN 2019kbj.

Table A1: Spectroscopic observations of SN 2023fyq UT Date | Julian Date (Days) | Phase (Days) | Telescope | Instrument
---|---|---|---|---
2023-07-27 | 2460152.743 | -1.3 | Gemini | GMOS
2023-07-27 | 2460152.85 | -1.1 | FTS | FLOYDS
2023-07-28 | 2460153.859 | -0.1 | FTS | FLOYDS
2023-07-31 | 2460156.851 | 2.9 | FTS | FLOYDS
2023-08-01 | 2460157.74 | 3.7 | Gemini | GMOS
2023-08-04 | 2460160.858 | 6.9 | FTS | FLOYDS
2023-08-04 | 2460161.392 | 7.4 | NOT | ALFOSC
2023-12-12 | 2460291.123 | 137.1 | Keck | LRIS
2024-01-23 | 2460332.761 | 178.8 | SOAR | GHTS
2024-03-11 | 2460380.865 | 226.9 | LBT | MODS
2024-05-01 | 2460431.943 | 277.9 | Keck | LRIS
Table A2: Spectroscopic observations of SN 2019kbj UT Date | Julian Date (Days) | Phase (Days) | Telescope | Instrument
---|---|---|---|---
2019-09-23 | 2458750.817 | 80 | Keck | DEIMOS

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