A&A, 685, A20 (2024)
https://doi.org/10.1051/0004-6361/202348166
c© The Authors 2024
Astronomy
&Astrophysics
SN2020zbf: A fast-rising hydrogen-poor superluminous supernova
with strong carbon lines
A. Gkini1 , R. Lunnan1 , S. Schulze2 , L. Dessart3 , S. J. Brennan1 , J. Sollerman1 , P. J. Pessi1 , M. Nicholl4 ,
L. Yan5, C. M. B. Omand1 , T. Kangas6,7 , T. Moore4 , J. P. Anderson8,9 , T.-W. Chen10 , E. P. Gonzalez11,12,
M. Gromadzki13 , C. P. Gutiérrez14,15, D. Hiramatsu16,17 , D. A. Howell11,12 , N. Ihanec13, C. Inserra18 ,
C. McCully11 , T. E. Müller-Bravo15,14 , C. Pellegrino11,12 , G. Pignata19 , M. Pursiainen20 , and D. R. Young4
1 The Oskar Klein Centre, Department of Astronomy, Stockholm University, Albanova University Center, 106 91 Stockholm,
Sweden
e-mail: anamaria.gkini@astro.su.se
2 The Oskar Klein Centre, Department of Physics, Stockholm University, Albanova University Center, 106 91 Stockholm, Sweden
3 Institut d’Astrophysique de Paris, CNRS-Sorbonne Université, 98 bis boulevard Arago, 75014 Paris, France
4 Astrophysics Research Centre, School of Mathematics and Physics, Queens University Belfast, Belfast BT7 1NN, UK
5 The Caltech Optical Observatories, California Institute of Technology, Pasadena, CA 91125, USA
6 Finnish Centre for Astronomy with ESO (FINCA), University of Turku, 20014 Turku, Finland
7 Tuorla Observatory, Department of Physics and Astronomy, University of Turku, 20014 Turku, Finland
8 European Southern Observatory, Alonso de Córdova 3107, Casilla 19, Santiago, Chile
9 Millennium Institute of Astrophysics MAS, Nuncio Monsenor Sotero Sanz 100, Of. 104, Providencia, Santiago, Chile
10 Graduate Institute of Astronomy, National Central University, 300 Jhongda Road, 32001 Jhongli, Taiwan
11 Las Cumbres Observatory, 6740 Cortona Dr. Suite 102, Goleta, CA 93117, USA
12 Department of Physics, University of California, Santa Barbara, CA 93106-9530, USA
13 Astronomical Observatory, University of Warsaw, Al. Ujazdowskie 4, 00-478 Warszawa, Poland
14 Institut d’Estudis Espacials de Catalunya (IEEC), Gran Capità, 2-4, Edifici Nexus, Desp. 201, 08034 Barcelona, Spain
15 Institute of Space Sciences (ICE, CSIC), Campus UAB, Carrer de Can Magrans s/n, 08193 Barcelona, Spain
16 Center for Astrophysics, Harvard & Smithsonian, 60 Garden Street, Cambridge, MA 02138-1516, USA
17 The NSF AI Institute for Artificial Intelligence and Fundamental Interactions, MIT, Cambridge, MA 02139, USA
18 Cardiff Hub for Astrophysics Research and Technology, School of Physics & Astronomy, Cardiff University, Queens Buildings,
The Parade, Cardiff CF24 3AA, UK
19 Instituto de Alta Investigación, Universidad de Tarapacá, Casilla 7D, Arica, Chile
20 Department of Physics, University of Warwick, Gibbet Hill Road, Coventry CV4 7AL, UK
Received 5 October 2023 / Accepted 23 January 2024
ABSTRACT
SN 2020zbf is a hydrogen-poor superluminous supernova (SLSN) at z = 0.1947 that shows conspicuous C ii features at early times,
in contrast to the majority of H-poor SLSNe. Its peak magnitude is Mg = −21.2 mag and its rise time (.26.4 days from first light)
places SN 2020zbf among the fastest rising type I SLSNe. We used spectra taken from ultraviolet (UV) to near-infrared wavelengths
to identify spectral features. We paid particular attention to the C ii lines as they present distinctive characteristics when compared to
other events. We also analyzed UV and optical photometric data and modeled the light curves considering three different powering
mechanisms: radioactive decay of 56Ni, magnetar spin-down, and circumstellar medium (CSM) interaction. The spectra of SN 2020zbf
match the model spectra of a C-rich low-mass magnetar-powered supernova model well. This is consistent with our light curve
modeling, which supports a magnetar-powered event with an ejecta mass Mej = 1.5 M. However, we cannot discard the CSM-
interaction model as it may also reproduce the observed features. The interaction with H-poor, carbon-oxygen CSM near peak light
could explain the presence of C ii emission lines. A short plateau in the light curve around 35–45 days after peak, in combination with
the presence of an emission line at 6580 Å, can also be interpreted as being due to a late interaction with an extended H-rich CSM.
Both the magnetar and CSM-interaction models of SN 2020zbf indicate that the progenitor mass at the time of explosion is between
2 and 5 M. Modeling the spectral energy distribution of the host galaxy reveals a host mass of 108.7 M, a star formation rate of
0.24+0.41−0.12 M yr
−1, and a metallicity of ∼0.4Z.
Key words. supernovae: general – supernovae: individual: SN 2020zbf
1. Introduction
Modern time-domain sky surveys with large fields of view
are able to detect and follow rare transient events. Superlu-
minous supernovae (SLSNe) are an extremely luminous class
of transients, 10–100 times brighter than canonical super-
nova (SN) explosions (Quimby et al. 2011; Gal-Yam 2012).
The need for a new class of SN arose due to the fact that
some events (Nugent et al. 1999; Ofek et al. 2007; Smith et al.
2007; Quimby et al. 2007; Miller et al. 2009; Barbary et al.
2009; Gal-Yam et al. 2009; Gezari et al. 2009) are much brighter
than the majority of previously discovered events and could
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A20, page 1 of 26
Gkini, A., et al.: A&A, 685, A20 (2024)
not fit into the conventional SN explosion scenario. SLSNe
are frequently detected in metal-poor dwarf host galaxies
(Neill et al. 2011; Chen et al. 2013, 2017; Lunnan et al. 2014;
Leloudas et al. 2015; Angus et al. 2016; Perley et al. 2016;
Schulze et al. 2018) and were originally defined as having an
absolute magnitude of M < −21 (Gal-Yam 2012). How-
ever, many SLSNe have peak magnitudes less than this thresh-
old (Inserra et al. 2013, 2018a; Angus et al. 2019; Lunnan et al.
2018; De Cia et al. 2018; Chen et al. 2023a), and SLSN clas-
sification is better determined using spectroscopic properties
(Quimby et al. 2011, 2018; Gal-Yam 2019b) rather than an arbi-
trary brightness cut.
The SLSN class is typically divided into two subgroups
based on the presence of hydrogen in their spectra around the
peak – type I (H-poor; SLSN-I hereafter) and type II (H-rich;
SLSN-II) – although a few SLSNe-I have Hα emission detected
in their late-time spectra (Yan et al. 2015, 2017a; Chen et al.
2018; Fiore et al. 2021; Pursiainen et al. 2022). In particular,
SLSNe-I whose spectra do not show He are characterized as
SLSNe-Ic. It has been proposed that SLSNe-I can be further
classified photometrically: “slow-evolving” SLSNe-I show a
rise of about 50 days and a decline consistent with the 56Co
decay rate, whereas “fast-evolving” SLSNe-I have rise times
of less than 30 days (Nicholl et al. 2016; Quimby et al. 2018;
Inserra 2019). However, with some events being in the interme-
diate regime (e.g. Gaia16apd, Kangas et al. 2017; SN 2017gci,
Fiore et al. 2021), there are claims of a continuous distribution
rather than distinct subclasses (Nicholl et al. 2016; De Cia et al.
2018).
The dominating powering mechanism for SLSNe-II is likely
interaction with a dense circumstellar medium (CSM; Ofek et al.
2014; Inserra et al. 2018b, but see Kangas et al. 2022). How-
ever, for H-poor SLSNe (Chevalier & Irwin 2011; Ofek et al.
2013; Inserra et al. 2018a), the powering mechanism is still
poorly understood. The amount of radioactive 56Ni produced in
the standard core collapse mechanism is insufficient to explain
the extreme brightness of SLSNe-I; thus, other mechanisms
have been proposed. One proposed mechanism is pair-instability
supernovae (PISNe), in which the formation of positron–electron
pairs in the CO core of a 140–260 M zero-age main sequence
(ZAMS) metal-poor star results in explosive O burning, and the
energy released reverses the collapse and entirely disrupts the
star (Woosley et al. 2002; Heger & Woosley 2002). This pro-
cess has the potential to generate the enormous amount of 56Ni
required to power a SLSN-I light curve. Despite the presence
of a few PISN candidates (Gal-Yam et al. 2009; Schulze et al.
2023), numerous earlier studies have demonstrated that 56Ni is
not the dominant source of energy for the majority of SLSNe-I
(Chatzopoulos et al. 2012; Chen et al. 2013, 2023b; Inserra et al.
2013, 2018b; Nicholl et al. 2017; Moriya et al. 2018). The
majority of the observed photometric features of SLSNe-
I (Inserra et al. 2013; Nicholl et al. 2017; Chen et al. 2023b;
Omand & Sarin 2024) can instead be attributed to the spin-down
of a rapidly rotating young neutron star (Ostriker & Gunn 1971;
Kasen & Bildsten 2010; Woosley 2010), in which the photons
from the newborn magnetar wind nebula are absorbed and ther-
malized in the SN ejecta, increasing the temperature of the
ejecta and the luminosity of the SN. The spectroscopic signa-
tures of magnetar-powered SLSNe-I have yet to be explored in
detail, but simulations demonstrate that magnetars can reproduce
the observed SLSN-I spectra (Dessart et al. 2012; Mazzali et al.
2016; Jerkstrand et al. 2017; Dessart 2019; Omand & Jerkstrand
2023).
The majority of the early SLSN-I spectra exhibit a steep
blue continuum (Yan et al. 2017b, 2018) and present distinct
spectroscopic key characteristics (Quimby et al. 2011, 2018;
Mazzali et al. 2016). The presence of O ii features dominates
the spectra at 3500–5000 Å, with the most significant W-shape
feature at 4350–4650 Å, which is not typically seen in nor-
mal SNe-Ic. However, numerous SLSNe-I in the literature do
not appear to have the W-shape O ii in their spectra (e.g.,
Gutiérrez et al. 2022), suggesting a further division of the SLSN-
I class (Könyves-Tóth & Vinkó 2021). The red part of the
optical SLSN-I spectra presents weak C and O lines, which
Dessart et al. (2012) and Howell et al. (2013) suggest come from
the explosion of the CO core. The spectra at ∼30 days resemble
those of SNe-Ic around maximum light (Pastorello et al. 2010).
Interestingly, several SLSNe-I in the literature do not fit into
this “standard” classification scheme. These events show strong
C ii lines in their spectra (Yan et al. 2017a; Quimby et al. 2018;
Anderson et al. 2018; Fiore et al. 2021; Gutiérrez et al. 2022).
Anderson et al. (2018) suggest that the persistent C ii features
in SN 2018bsz are produced by a magnetar-powered explosion
of a massive C-rich Wolf-Rayet (WR) progenitor. The mod-
els of Fiore et al. (2021) suggest that the C-rich SN 2017gci
was powered by either a magnetar or CSM interaction with
a 40 M progenitor. Additionally, Gutiérrez et al. (2022) find
that SN 2020wnt requires over 4 M of 56Ni to produce the
observed light curve, which is consistent with the PISN scenario,
while Tinyanont et al. (2023) favor a magnetar model. Various
ideas have been proposed to explain these characteristics, but
the reasons for the presence of the C ii lines are still poorly
understood.
In this paper we present an extensive dataset for SN 2020zbf,
a fast-rising SLSN-I with peculiar features in its early spectrum
that initially led to an incorrect redshift estimation. A medium-
resolution X-shooter spectrum displays three strong C ii lines,
indicating that SN 2020zbf belongs to the C-rich SLSN-I sub-
class. We performed an extensive investigation of the observed
data, modeling the light curve and comparing the high-quality
spectra with synthetic models. This allowed us to explore var-
ious combinations of power sources and progenitor stars that
could result in these spectral signatures. This object, along with
other rare SLSNe, may challenge the conventional classification
scheme by demonstrating how diverse even the SLSN-I class
may be, with implications for both progenitor populations and
explosion mechanisms.
This paper is structured as follows. In Sect. 2 we present
the photometric and spectroscopic observations of SN 2020zbf
as well as the photometric measurements of the host galaxy.
In Sect. 3 we analyze the light curve properties, compare them
with those of well-studied SLSNe-I, and apply blackbody fits
to derive the photospheric temperatures and radius. In Sect. 4
we analyze the spectral properties of SN 2020zbf. We compare
the light curves and the early and the late photospheric spec-
tra with those of typical SLSNe-I in the literature as well as C-
rich objects in Sect. 5. In Sect. 6 we compare existing synthetic
spectra with our high-quality X-shooter spectrum. We model
the multiband light curves of SN 2020zbf under the assump-
tion that they are powered by three distinct power sources
in Sect. 7. In Sect. 8 we discuss the properties of the host
galaxy. We discuss the results and provide possible scenarios
in Sect. 9 and summarize our findings in Sect. 10. Throughout
this paper we assume a flat Lambda cold dark matter cosmol-
ogy with H0 = 67.4 km s−1 Mpc−1, ΩM = 0.31, and ΩΛ = 0.69
(Planck Collaboration VI 2020).
A20, page 2 of 26
Gkini, A., et al.: A&A, 685, A20 (2024)
Legacy Survey g/r/i
(pre-explosion)
20"
N
E
LCO g/r/i
11 November 2020
Fig. 1. Images of the field of SN 2020zbf. Left: Legacy Survey DR10 image of the field of SN 2020zbf before explosion. A faint host galaxy at the
SN position is visible, marked by the white crosshairs. Right: gri composite image of the SN near peak from Las Cumbres Observatory (LCO).
Both images have a size of 2 × 2 arcmin and have been combined following the algorithm in Lupton et al. (2004).
2. Observations
2.1. Discovery and classification
SN 2020zbf was discovered by the Asteroid Terrestrial-impact
Last Alert System (ATLAS; Tonry et al. 2020) on Novem-
ber 8, 2020, as ATLAS20bfee at an orange-band magnitude of
18.92 mag at right ascension, declination (J2000) 01h58m01.65s,
−41◦20′51.89′′. It was classified by the extended Public
ESO Spectroscopic Survey for Transient Objects (ePESSTO+;
Smartt et al. 2015) as a SLSN-I (see Sect. 4.1) on November 9,
2020 (Ihanec et al. 2020). An image of the field showing the host
galaxy from the Legacy Survey (Dey et al. 2019), as well as an
image showing the SN near peak, are shown in Fig. 1.
We adopted a redshift of z = 0.1947 (see Sect. 4.1) and
computed the distance modulus to be 39.96 mag. In order to
compute the Milky Way (MW) extinction, we used the dust
extinction model of Fitzpatrick (1999) based on RV = 3.1 and
E(B − V) = 0.014 mag (Schlafly & Finkbeiner 2011). As for
the host galaxy extinction, we find that the host of SN 2020zbf
is a faint, blue dwarf galaxy, quite typical of SLSN-I host galax-
ies (Sect. 8; Lunnan et al. 2014; Perley et al. 2016; Schulze et al.
2018). The host galaxy analysis supports moderate host extinc-
tion [E(B−V)host = 0.22+0.20−0.22 mag]. However, given that the host
properties are consistent with no extinction within the uncer-
tainties, we did not apply any host galaxy extinction correction
to the light curves. The estimated epoch of maximum light is
November 11, 2020, MJD = 59 164.8 (see Sect. 3.2).
2.2. Photometry
2.2.1. ATLAS
ATLAS is a wide-field survey consisting of four telescopes
that scan the whole sky with daily cadence (Tonry et al. 2018).
ATLAS observes in two wide filters, the cyan (c) and orange (o)
down to a limiting magnitude of ∼19.7 mag (Tonry 2011) and the
data are processed using the pipeline described in Stalder et al.
(2017).
We retrieved forced photometry from the ATLAS forced
photometry server1 (Tonry et al. 2018; Smith et al. 2020;
1 https://fallingstar-data.com/forcedphot/
Shingles et al. 2021) for both c and o filters. We computed the
weighted average of the fluxes of the observations on nightly
cadence. We performed a quality cut of 3σ in the resulting flux
of each night for each filter and converted them to the AB magni-
tude system. The resulting data span from −22 to +68 rest-frame
days post maximum light and the observed photometry is listed
in Table A.1.
2.2.2. Las Cumbres Observatory
SN 2020zbf was monitored by ePESSTO+ between November
2020 and December 2021 using the Las Cumbres Observatory
(LCO) in the BgVriz filters. The data were collected using the
1-m telescopes on South African Astronomical Observatory,
Cerro Tololo Inter-american Observatory and Siding Spring
Observatory. Reference images to perform image subtraction
were taken in September 2022.
We performed photometry using the AUTOmated Pho-
tometry of Transients (AUTOPhoT2) pipeline developed by
Brennan & Fraser (2022). AUTOPhoT removes host galaxy con-
tamination through image subtraction using the HOTPANTS
(Becker 2015) software. The instrumental magnitude of the SN
is measured through point-spread function fitting and the zero
point in each image is calibrated with stars from the Legacy
Survey (Dey et al. 2019) and SkyMapper Southern (Onken et al.
2019) catalogs. The LCO light curve covers the range from 0 to
180 rest-frame days post maximum light. For nights with mul-
tiple exposures, we computed the weighted average. We do not
discuss the z-band photometry because of the poor quality of
these images. The final photometry is listed in Table A.1.
2.2.3. Neil Gehrels Swift Observatory
SN 2020zbf was observed with the UV/Optical Telescope
(UVOT; Roming et al. 2005) on the Neil Gehrels Swift Obser-
vatory (Gehrels et al. 2004) in all six filters, ranging from ultra-
violet (UV) to visible wavelengths. The UVOT photometry is
2 https://github.com/Astro-Sean/autophot
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Gkini, A., et al.: A&A, 685, A20 (2024)
Table 1. Host galaxy photometry.
Instrument/Filter Brightness
(mag)
LS/g 22.58 ± 0.03
LS/r 22.02 ± 0.03
LS/i 21.85 ± 0.05
LS/z 21.69 ± 0.06
DES/y 21.56 ± 0.24
Notes. All magnitudes are reported in the AB system and are not cor-
rected for MW extinction.
retrieved from the NASA Swift Data Archive3 and processed
using UVOT data analysis software HEASoft version 6.194. The
reduction of the images is achieved by extracting the source
counts from the images within a radius of 3 arcseconds and the
background was estimated using a radius of 48 arcseconds. We
used the Swift tool UVOTSOURCE to extract the photometry using
the zero points from Breeveld et al. (2011) and the calibration
files from September 2020.
The four UVOT epochs cover the range +12 to +24 rest-
frame days past maximum, in all six UVOT filters. Since we
have LCO B- and V-band data with better coverage, we omitted
UVOT B- and V-band data from further analysis. The photome-
try is listed in Table A.1.
2.3. Host galaxy photometry
We retrieved the science-ready images from the DESI Legacy
Imaging Surveys (Dey et al. 2019) Data Release (DR) 10 and
complemented the dataset with archival y-band observations
from the Dark Energy Survey DR 1 (Abbott et al. 2018). The
photometry was extracted with the aperture-photometry tool pre-
sented by Schulze et al. (2018)5. Table 1 summarizes all mea-
surements.
2.4. Spectroscopy
We obtained six low-resolution spectra of SN 2020zbf between
November 9, 2020, and January 19, 2021, with the ESO Faint
Object Spectrograph and Camera 2 (EFOSC2; Buzzoni et al.
1984) on the 3.58 m ESO New Technology Telescope (NTT)
at the La Silla Observatory in Chile under the ePESSTO+ pro-
gram (Smartt et al. 2015). We complemented this dataset with
one medium-resolution spectrum on November 18, 2020, with
the X-shooter spectrograph (Vernet et al. 2011) on the ESO Very
Large Telescope (VLT) on Paranal, Chile. The spectral log is pre-
sented in Table B.1.
The NTT spectra were reduced with the PESSTO6 pipeline.
The observations were performed with grisms #11, #13, and
#16 using a 1′′.0 wide slit. The integration times varied between
900 and 5400 s. The spectrum taken on December 8, 2020, is
excluded from the analysis due to the poor signal-to-noise ratio
(S/N).
The X-shooter observations were performed in nodding
mode using 1′′.0, 0′′.9, 0′′.9 wide slits for the UV, visible (VIS),
3 https://heasarc.gsfc.nasa.gov/cgi-bin/W3Browse/
swift.pl
4 https://heasarc.gsfc.nasa.gov/
5 https://github.com/steveschulze/Photometry
6 https://github.com/svalenti/pessto
and near-infrared (NIR) arms, respectively and were reduced
using the ESO X-shooter pipeline. The procedure is the fol-
lowing; first the removal of cosmic-rays is done using the tool
astroscrappy7, based on the algorithm of van Dokkum (2001),
then the data were processed with the X-shooter pipeline v3.6.3
and the ESO workflow engine ESOReflex (Goldoni et al. 2006;
Modigliani et al. 2010) and finally telluric absorption features
in the VIS arm were removed with the Molecfit version 4.3.1
(Smette et al. 2015; Kausch et al. 2015). The wavelength cali-
bration of all spectra was adjusted to account for barycentric
motion. The spectra of the individual arms were stitched together
by averaging the overlap regions.
Each spectrum was flux calibrated against standard stars.
The spectral evolution from −2.4 to +57 rest-frame days past
maximum brightness are depicted in Fig. 2. All the spectra are
uploaded on the WISeREP8 archive (Yaron & Gal-Yam 2012).
3. Light curve analysis
We estimated the absolute magnitudes in each filter using the
following expression:
M = m − µ − AMW − Kcorr, (1)
where m is the apparent magnitude, µ is the distance modulus,
AMW is the extinction caused by the MW and the last term is
associated with the K-correction. The K-correction relates the
photometric bandpasses in the rest frame and observer frame.
It can be separated into two terms; the first term corrects for
the redshift and the second term also for the shape of the spec-
trum (Hogg et al. 2002). In this case, we considered only the
first term, −2.5 log(1 + z), which is a good approximation for the
total K-correction as shown in Chen et al. (2023a). We estimate
Kcorr = −0.19 mag for all bands and epochs. The multiband light
curve in apparent and absolute magnitude systems are shown in
Fig. 3.
3.1. Time of first light
The rising part of the light curve was only observed with
ATLAS, since LCO follow-up was triggered only after the
SN was classified near peak light. Figure 3 shows the most
recent upper limits in the ATLAS c and o filters before the
first detections (from forced photometry) at MJD 59 137.5 and
MJD 59 147.3, respectively. Initially, we fit both the c and o fil-
ters separately to calculate the time of first light. However, we
find that the estimated best-fit time of first light in the o band is
later than the first c-band detection. This can be understood from
Fig. 3, since the last non-detection in the o band (20.54 mag) is
also after the first c-band detection. We therefore used the bluer
c band for the calculation of first light, despite the o band being
better sampled. The contemporaneous detection in the c band
and the non-detection in the o band sets a limit on the color at
this time to c − o < 0.3 mag.
Following Miller et al. (2020), we fit a Heaviside step func-
tion multiplied by a power law to simultaneously fit the pre-
explosion baseline and the rising light curve in the ATLAS c
filter. Using the PYTHON module emcee (Foreman-Mackey et al.
2013) the power-law index αc is estimated to be 0.68+0.15−0.14 and
the time of first light to be MJD 59 135.4+1.3−2.1. We note that these
error bars only account for the statistical errors in the fit and
7 https://github.com/astropy/astroscrappy
8 https://www.wiserep.org
A20, page 4 of 26
Gkini, A., et al.: A&A, 685, A20 (2024)
3000 4000 5000 6000 7000 8000
Rest-frame wavelength (Å)
2
1
0
1
2
Sc
al
ed
 F
 +
 o
ffs
et
 (e
rg
 s
1  c
m
2  Å
1 )
 -2.4 d
+2.7 d
+4.3 d
+33.6 d
+43.7 d
+57 d
Fig. 2. Spectral sequence of SN 2020zbf from −2.4 to +57 rest-frame days. We highlight the X-shooter spectrum in purple. An offset in flux was
applied for illustration purposes. The spectra are corrected for the MW extinction and are smoothed using a Savitzky-Golay filter. The original
data are presented in lighter colors. See Sect. 4.2 for details on line identification.
not for any systematic errors associated with the method chosen.
The approach in Miller et al. (2020) is based on the modeling of
a different type of SN and the uncertainty in the explosion date
in SN 2020zbf may be larger due to qualitative differences in the
rise of SLSNe-I (Nicholl et al. 2015).
3.2. Peak magnitude, timescales, and color evolution
To estimate the epoch of the maximum light as well as vari-
ous light-curve timescales, we interpolated the c- and the g-band
light curves. We applied the method from Angus et al. (2019)
for the light-curve interpolation and fit a Gaussian process (GP)
regression. To do this, we utilized the PYTHON package GEORGE
(Ambikasaran et al. 2015) with a Matern 3/2 kernel.
The c- and g-band photometric data with the resulting inter-
polations are shown in Fig. 4. The g-band light curve is already
declining by the first observation, and we took the first data point
as a lower limit on the g-band peak absolute magnitude: Mg is
−21.18±0.07 mag, observed at MJD 59164.8. This in turn gives
an upper limit in the rise time of .26.4 rest-frame days, includ-
ing the 2.1 days statistical error on the estimated time of first
light.
A20, page 5 of 26
Gkini, A., et al.: A&A, 685, A20 (2024)
26
24
22
20
18
16
14
Ab
so
lu
te
 M
ag
ni
tu
de
 (m
ag
)
100 50 0 50 100 150 200
Rest-frame days past peak
12
14
16
18
20
22
24
26
Ap
pa
re
nt
 M
ag
ni
tu
de
 (m
ag
)
UVOT-UVW2  -8
UVOT-UWM2  -6
UVOT-UVW1  -4.5
UVOT-U  -2
LCO-B  -1
LCO-g
ATLAS-c  +1
LCO-V  +2
LCO-r  +3
ATLAS-o  +4
LCO-i  +5
NTT/EFOSC2
VLT/X-shooter
Fig. 3. UV and optical light curves of SN 2020zbf. The magnitudes are corrected for MW extinction and cosmological K-correction. Upper limits
are presented as downward-pointing triangles in a lighter shade. The phase of the first detection is marked with a vertical dashed line and the
epochs of the spectra are marked as thick lines at the top of the figure. The zero value on the x-axis is with respect to the g-band maximum (MJD
59 164.8).
The rise and decline timescales of the light curve described in
Chen et al. (2023a) are determined using the c-band interpolated
light curve and the maximum Mg. The rest-frame rise time from
the half maximum flux (Fg,peak/2) is 12.2+1.2−2.4 days and from 1/e
maximum flux (Fg,peak/e) is 15.9+1−1.1 days. We estimated the rest-
frame decline times using the g-band interpolated light curve.
The decline time to the half maximum flux is 26.86+1.44−1.44 days
and to the 1/e maximum flux is 41.9+2−1.9 days. These are shown
in gray lines in Fig. 4.
In Fig. 5, we put the light curve properties of SN 2020zbf
in the context of the homogeneous Zwicky Transient Facil-
ity (ZTF; Bellm et al. 2019) SLSN-I sample from Chen et al.
(2023a). This paper studied the UV and optical photometric
properties of 78 H-poor SLSNe-I. In the three different panels,
we show the kernel density estimates (KDEs) of the ZTF sample,
which are an outcome of a Monte Carlo simulation accounting
for the asymmetric errors, and indicate by the red vertical lines
the measurements for SN 2020zbf. The peak absolute magnitude
is fairly average, being slightly fainter than the median value of
the SLSNe-I. In contrast, the rise time of SN 2020zbf is among
the fastest seen for SLSNe-I, whereas the decline is again rather
average.
To construct the g−r color evolution of SN 2020zbf, we used
the g- and r-band interpolated light curves and plot the results
in Fig. 6. For comparison, we present the reddening corrected
g − r colors of the SLSNe-I from the ZTF sample of Chen et al.
(2023a) with redshifts within ±0.02 of SN 2020zbf’s redshift (in
order to facilitate comparison at similar effective wavelengths).
Although the g − r color evolution of SN 2020zbf follows the
general trend of the ZTF sample, by getting redder over time, it
evolves more slowly than other SLSNe-I showing a consistently
bluer color.
3.3. Photospheric temperature and radius
We interpolated all the UVOT and LCO light curves using the
GP method described in Sect. 3.2 and extracted the magni-
tudes using the V-band epochs as reference since it has the
most observed epochs. We excluded the ATLAS filters because
they are significantly broader than the LCO filters. We con-
structed the spectral energy distributions (SEDs) by calculat-
ing the spectral luminosities Lλ for each band, at each of
the 14 past-peak epochs, and fit a blackbody utilizing the
scipy.optimize.curvefit9 module. Due to line blanket-
ing (Yan et al. 2017b), we excluded the UVOT data from the
9 https://docs.scipy.org/doc/scipy/reference/generated/
scipy.optimize.curve_fit.html
A20, page 6 of 26
Gkini, A., et al.: A&A, 685, A20 (2024)
40 20 0 20 40 60 80 100
Rest-frame days past peak
21.5
21.0
20.5
20.0
19.5
19.0
18.5
Ab
so
lu
te
 m
ag
ni
tu
de
Mg peak
Mg at Fg, peak/2
Mg at Fg, peak/e
LCO-g
ATLAS-c
Fig. 4. Gaussian process interpolation of ATLAS c- and LCO g-band
light curves. The peak magnitude in the g band, Mg, is indicated with the
vertical line and the various rise and decline times with the horizontal
lines.
23.022.522.021.521.020.520.019.5
Mg peak magnitude
0.2
0.4
SN 2020zbf
Median
0 20 40 60 80 100 120
trise, 1/e (days)
0.00
0.01
0.02
0.03
De
ns
ity
0 20 40 60 80 100 120
tdecline, 1/e (days)
0.00
0.01
0.02
Fig. 5. Comparison of the photometric properties of SN 2020zbf with
the ZTF SLSN-I sample (Chen et al. 2023a). Top: KDE distribution of
the Mg peak magnitudes for 78 ZTF SLSNe-I. Middle: KDE plot of the
e-folding rise time for 69 ZTF SLSNe-I. Bottom: e-folding decline time
distribution for 54 ZTF SLSNe-I. The vertical red line along with the
errors (shaded red regions) illustrate the position of SN 2020zbf and the
black vertical lines the median values.
blackbody fits. The resulting photospheric BgVri temperature
and radius evolution are plotted in Fig. 7.
To check whether there is consistency with the spectral mea-
surements, we estimated the temperature and the radius by fit-
ting a blackbody to the spectra taken in the early photospheric
phase. We first absolute-calibrated the spectra with the photo-
metric data before template subtraction and then corrected them
for the MW extinction. The results are shown in Fig. 7 along
with the SLSNe-I from the Chen et al. (2023a) sample in light
gray. We chose to compare with the ZTF SLSNe-I, which are
characterized as normal events by Chen et al. (2023a), excluding
40 20 0 20 40 60 80 100
Rest-frame days past peak
0.5
0.0
0.5
1.0
1.5
g
r c
ol
or
 (m
ag
)
SN 2020zbf
Chen et al.(2023a) with 0.18<z<0.22
Fig. 6. g − r color evolution of SN 2020zbf along with the GP interpo-
lation in light purple. For comparison, the ZTF g − r color evolution of
the Chen et al. (2023a) SLSN sample with 0.18 < z < 0.22 is presented
in gray.
4
6
8
10
12
14
16
18
20
Te
m
pe
ra
tu
re
 (1
03
 K
)
SN 2020zbf
Chen et al. (2023a)
40 20 0 20 40 60 80 100
Rest-frame days past peak
0
2
4
Ra
di
us
 (1
01
5  c
m
)
Fig. 7. Blackbody temperatures and radii of SN 2020zbf. Top: tempera-
ture evolution of SN 2020zbf derived from the blackbody fits of the pho-
tometric (green) and the spectroscopic (blue) data. The gray background
points present the temperature evolution of the ZTF sample (Chen et al.
2023a). Bottom: blackbody radius evolution of SN 2020zbf using pho-
tometry (green) and the early spectra (blue).
the objects with fewer than two epochs as well as three objects
marked in Chen et al. (2023a) as extraordinary events. Overall,
the temperature evolution of SN 2020zbf is comparable to those
of the ZTF sample but the temperatures are higher than for the
bulk of the population, which is in agreement with the color
evolution in Fig. 6. The radius of the photosphere shows a rise
trend up to ∼45 days, which is consistent with what is seen in the
ZTF sample (Chen et al. 2023a). After 45 days, the photospheric
radius seems to decline.
3.4. Bolometric light curve
To construct a bolometric light curve, we started by integrating
the observed SED (BgVri) at the epochs we have LCO data.
The result is shown as green points in Fig. 8, and constitutes a
strict lower limit on the bolometric luminosity from the observed
A20, page 7 of 26
Gkini, A., et al.: A&A, 685, A20 (2024)
flux over the optical bands only. For a better estimate on the
total bolometric luminosity, we considered different methods to
account for the missing flux. For the NIR correction, we fit a
blackbody to the LCO data as before and integrated the black-
body tail up to 24 400 Å, beyond which the contribution to the
bolometric light curve is negligible in the photospheric phase
(∼1%; Ergon et al. 2013).
For the UV correction, which constitutes a significant frac-
tion of the bolometric flux at early times when the temperature
is high, we used two different methods for comparison. First, we
considered a blackbody method, integrating the same blackbody
fit from 0 Å to the B band, and adding up the UV, observed opti-
cal, and NIR flux. This is shown as the purple curve in Fig. 8, and
can be considered an upper limit on the bolometric luminosity,
since it does not take into account UV line blanketing. To better
capture this effect, we finally linearly extrapolated the SED from
the B band to 2000 Å (where Lλ is assumed to be zero), follow-
ing Lyman et al. (2014); the pseudo-bolometric light curve using
this UV correction is shown as black points in Fig. 8. Near peak,
the linear extrapolation method adds ∼0.31 dex to the observed
flux while the blackbody method adds ∼0.62 dex. In later epochs,
when the UV emission is small, the two luminosities have con-
sistent values.
We tested the validity of our UV corrections against the
four epochs for which we have UVOT data. We integrated the
observed SED using the BgVri and UVOT filters and added the
NIR flux we calculated above. The four-epoch bolometric light
curve is shown in orange in Fig. 8. The data points are placed
between the UV linear extrapolation method (Lyman et al. 2014)
and the blackbody fit corrected light curves, indicating that the
flux is overestimated by ∼0.1 dex when integrating the full black-
body in the UV and underestimated by ∼0.1 dex when using the
UV linear correction. Since we know that SLSNe-I do have sig-
nificant UV absorption (Yan et al. 2017b), we used the linear
correction for the rest of the analysis but note that this also some-
what underestimates the total flux. The peak bolometric luminos-
ity Lpeakbol is estimated to be &8.3× 1043 erg s−1.
At our earliest and latest epochs, we did not have the SED
information to perform the same analysis as above. In order to
include these points in the bolometric light curve (for purposes
of integrating the total energy), we instead assumed a constant
bolometric correction. For the data on the rise, we used the c
band and used the same ratio of the c-band flux to total flux that
we measured at the first epoch with multiband data. Since we
expect the temperature to be higher at these earlier epochs, this
approximation will be a lower limit. Similarly, our final epoch
only has r and i measurements, and we scaled the bolometric
luminosity to the latest multiband measurement. These points
are shown as square symbols in Fig. 8. Integrating the pseudo-
bolometric light curve, including these early and late measure-
ments, we find the total radiated energy of SN 2020zbf to be
Erad & 4.68 ± 0.14 × 1050 erg, consistent with what we would
expect for SLSNe-I (Quimby et al. 2011).
4. Spectral analysis
4.1. Host galaxy redshift and spectral classification
The first spectrum of SN 2020zbf was obtained on
November 9, 2020, and and was classified by Ihanec et al.
(2020) as a SLSN-Ic at z = 0.35 using the cross matching tool
SNID (Blondin & Tonry 2007); it was found to be a good match
to the SLSN-I PTF12dam (Nicholl et al. 2013). Based on this
classification we triggered our X-shooter program; the higher-
50 0 50 100 150 200
Rest-frame days past peak
42.50
42.75
43.00
43.25
43.50
43.75
44.00
44.25
44.50
lo
g 1
0 (
L b
ol
) (
er
g 
s
1 )
Observed L (BgVri)
Observed L + corr (UV linear extrapolation)
Observed L + corr (BB fit)
Observed L + corr (UVOT epochs)
Fig. 8. Bolometric light curve of SN 2020zbf with different corrections
applied. The circles correspond to the derived luminosities using all
LCO filters, the crosses are the Vri bands and the square symbols illus-
trate the bolometric luminosity assuming the same bolometric correc-
tion as the epochs with multiband data. The green light curve corre-
sponds to the integrated observed BgVri flux and constitutes a lower
limit on the total bolometric luminosity. The black symbols include a
correction to the NIR and the UV using the UV linear extrapolation
method (Lyman et al. 2014), whereas the purple light curve considers a
UV correction by integrating the blackbody fit. The orange data points
show the observed bolometric luminosity including the NIR correc-
tion if we consider also the UVOT data. The errors represent statistical
errors. For our analysis we used the black curve, and the dashed lines
connect the correction data points for illustration purposes.
quality spectrum taken on November 18, 2020, reveals narrow
emission and absorption lines consistent with a galaxy at a
redshift of z = 0.195 (Lunnan & Schulze 2021). We reexamined
the spectrum and identified emission lines from the interstellar
medium and H ii regions in the host galaxy at a common redshift
of z = 0.1947 ± 0.0001. Figure 9 shows the host lines that
were used for the redshift determination: the galaxy’s narrow
Mg ii doublet λ2796, 2803, the Mg i λ2852, the [O ii] doublet
λλ3727, 3729, the narrow Hα λ6563, the Hβ λ4861, and the
forbidden [O iii] doublet λλ4959, 5007. The initial redshift
misclassification highlights both the peculiarity of SN 2020zbf
in comparison with typical SLSNe-I, as well as the challenge of
classifying unusual events based on lower-quality data.
Having established a robust redshift, we tentatively identified
the two absorption features around 4500 Å as O ii, supporting a
SLSN-I spectroscopic classification (e.g., Quimby et al. 2018),
although the velocities would be quite low (see Sect. 4.2). Even
without the O ii line identification, we argue that SN 2020zbf
is best classified as a SLSN-Ic. The absolute magnitude of
SN 2020zbf in the g band is −21.18 mag (see Sect. 3.2),
which is well within the SLSN brightness regime (Gal-Yam
2012; De Cia et al. 2018; Inserra et al. 2018b; Angus et al. 2019;
Chen et al. 2023a). There is no evidence of H lines in the SN
spectrum (the feature at 6578 Å is more likely to be C ii; see
Sect. 4.2) and also no obvious He lines. Finally, as the ejecta
cool, the spectrum evolves to look like a typical SN-Ic.
4.2. Line identification
Line identification of SN 2020zbf was performed using the
medium-resolution X-shooter spectrum at +4.3 days and the
A20, page 8 of 26
Gkini, A., et al.: A&A, 685, A20 (2024)
3317.7 3383.4 3449.1
2777 2827 2877
0
1
2
3
MgII MgI
4420.4 4455.0 4489.6
Observed wavelength (Å)
3700 3729 3758
1.4
1.5
1.6
1.7
1.8 [O II]
5774.0 5809.7 5845.4
4833 4863 48930.6
0.7
0.8
0.9
1.0
H
5913.8 5934.6 5955.4 5976.2 5997.0
4950 4967 4984 5001 5018
0.6
0.7
0.8
0.9 [O III][O III]
7806.2 7842.0 7877.8
6534 6564 65940.3
0.4
0.5
0.6
0.7
0.8
0.9 H
Rest-frame wavelength (Å)
F
 (1
0
16
 e
rg
 s
1  c
m
2  Å
1 )
Fig. 9. Host galaxy absorption and emission lines in the X-shooter spec-
trum of SN 2020zbf +4.3 days after peak. The absorption and emission
lines give a consistent host redshift of z = 0.1947.
low-resolution NTT-EFOSC2 spectrum at +43.7 days after
maximum brightness; these spectra have the highest S/N.
The identification was done by comparing with other SLSNe
from the literature (Inserra et al. 2013; Anderson et al. 2018;
Quimby et al. 2018; Gal-Yam 2019b; Pursiainen et al. 2022;
Tinyanont et al. 2023), with the predictions of spectral mod-
els (Dessart et al. 2012; Mazzali et al. 2016; Dessart 2019), and
finally by searching the National Institute of Standards and Tech-
nology (NIST; Kramida et al. 2023) atomic spectra database for
lines above a certain strength, similar to what was done in
Gal-Yam (2019b). Figures 10 and 11 show the +4.3 and +43.7
days spectra, respectively, along with the most conspicuous fea-
tures blueshifted by 9000–12 000 km s−1 (see Sect. 4.4). The C ii
lines are shown at zero rest-frame velocity.
Early spectra of SLSNe-I are characterized mostly by
the presence of strong O ii lines between 3500 and 5000 Å
(Quimby et al. 2011, 2018; Mazzali et al. 2016). In the case
of SN 2020zbf, only the O ii λ4358 and λ4651 W-shape is
visible with absorption troughs consistent with a velocity of
4000 km s−1 (see Sect. 4.4), which is relatively low for SLSNe-I
(median value of 9700 km s−1; Chen et al. 2023a). The absorp-
tion at 4300 Å seems stronger and broader compared to the
one at 4600 Å, indicating the presence of other ions at these
wavelengths (Nicholl et al. 2016). In particular, we associate
the feature at 4300 Å with a blend of O ii λ4358, Fe iii λ4432
and Mg ii λ4481. Redward of the O ii λ4651, the Fe ii triplet
λλ4923, 5018, 5169 is present, but the Fe ii λ5169 is likely mixed
with Fe iii λ5129 (Liu et al. 2017).
Blueward of 3000 Å, the UV part of the spectrum is very
blended and it is hard to identify individual lines. There are two
strong absorption components at 2670 Å and 2880 Å. The first
trough at 2670 Å is associated with Mg ii, C ii, and C iii, which
is consistent with what is seen in other SLSNe-I (Quimby et al.
2011; Lunnan et al. 2013; Howell et al. 2013; Vreeswijk et al.
2014; Yan et al. 2017b, 2018; Smith et al. 2018) and in spec-
tral models (Dessart et al. 2012; Mazzali et al. 2016; Dessart
2019). The second absorption at 2880 Å has been observed in a
number of SLSNe (iPTF 13ajg; Vreeswijk et al. 2014, PTF09atu
and PTF12dam; Quimby et al. 2018) but not as strong as in
SN 2020zbf and has never been conclusively identified. Accord-
ing to Mazzali et al. (2016) and Quimby et al. (2018) these fea-
tures are mostly attributed to Ti iii. Dessart et al. (2012) also sug-
gested Fe iii, Si iii, and S iii. Searching the NIST library, we find
that other possible contributions in this region could be C ii and
Mg ii.
Between 3000 and 3600 Å, the absorption at 3200 Å could be
attributed to Fe ii λ3325 and Fe iii λ3305. The feature at 3550 Å
is associated with Fe iii λ3691, but we were unable to identify
the ions that may contribute to the feature at 3410 Å. We also
see an absorption feature at 7550 Å that could be formed by the
O i triplet λλ7772, 7774, 7775 and a small contribution of Mg ii
λ8234.
Figure 11 shows the spectrum at +43.7 days. The spec-
trum at this stage resembles that of a normal SN-Ic at max-
imum light (Pastorello et al. 2010; Quimby et al. 2011). The
O ii and Fe iii lines, which dominated the early spectra, have
disappeared and elements from further in are revealed as the
ejecta cool down. We see Ca ii λλ3966, 3934, Mg i] λ4571
and strong Fe ii lines between 4000 and 5200 Å blueshifted by
9000 km s−1 (see Sect. 4.4) to match the absorption component.
We cannot distinguish the O i triplet λλ7772, 7774, 7775 due to
the low S/N. At 3600 Å, there is a broad emission component
that has been observed in the SLSNe-Ic LSQ12dlf (Nicholl et al.
2014, 2015), SN 2007bi (Young et al. 2010; Nicholl et al. 2013),
and SN 2017egm (Bose et al. 2018). In Nicholl et al. (2014) and
Bose et al. (2018) the result from spectrum synthesis code for
LSQ12dlf and SN 2017egm, respectively, shows that this fea-
ture is associated with Fe ii while in SN 2007bi this line is not
identified.
4.3. C II lines
In Fig. 10, the most prominent features in the red part of the
optical spectrum are likely attributed to C ii λ5890, λ6580, and
λ7234. These lines have been predicted by various SLSN-I mod-
els (Dessart et al. 2012; Mazzali et al. 2016; Dessart 2019) and
have been seen in several SLSNe-Ic (see Sect. 5) but are typi-
cally much weaker than seen in SN 2020zbf. If we consider that
the broad emission component at 6580 Å at +4.3 days spectrum
is associated with Hα rather than C ii, we should be able to detect
a broad Hβ emission at 4861 Å. The region around Hβ is dom-
inated by Fe ii absorption, complicating the analysis, but we do
not find evidence for any broad Balmer emission.
For comparison, the three C ii profiles in the X-shooter spec-
trum at +4.3 days post maximum are shown in Fig. 12. The
emission line profiles appear to be consistent with one another
supporting the hypothesis of C ii at 6580 Å over Hα. The C ii
profiles show an extended tail in the red wing and their peaks are
blueshifted by ∼1000 km s−1 relative to zero rest-frame veloc-
ity, which can be explained by multiple electron scatterings
(Branch & Wheeler 2017; Jerkstrand 2017). However, the lines
have a triangular shape, which could indicate a different forma-
tion process than for the rest of the emission lines in the spec-
trum. The C ii lines do not show obvious P-Cygni profiles since
there is no evident absorption component.
Figure 13 illustrates the evolution of the C ii lines. The
emissions at 5890 Å and 7300 Å become weaker than in the
A20, page 9 of 26
Gkini, A., et al.: A&A, 685, A20 (2024)
3000 4000 5000 6000 7000 8000
Rest-frame wavelength (Å)
0.0
0.5
1.0
1.5
2.0
2.5
3.0
3.5
Sc
al
ed
 F
 (e
rg
 c
m
2  s
1  Å
1 )
CII
CII
CII
CIICII
FeIII
FeIII
FeIII
FeIIIOII
OII
OI
MgII
MgII
MgII
FeII
FeII
MgII?
CIII?
TiIII?
FeIII?
SiIII?
SIII?
Fig. 10. X-shooter spectrum of SN 2020zbf at +4.3 days after maximum light. The spectrum is corrected for MW extinction and is smoothed using
a Savitzky-Golay filter. The original spectrum is shown in lighter colors. The most conspicuous features are labeled. Uncertain line identifications
are denoted with question marks. The ions beneath the spectrum are shown at v = 0, whilst those above have been shifted to match the absorption
component.
3000 4000 5000 6000 7000 8000
Rest-frame wavelength (Å)
0.0
0.2
0.4
0.6
0.8
1.0
Sc
al
ed
 F
 (e
rg
 c
m
2  s
1  Å
1 )
Ca H&K
MgI]
FeII
FeII
CII?CII?
CII?
NaI
FeII?
H ?
Fig. 11. NTT-EFOSC2 spectrum of SN 2020zbf at +43.7 days after
maximum light. The spectrum is corrected for MW extinction and is
smoothed using a Savitzky-Golay filter. The original spectrum is shown
in lighter colors. The colors depict the various elements and the question
marks indicate that their presence in the spectrum is unclear. The ions
beneath the spectrum are shown at the rest-frame wavelength, whilst
those above have been shifted to match the absorption component.
early photospheric phase, whereas the component at 6580 Å
remains strong. The weak emission component at 5890 Å can
be attributed mostly to Na i λλ5890, 5896. The most intriguing
characteristic of the spectrum at +43.7 days is the persistent sig-
10 5 0 5 10
Velocity (103 km s 1)
0.25
0.00
0.25
0.50
0.75
1.00
1.25
1.50
Sc
al
ed
 F
CII 5890 CII 6580 CII 7234
Fig. 12. Three C ii emission lines detected in the X-shooter spectrum
+4.3 days after peak. A local linear continuum subtraction and a nor-
malization to the peak flux has been applied for illustration purposes.
The spectra have been smoothed using a Savitzky-Golay filter while the
original data are shown in a lighter color. The rest-frame wavelength of
the C ii emission is denoted by the dashed vertical line. Note that the
narrow emission component at 6580 Å is the Hα emission from the host
galaxy.
nal at around 6580 Å, which could be attributed to the late-time
appearance of Hα (see Sect. 9.1.2). In Fig. 13 we note that the
smoothed data at −2.4 days for C ii λ5890 and λ6580 and at +2.7
A20, page 10 of 26
Gkini, A., et al.: A&A, 685, A20 (2024)
10 5 0 5 10
Velocity (103 km s 1)
5600 5800 6000
Sc
al
ed
 F
 +
 o
ffs
et
CII 5890
10 5 0 5 10
Velocity (103 km s 1)
6400 6600 6800
Rest-frame wavelength (Å)
 -2.4 d
 +2.7 d
 +4.3 d
 +33.6 d
 +43.7 d
 +57 d
CII 6580
10 5 0 5 10
Velocity (103 km s 1)
7000 7200 7400
CII 7234
Fig. 13. Evolution of the three C ii lines of SN 2020zbf from −2.4 to
+57 rest-frame days. The telluric absorptions are indicated with lighter
gray vertical bands. The spectra have been corrected for extinction. A
smoothing has been applied utilizing the Savitzky-Golay filter and the
original data are shown in lighter colors. The regions corrected for tel-
luric absorptions are indicated with lighter gray vertical lines. A local
linear continuum subtraction and an offset were applied for display pur-
poses. The rest-frame wavelengths of the C ii lines are illustrated with
dashed vertical lines.
days for C ii λ6580 show an absorption component blueshifted
by 5000, 8000, and 6000 km s−1, respectively. However, these
components are not clearly seen in the original data and the
velocity values are not consistent across the three C ii lines. In
addition, any feature blueward C ii λ5890 and λ6580 could be an
artifact from the telluric removal. Thus, we cannot conclusively
consider them as the absorption troughs of a P-Cygni profile.
The absence of absorption troughs in C ii λ7234 as well as in
Fig. 12 strengthens this argument. For the reasons stated above,
we conclude that the C ii lines are more likely pure emission.
4.4. Line velocities
There are two main absorption line features commonly used
to derive the velocity of the photosphere from the spectra
in SLSN-I. The first method, shown in Quimby et al. (2018)
and Gal-Yam (2019a,b), is to use the O ii absorption lines
at 3500–5000 Å at early phases. As mentioned in Sect. 4.1,
the only visible O ii lines in the spectra of SN 2020zbf are
the λ4358 and λ4651, and the absorption troughs are shifted
by −4000 km s−1, which is low compared to other SLSNe-I
(Quimby et al. 2018; Gal-Yam 2019a). Chen et al. (2023b)
studying a sample of 77 SLSNe-I finds the median O ii velocity
of the ZTF sample to be 9700 km s−1 with the distribution going
down to 3000 km s−1. Nicholl et al. (2015) studied a sample of
24 SLSNe-Ic and showed that the median velocity of the SLSNe-
I is 10 500± 3000 km s−1. Our estimated value of 4000 km s−1 is
lower than the median, but since it has been observed in at least
one SLSN in the Chen et al. (2023b) sample, this value is not
unprecedented.
The second method for measuring the photospheric veloc-
ity is to use the Fe ii triplet λλ4923, 5018, 5169 as tracers
(Branch et al. 2002; Nicholl et al. 2015; Modjaz et al. 2016;
Liu et al. 2017). We note that Fe ii λ5169 seems blended
with Fe iii λ5129 in the hot photospheric phase and Fe ii
4700 4800 4900 5000 5100
Rest-frame wavelength (Å)
1.5
1.0
0.5
0.0
0.5
Sc
al
ed
 F
+ 
of
fs
et
FeII
FeII at 12000 km s 1
FeII
FeII at 9000 km s 1 ?
 +4.3 d
 +43.7 d
Fig. 14. Fe ii triplet λλ4923, 5018, 5169 region of the +4.3 and +43.7
day spectra. A linear continuum subtraction and an arbitrary offset has
been applied for illustration purposes. The spectra have been smoothed
using the Savitzky-Golay filter and the original data are shown in
lighter colors. The absorption features that correspond to the blueshift
of the Fe ii lines are marked in blue along with the velocity. At +43.7
days, the absorption feature at 4740 Å could imply that Fe ii is still at
12 000 km s−1.
λ4923 is poorly detected. Thus, we used Fe ii λ5018 to
estimate the velocity. In Fig. 14 a zoomed-in view of the
Fe ii triplet region at +4.3 and +43.7 days post-maximum
is shown. In the X-shooter spectrum at +4.3 days after
peak, the marked absorption components match well with the
Fe ii triplet shifted by −12 000 km s−1 even though the Fe ii
λ4923 is poorly resolved and the Fe ii λ5169 is blended.
This velocity is consistent within the errors with the values
for other SLSNe-I from Nicholl et al. (2015) and Chen et al.
(2023b), the latter of whom found a median velocity of
12 800 km s−1 when studying the Fe ii lines in the ZTF SLSN-I
sample.
We note that the velocity measured from the Fe ii lines is
higher than that measured from the tentative O ii lines, which
is not unprecedented as shown in Quimby et al. (2018) and
Chen et al. (2023b). However, the difference of 8000 km s−1
has never been observed as the average difference between the
estimated velocities using Fe ii and O ii is ∼3000 km s−1 (e.g.,
Chen et al. 2023b). A more likely explanation is that the absorp-
tion tentatively identified as O ii are dominated by different lines.
The low S/N in the cold photospheric phase spectra pre-
vents us from tracking the evolution of the velocity. The strong
feature at 4870 Å in the spectrum at +43.7 days suggests that
the Fe ii λ5018 may be blueshifted by ∼9000 km s−1, in which
case the mismatch of the absorption at 4730 Å can be explained
by a blend of Fe ii with other ions. On the other hand, the
A20, page 11 of 26
Gkini, A., et al.: A&A, 685, A20 (2024)
75 50 25 0 25 50 75 100
Rest-frame days past peak
21
20
19
18
17
M
g A
bs
ol
ut
e 
M
ag
ni
tu
de
SN 2020zbf
PTF10aagc
SN 2017gci
SN 2018bsz
SN 2020wnt iPTF16bad
Fig. 15. Rest-frame absolute magnitude g-band light curve of
SN 2020zbf in comparison with C-rich SLSNe-I from the literature. The
magnitudes are K-corrected and corrected for MW extinction.
absorption feature at 4740 Å could be associated with the Fe ii
λ4923 at −12 000 km s−1, which would result in a constant veloc-
ity over a period of 40 days (Nicholl et al. 2013, 2016, 2015;
Liu et al. 2017).
5. Comparison to other C-rich SLSNe
In Sect. 3.2, we compared the light curve properties of
SN 2020zbf with a homogeneous sample of SLSNe-I, conclud-
ing that the general photometric characteristics of SN 2020zbf
are overall average aside from the fast rise. However, the spec-
tral properties, notably the strong C ii lines, are unusual in
SLSNe though not entirely unprecedented. We inspected the
available spectra in the literature and find a number of objects
also noted for their strong C ii features. We note that the com-
parison is not with the full sample of all the C-rich objects
but rather with a carefully selected sample of publicly avail-
able C-dominated SLSNe-I studied in individual papers as well
as in sample papers. This includes SN 2018bsz (Anderson et al.
2018; Pursiainen et al. 2022), SN 2017gci (Fiore et al. 2021),
SN 2020wnt (Gutiérrez et al. 2022; Tinyanont et al. 2023),
iPTF16bad (Yan et al. 2017a) and PTF10aagc (De Cia et al.
2018; Quimby et al. 2018). In this section, we compare the prop-
erties of SN 2020zbf to these other C-rich objects.
5.1. Light curve comparisons
In Fig. 15, we compare the g-band light curve with g-band light
curves of SLSNe-I that have been found to have strong C ii fea-
tures in their early spectra. The absolute magnitudes of all the
objects are corrected for the MW extinction and cosmological
K-correction. C-rich SLSNe show a peak absolute magnitude
distribution from −20.1 to −21.5 with a mean of −20.8, and
SN 2020zbf is in the upper half of the distribution. SN 2020zbf
fades by 1.8 mag in 79 days (∼0.02 mag d−1) and declines more
slowly than the other SLSNe with C ii features. We notice a
large diversity in the light curves of the C-rich sample, possi-
bly reflecting the variety of mechanisms powering them and the
general diversity among the whole population of SLSNe-I.
5.2. Spectra comparisons
In Fig. 16 (left panel), we compare the spectra of SN 2020zbf
with those of the C-rich sample. We used two reference epochs:
around peak (top) and ∼+40 days post-maximum light (bot-
tom). As mentioned above, the spectrum of SN 2020zbf at +4.3
days presents strong C ii λ5890, λ6580, λ7234, and Fe iii λ4432
and λ5129. These characteristics match well with the near peak
spectrum of PTF10aagc (Quimby et al. 2018) but are stronger in
SN 2020zbf. Even though the C ii profiles of PTF10aagc show
similarities, the peak of the lines is blueshifted by a higher veloc-
ity (3000 km s−1) for that SN than for SN 2020zbf. The C ii lines
in PTF10aagc also show an absorption component in the blue;
in SN 2020zbf absorption components are not obviously visible.
Strong C ii lines are also observed in SN 2017gci (Fiore et al.
2021) though the shapes are different compared to SN 2020zbf
and the 5890 Å emission is absent. The C ii λ6580 and λ7234
in SN 2017gci are broader than in SN 2020zbf and P-Cygni pro-
files are present. SN 2018bsz is a SLSN-I that has been stud-
ied particularly for the strong C ii features it shows in its spec-
tra; however, the C ii features in SN 2018bsz are weaker than
those of SN 2020zbf, and the general structure of the spectrum
is different. SN 2020wnt (Gutiérrez et al. 2022; Tinyanont et al.
2023) and iPTF16bad (Yan et al. 2017a) also present C ii λ6580
and λ7234 lines, yet their spectra are significantly distinct from
those of the other objects in the sample. SN 2020wnt is the only
object of the sample that shows Si ii at 6300 Å and a strong fea-
ture at 5200 Å mostly attributed to Fe ii (see Sect. 5.1).
A large variety regarding the presence of the O ii features
in the spectra exist in the C-rich sample. In SN 2020zbf we
can distinguish only the O ii λ4358 and λ4651 W-shape, which
appears to match well with the O ii lines of PTF10aagc. For
SN 2017gci, even though it exhibits these features, they are
weaker than in SN 2020zbf. In contrast, SN 2018bsz seems to not
show any features at these wavelengths at the considered epoch,
and Gutiérrez et al. (2022) and Tinyanont et al. (2023) demon-
strate that the spectrum of SN 2020wnt lacks the O ii feature. In
the case of iPTF16bad it is unclear whether it presents O ii since
only the feature at 4651 Å is visible. We conclude that apart from
PTF10aagc, which is a relatively good match, none of the C-rich
SLSNe-I in this sample resembles SN 2020zbf in terms of the
shape and the intensity of the C ii or O ii ion lines or resembles
some other object of the sample. This diversity in the spectra
around peak is in concordance with the variety of the light curve
shapes we found in Sect. 5.1.
In Fig. 16 (right panel), we compare SN 2020zbf with a sam-
ple of well-studied SLSNe-I including SN 2015bn (Nicholl et al.
2016), SN 2010gx (Pastorello et al. 2010), LSQ12dlf (Nicholl
et al. 2014), Gaia16apd (Kangas et al. 2017), and SN 2011ke
(Inserra et al. 2013). Typically, the red part of the spectrum of
SLSNe-I shows weak features of O ii and C ii (Gal-Yam 2019a),
which we observe in this sample. However, the C ii lines of
SN 2020zbf have a unique shape and strength that set it apart
from the other objects with a similar general spectral shape. In
addition, the O ii W-shape appears to be present in these SLSNe-
I though it is typically shifted by different velocities. SN 2011ke
is an exception since these features are not present in the spectra
at this particular epoch.
SN 2020zbf has similar properties in the late-time spec-
tra as typical SLSNe-I at these phases and also to SNe-Ic at
around peak (Gal-Yam 2019a). In contrast to SN 2020zbf, which
exhibits no Si ii λ6355, all the comparison objects show a strong
Si ii emission line, but none of the events presents the strong
feature at around ∼6580 Å, which is persistent in all the C-rich
A20, page 12 of 26
Gkini, A., et al.: A&A, 685, A20 (2024)
SN 2020zbf
PTF10aagc
SN 2017gci
SN 2018bsz
SN 2020wnt
iPTF16bad
CII
4000 5000 6000 7000 8000
SN 2020zbf SN 2015bn SN 2010gx
SN 2020zbf
15bn
SN 2010gx LSQ12dlf
LSQ12dlf
Gaia16apd
1ke
CII
4000 5000 6000 7000 8000
Rest-frame wavelength (Å)
Sc
al
ed
 F
 +
 o
ffs
et
Fig. 16. Spectral comparison of SN 2020zbf with SLSNe-I from the literature. Left: SN 2020zbf spectra in comparison with C-rich SLSNe-I at
around peak (top) and ∼30–50 days after peak (bottom). Right: comparison of SN 2020zbf spectra with typical well-studied SLSNe-I at the same
epochs. The spectra are corrected for MW extinction. The spectra of SN 2020zbf, SN 2018bsz, iPTF16bad and LSQ12dlf have been smoothed
using a Savitzky-Golay filter. The vertical gold dashed lines indicate the rest wavelengths of the C ii emission lines.
objects and evolves to Hα in the majority of them (Yan et al.
2015, 2017a; Fiore et al. 2021; Pursiainen et al. 2022).
6. Comparison to model spectra
We compared the X-shooter spectrum of SN 2020zbf at +4.3
days after maximum with synthetic model spectra presented in
Dessart (2019). These spectra are the result of time-dependent
radiative transfer simulations based on magnetar-powered
SNe-Ic using the numerical approach of Dessart (2018). The
SNe are followed with the nonlocal thermodynamic equilibrium
radiative transfer code Cmfgen (Hillier & Dessart 2012) from
day one until one or two years after the explosion (Dessart et al.
2015, 2016, 2017).
Exploring the whole range of models presented in Dessart
(2019), we find that the best match to the observed data is the
model 5p11Bx2. The progenitor of this model is 5p11, which
is described in Yoon et al. (2010) and corresponds to a 60 M
ZAMS star with solar metallicity; it is a primary star of a close
binary system (orbital period of 7 days) that evolves to a WR star
with a final mass of 5.11 M (rather than the 4.95 M quoted
in Yoon et al. 2010; see Dessart et al. 2015). This model has
lost its He-rich envelope and the surface C mass fraction XC,s
is computed to be 0.51 (Dessart et al. 2016). As suggested by
Yoon et al. (2010), the end fate of this progenitor is more likely
to be a SN-Ic.
The explosion is then remapped with v1D (Livne 1993;
Dessart et al. 2010a,b) by means of a piston producing a kinetic
energy of 2.49× 1051 erg (suffix B in 5p11Bx2), which accounts
for the additional kinetic energy to match SLSN ejecta velocities.
This model also uses a strong mixing process in the ejecta (suffix
x2 in 5p11Bx2; Dessart et al. 2015). According to Dessart et al.
(2016), the 5p11Bx2 model corresponds to a H-deficient ejecta
of 3.63 M expanding with 11 500 km s−1 that contains 0.34 M
of He, 0.89 M of C, 1.40 M of O, 0.15 M of Si, and 0.19 M
of 56Ni (mass prior to the decay). The magnetar in this model
A20, page 13 of 26
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3000 4000 5000 6000 7000 8000
Rest-frame wavelength (Å)
0.0
0.1
0.2
0.3
0.4
0.5
0.6
Sc
al
ed
 F
 (e
rg
 s
1  c
m
2  Å
1 )
5p11Bx2
SN 2020zbf
Fig. 17. Comparison of the high-quality X-shooter spectrum of
SN 2020zbf at ∼+4.3 days post maximum (green) with a synthetic spec-
trum at 12 days after peak from Dessart (2019; red). The spectra are
corrected for MW extinction and the X-shooter spectrum is smoothed
using a Savitzky-Golay filter. The original spectrum is shown in lighter
colors, whereas the telluric absorptions are denoted by lighter gray ver-
tical bands.
has an initial rotational energy of 0.4× 1051 erg, a magnetic field
of 3.5× 1014 G, an initial spin period of 7 ms and a spin-down
timescale of 19.1 days.
In Fig. 17, the observed X-shooter spectrum is compared
to the synthetic one. The best match for the observed spec-
trum (+4.3 days past maximum) is obtained for a model at
38.4 days after explosion corresponding to +11.7 days after max-
imum. We highlight that the model spectrum has not been devel-
oped or fine-tuned to match SN 2020zbf, but rather employs
the previously described grid of parameters. However, it qual-
itatively reproduces the overall shape and the main features of
the observed spectrum, despite the fact that the model lines are
slightly blueshifted and the temperature of the model is higher.
In the red part of the spectrum, the model matches almost per-
fectly the profiles associated with Fe ii λ5169 and Fe iii λ5129,
C ii λ5890, λ6580, λ7234, and O i λλ7772, 7774, 7775. In the
synthetic spectra, the C ii form P-Cygni profiles that we do
not observe in the spectra of SN 2020zbf (see the discussion in
Sect. 4.3). Although weak He i λ5875 is present in the model, the
line is mostly dominated by C ii, which is in agreement with our
identification scheme. The computed spectrum in the blue part
shows contributions of Mg ii, C ii, and C iii at 2670 Å, Mg ii at
2880 Å, Fe iii in the region 3000–3600 Å, O ii between 3500 and
5000 Å, and Mg ii at 4330 Å.
The 5p11Bx2 model predicts a rise time of 26.7 days
and a maximum bolometric luminosity of 4.86× 1043 erg s−1
(see Table 1 in Dessart 2019). In the photometric analysis of
SN 2020zbf (see Sects. 3.2 and 3.4), we estimated the rise time
to be .26.4 rest-frame days and the peak bolometric luminosity
to be &8.34± 0.49× 1043 erg s−1. Since 5p11Bx2 has not been
modeled to match SN 2020zbf, we expect some differences in
the light curve properties. However, both the observed upper
limit in the rise time and the modeled value are placed in the
fastest regime of the SLSNe-I. On the other hand, the estimated
peak bolometric luminosity of SN 2020zbf is 1.7 times higher
than that of the model. Other works that find good matches to
Dessart (2019) spectral models also find discrepancies in the
light curve behaviors. For example, Anderson et al. (2018) find
a good match between spectral models and the observed spectra
of SN 2018bsz but the light curve models did not agree with the
observations. They interpreted this fact as due to the model miss-
ing a mechanism that could lead to the slow rise of SN 2018bsz.
Anderson et al. (2018) also point out that the input kinetic
energy of the model and/or the magnetar energy deposition pro-
file can affect both the rise time and the maximum bolometric
luminosity.
In conclusion, the 5p11Bx2 model produces a synthetic spec-
trum that matches the general shape of the X-shooter data very
well and reproduces the majority of the features we observe in
SN 2020zbf. However, the selection of other parameters such as
the mass of the progenitor, the ejecta mass, the amount of C in
the ejecta and the kinetic energy might lead to a light curve and
spectral properties closer to the observed ones. We explore the
magnetar models for the light curve of SN 2020zbf in Sect. 7.2.
7. Light curve modeling
In this section, we model the observed multiband light curves of
SN 2020zbf using the Python-based Modular Open Source Fitter
for Transients (MOSFiT) code (Guillochon et al. 2018) to deter-
mine the most likely powering source. Giving as input the SN
redshift, the luminosity distance, the E(B−V) from the MW, the
light curve data in the different bands and a set of priors listed in
Table 2, MOSFiT outputs the posterior distribution of the mod-
eled light curve parameters. The prior for the ejecta velocity was
set to be a Gaussian distribution with µ = 10 000 km s−1 and
σ = 2000 km s−1.
We selected the Default 56Ni model based on the paper of
Arnett (1982) and Nadyozhin (1994), the Slsn magnetar model
described in Nicholl et al. (2017) and the csm model based on
the semi-analytic model from Chatzopoulos et al. (2013). These
models are fitted using the dynamic nested sampling package
Dynesty (Speagle 2020). The quality of the fit can be quantified
by the likelihood score (logZ; Watanabe 2010) for each model.
7.1. Radioactive source model
We tested the hypothesis that the light curve of SN 2020zbf is
powered by the radioactive decay of 56Ni (Arnett 1982). The
resulting fits to the observed data are plotted in Fig. 18a and the
values of the parameters are listed in Table 2. The corner plot is
presented in Fig. C.1.
By a visual inspection, the model does not reproduce the
multiband light curves very well. It fits the rising part in the c
filter but it does not capture the observed peak magnitudes in the
g and B bands and it does not fit the UVOT data. In addition, the
model light curves decline more slowly in the bluer bands than
the observed ones, and at ∼+180 days the pure 56Ni model fails
to describe the data.
According to the results from MOSFiT, the high luminosity
of SN 2020zbf requires an unrealistic amount of 90% 56Ni in the
7.9 M ejecta. We conclude that radioactive decay is unlikely to
be the main contributor to the luminosity for SN 2020zbf, which
is consistent with the conclusions for other H-poor SLSNe (e.g.,
Chomiuk et al. 2011; Quimby et al. 2011; Inserra et al. 2013;
Nicholl et al. 2017; Chen et al. 2023b).
7.2. Magnetar source model
Assuming that SN 2020zbf is powered by the spin-down of a
rapidly rotating newly formed neutron star (Ostriker & Gunn
1971; Arnett & Fu 1989; Kasen & Bildsten 2010; Chatzopoulos
et al. 2012; Inserra et al. 2013), we ran the slsn model in
A20, page 14 of 26
Gkini, A., et al.: A&A, 685, A20 (2024)
Table 2. Priors and posterior of the parameters fitted with MOSFiT for
the Default , SLSN, and CSM model.
Parameters Priors Best-fit values
56Ni model
f56Ni U (10−3, 1) 0.9+0.05−0.1
Mej (M) L (0.1, 100) 7.9+1−0.5
κ (cm2 g−1) U (0.05, 0.2) 0.06+0.01
κγ (cm2 g−1) L (10−4, 104) 0.05+0.02−0.01
Tmin (K) U (3× 103, 2× 104) 10 178+169−162
vej (km s−1) G (µ = 104, σ = 2× 103) 13 535+918−875
Score (log Z) 86.12
Magnetar-powered model
Mej (M) L (0.1, 100) 1.5+0.6−0.4
MNS (M) U (1, 2) 1.6+0.2−0.2
B (1014 G) U (0.1, 10) 2.2+3.4−1
Pspin (ms) U (1, 10) 5.1+0.5−0.7
κ (cm2 g−1) U (0.05, 0.2) 0.12+0.04−0.04
κγ (cm2 g−1) L (10−4, 104) 0.7+35.6−0.3
Tmin (K) U (3× 103, 2× 104) 10 844+271−387
vej (km s−1) G (µ = 104, σ = 2 × 103) 9603+1304−952
Score (log Z) 129.29
CSM model
Mej (M) L (0.1, 100) 0.2+0.1−0.1
MCSM (M) L (0.1, 30) 4.8+0.7−0.5
ρ (10−12 g cm−3) L (10−15, 10−11) 1.15+0.66−0.34
s U (0, 2) 0.21+0.21−0.13
Tmin (K) U (3× 103, 2× 104) 10 137+122−121
vej (km s−1) G (µ = 104, σ = 2 × 103) 12 343+948−1068
Score (log Z) 115.97
Notes. The U stands for uniform, L for log-uniform, and G for
Gaussian.
MOSFiT setting similar priors to those used by Nicholl et al.
(2017). This model assumes a modified SED accounting for
the line blanketing in the UV part of the SLSN spectra
(Chomiuk et al. 2011). The MOSFiT light curve fits are shown
in Fig. 18b and the resulting values of the posteriors in
Table 2. The corner plot is presented in Fig. C.2. Visually,
the model captures the rise, peak and decline up to ∼+50
days and it fits the data point at ∼+180 days. However, after
∼+50 days, the light curves in the bluer bands decline faster than
the model.
The inferred spin period of 5.14 ms is higher compared to
the typical values (∼2.5 ms) found in the literature, approach-
ing the slowest limit of the observed SLSNe thought to be
powered by a magnetar (Nicholl et al. 2017; Chen et al. 2023b;
Blanchard et al. 2020). Similarly, the value of the magnetic
field (2.18× 1014 G) falls in the highest regime but within
the SLSN magnetar properties (Nicholl et al. 2017; Chen et al.
2023b). The resulting ejecta mass from the model (1.51 M)
is significantly lower than the results of Chen et al. (2023b),
who found that the median ejecta mass is 5.03+4.01−2.39 M. Given
that SN 2020zbf is placed among the fastest rising SLSNe in
the Chen et al. (2023a) sample, we expect the ejecta mass of
SN 2020zbf to be significantly lower than the median value
of the ZTF sample. This is also lower than the findings
of Nicholl et al. (2017), who found a median of 4.8+8.1−2.6 M
studying a more heterogeneous sample of SLSNe-I. However,
in the sample (Chen et al. 2023b), five SLSNe have ejecta
masses of <2 M. In addition, there are two further SLSNe
with Mej < 2 M: PS1-10bzj (Mej = 1.65 M; Lunnan et al.
2013; Nicholl et al. 2017) and SNLS-07D2bv (Mej = 1.55 M;
Howell et al. 2013). Consequently, even though the predicted
ejecta mass for SN 2020zbf is lower than the median value, there
are a few magnetar-powered SLSNe-I with similar derived ejecta
masses.
Based on the results of the model, we estimate the
kinetic energy of the ejecta [Ek = 0.6 × 1051 (Mej/1 M)
(vej/104 km s−1)2 erg, uniform density] to be 0.84 × 1051 erg, the
rotational energy of the magnetar [Erot = 2 × 1052P−2ms(MNS/
1.4 M)3/2 erg; Kasen 2017] to be 1.17 × 1051 erg, the magnetar
spin-down timescale [τspin = 1.3×105(MNS/1.4 M)3/2P2msB−214 s;
Nicholl et al. 2017] to be 10.4 days and the diffusion time [τdiff =
8.6 × 105 (Mej/1 M)3/4 (κ/0.1 cm2 g−1)1/2 (Ek/1051 erg)−1/4 s;
Arnett 1982] to be 17.4 days. The ratio of the two timescales
is ∼1.7, which according to Suzuki & Maeda (2021) indicates
that a large fraction of the rotational energy is contributing to
the SN luminosity, typical for SLSNe. This ratio is similar to
the values found in Chen et al. (2023a,b). The estimated kinetic
energy is lower than the median value of 2.13× 1051 erg pre-
sented in Chen et al. (2023b), which is a result of the lower
ejecta mass. We estimate the total radiated energy of the SN
[Ek = 1051 + 1/2(Erot − Erad) erg; Inserra et al. 2013] to be
1.5× 1051 erg, higher than our estimated integrated radiated
energy of 0.47± 0.014× 1051 erg. This discrepancy is mostly
attributed to the energy emitted by the model in the UV around
peak, which is not captured in our estimate since it is before our
earliest UVOT data point.
We compared the results of MOSFiT with the magne-
tar parameters of the 5p11Bx2 model discussed in Sect. 6.
The model refers to a slow rotating magnetar (7 ms) with
strong magnetic field (3.5× 1014 G), properties that resemble the
MOSFiT results but are placed in the most extreme limits of the
Chen et al. (2023b) magnetar properties. The Mej is 3.63 M in
the 5p11Bx2 model instead of 1.51 M from MOSFiT. The over-
all good match of the magnetar models to the observed spectra
and light curve allows us to favor a low ejecta mass explosion
powered by a magnetar engine.
7.3. CSM source model
We modeled the multiband light curve of SN 2020zbf using the
csm model (Chatzopoulos et al. 2013) in MOSFiT under the
assumption that the light curve is dominated by CSM interac-
tion. The light curve fits are depicted in Fig. 18c and the result-
ing posteriors are listed in Table 2. The corner plot is presented
in Fig. C.3. The model captures the rise but not the peak in the
B and g band, declines more slowly than the data and fails to fit
the single data point at ∼+180 days.
The resulting CSM mass is estimated to be 4.79 M, and the
ejecta mass is 0.22 M, which is even lower than the ejecta mass
inferred by the slsn model. In Chen et al. (2023b) the median
value of the CSM mass is 4.67+6.90−2.56 M, similar to our results
(4.79 M), whereas the ejecta mass derived from the CSM model
for their sample is 11.92+24.98−10.65 M. The wide distribution of the
mass reflects the relatively large amount of events with very low
Mej. In particular, 11 out of 70 SLSNe-I have ejecta masses sim-
ilar to SN 2020zbf and the corresponding CSM masses span the
whole range of the distribution. Three out of these 11 objects
favor the magnetar over the CSM model, six events can be
fit equally well by both the magnetar and CSM models and
two clearly prefer the CSM fit. Hence, there are SLSNe-I that
can be well fit by an CSM model with very low ejecta masses
(<1 M).
A20, page 15 of 26
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50 0 50 100 150 200
Rest-frame days past peak
12
14
16
18
20
22
24
26
28
Ap
pa
re
nt
 M
ag
ni
tu
de
 +
 o
ff
se
t
Ni model
(a)
50 0 50 100 150 200
Rest-frame days past peak
12
14
16
18
20
22
24
26
28
Ap
pa
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nt
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Magnetar model
(b)
50 0 50 100 150 200
Rest-frame days past peak
12
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22
24
26
28
Ap
pa
re
nt
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ag
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 +
 o
ff
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t
CSM model
(c)
Fig. 18. Multiband light curves of SN 2020zbf and their fits with the three different models in MOSFiT: Ni model (a), the magnetar model (b), and
the csm model (c). The colored lines indicate the range of the most likely fits.
It has been demonstrated that using the same CSM struc-
ture, the semi-analytic model from Chatzopoulos et al. (2013)
and hydrodynamic simulations can generate conflicting results
(Moriya et al. 2013, 2018; Sorokina et al. 2016) and the quan-
titative values of the CSM parameters may only be an order
of magnitude approximation. Additionally, it is not clear what
kind of progenitor scenario would result in such a configura-
tion, with an extremely stripped star (0.2 M ejecta) exploding
into massive (>4 M) carbon-oxygen CSM. We conclude that
the light curve shape is also consistent with a CSM model, but
more careful modeling outside the scope of this paper is neces-
sary to explore the possible progenitor and CSM structure.
8. Host galaxy
The left panel of Fig. 1 shows the Legacy Survey image of the
field around SN 2020zbf. A small galaxy is visible at the SN
location; the reported LS photometry corresponds to an abso-
lute magnitude Mg = −17.1 mag (at an effective wavelength of
∼4000 Å), similar to the Large Magellanic Cloud.
8.1. Galaxy SED modeling
We modeled the observed SED (black data points in Fig. 19,
tabulated in Table 1) with the software package Prospector
version 1.1 (Johnson et al. 2021)10. We assumed a Chabrier ini-
tial mass function (Chabrier 2003) and approximated the star
10 Prospector uses the Flexible Stellar Population Synthesis
(fsps) code (Conroy et al. 2009) to generate the underlying
physical model and Python-fsps (Foreman-Mackey et al. 2014)
to interface with Fsps in Python. The Fsps code also accounts for the
contribution from the diffuse gas based on the Cloudy models from
Byler et al. (2017). We used the dynamic nested sampling package
Dynesty (Speagle 2020) to sample the posterior probability.
103 104
Observed wavelength (Å)
19
20
21
22
23
Br
ig
ht
ne
ss
 (m
ag
)
Fig. 19. SED of the host galaxy from 1000 to 60 000 Å (black data
points). The solid line represents the best-fitting model of the SED.
The red squares represent the model-predicted magnitudes. The fitting
parameters are shown in the upper-left corner. The abbreviation “n.o.f.”
stands for the number of filters.
formation history (SFH) as a linearly increasing SFH at early
times followed by an exponential decline at late times [func-
tional form t × exp (−t/τ), where t is the age of the SFH episode
and τ is the e-folding timescale]. The model is attenuated with
the Calzetti et al. (2000) model. The priors of the model param-
eters are set identically to those used by Schulze et al. (2021).
Figure 19 shows the observed SED (black dots) and its best
fit (gray curve). The SED is adequately described by a galaxy
template with a log10 mass of 8.68+0.18−0.22 M, and a star formation
rate of 0.24+0.41−0.12 M yr
−1.
A20, page 16 of 26
Gkini, A., et al.: A&A, 685, A20 (2024)
8.2. Emission line diagnostics
The X-shooter spectrum shows a number of narrow emission
lines from the galaxy seen atop the SN continuum (see Fig. 9).
After calibrating this spectrum to the photometry and correct-
ing for MW extinction, we measured the line fluxes by fitting
Gaussian line profiles. The resulting flux values are listed in
Table 3.
We measured the host galaxy extinction using the Balmer
decrement, finding a value of Hα/Hβ = 3.7 ± 1.0. This is
larger than the theoretical ratio of 2.87 (assuming Case B recom-
bination and a temperature of 10 000 K; Osterbrock & Ferland
2006); however, the noisy Hβ measurement means it is also con-
sistent with the theoretical value within the error. Assuming a
Calzetti et al. (2000) reddening law with RV = 3.1, we estimate
E(B − V)host = 0.22+0.20−0.22 mag. This is consistent with the value
from the SED modeling. Given the large uncertainty in the extinc-
tion and the dwarf nature of the host galaxy, we did not apply any
host galaxy extinction correction to the SN photometry.
No auroral lines are detected in the spectrum of SN 2020zbf,
so we are limited to strong-line metallicity indicators. Given
the large uncertainty in the extinction, we prefer indicators
that use ratios of nearby lines and are therefore less sensitive
to the host extinction, such as N2 or O3N2 (Pettini & Pagel
2004). Using the calibration of Marino et al. (2013) as imple-
mented in the PyMCZ package (Bianco et al. 2016), we calcu-
lated 12 + log(O/H) = 8.31+0.04−0.05 dex using O3N2. Taking the
solar value to be 12 + log(O/H) = 8.69 dex (Asplund et al.
2021), this corresponds to a metallicity Z = 0.4Z.
We converted the Hα flux to a star formation rate following
SFR(M yr−1) = 7.9 × 10−42× L (Hα) erg s−1 (Kennicutt 1998).
This gives a star formation rate of 0.12 M yr−1 after correcting
for extinction, and 0.07 M yr−1 if assuming zero host extinction.
This is slightly lower than, but consistent with, what is inferred
from the SED modeling.
Taken together, the properties of the host galaxy of
SN 2020zbf are quite typical of SLSN-I host galaxies (e.g.,
Lunnan et al. 2014; Leloudas et al. 2015; Angus et al. 2016;
Perley et al. 2016; Chen et al. 2017; Schulze et al. 2018). The
luminosity, mass, and metallicity are all consistent with a dwarf
galaxy and are near the center of the distribution seen for the
hosts of these transients. Thus, even though some of the SN prop-
erties are unusual, the host galaxy environment is not.
9. Discussion
9.1. Powering mechanisms
9.1.1. Can SN2020zbf be powered by a magnetar?
In Sects. 6 and 7, comparisons of the X-shooter spectrum with
tabulated spectra from physical SLSN models, as well as pho-
tometric modeling with MOSFiT, result in a preference toward
a magnetar-powered model with a low ejecta mass (3.63 M
and 1.51 M, respectively). We note that the synthetic spectra
of Dessart (2019) have not been modeled to match SN 2020zbf
or any other SLSN, but qualitatively match SN 2020zbf well,
which strengthens the hypothesis of a low ejecta-mass magnetar-
powered event for SN 2020zbf. The similarity with the mag-
netar properties estimated for typical SLSNe-I in Chen et al.
(2023b) shows that the existence of low ejecta mass SLSNe-I
powered by a magnetar is feasible. The correlation between
the low ejecta mass and the high spin period (P = 5.6 ms) in
SN 2020zbf is in agreement with the studies of Blanchard et al.
(2020), Hsu et al. (2021) and Chen et al. (2023b), and indicates
Table 3. Observed host galaxy emission line fluxes (corrected for MW
extinction).
Line Flux
(10−17 erg s−1cm−2)
[S ii] λ6731 1.61 ± 0.28
[S ii] λ6717 2.27 ± 0.37
[N ii] λ6584 1.18 ± 0.46
Hα λ6563 8.35 ± 0.57
[O iii] λ5007 3.71 ± 0.68
[O iii] λ4959 1.22 ± 0.46
Hβ λ4861 2.25 ± 0.58
[O ii] λ3729 5.65 ± 0.93
[O ii] λ3727 2.63 ± 0.63
that low-mass ejecta SLSNe-I require less central power with
slower-spinning neutron stars.
SN 2020zbf presents unusually strong C ii emission lines in
its early spectra, which likely require the energy from the mag-
netar to be absorbed in the outer layers where the C is more
abundant. The low-mass ejecta could be a critical factor in pro-
ducing these lines, for which the extra energy from the magnetar
is not completely absorbed in the inner O-rich region, diffuses
out and thermally excites the C .
If the ejecta mass was the only key for the presence of strong
C ii lines in the early spectra, we would expect all the low ejecta
mass SLSNe-I to present strong C ii features, and these features
to be absent for higher ejecta masses. In Dessart (2019), some
of the models with higher ejecta masses (Mej = 9.6 M) and
similar magnetar properties as SN 2020zbf still present strong
C ii lines, which Dessart (2019) explains with the presence of a
C-rich shell in the outermost layers of the ejecta. These lines also
vanish as the photosphere recedes into the inner layers where C
is less abundant. In addition, the three objects in the C-rich sam-
ple (SN 2018bsz, SN 2020wnt and SN 2017gci) that have been
fit with a magnetar model, cover a wide range of magnetar prop-
erties (B = 2–6× 1014 G and P = 2.8–7 ms) and show a prefer-
ence for higher ejecta masses (Mej = 9–26 M) than SN 2020zbf.
Thus, we speculate that different scenarios including the ejecta
mass, the magnetar properties and, the mass fraction of C in the
progenitor might be the keys in producing strong C ii lines in the
spectra.
The main discrepancy in the C ii lines between the magnetar
model and our observations is the lack of P-Cygni profiles seen
in our spectra. However, based on the similarities of both the
light curves and spectra, we conclude that the magnetar-powered
explosion of a low-mass, C-rich progenitor star is a viable expla-
nation for the observed properties of SN 2020zbf.
9.1.2. Can SN2020zbf be powered by CSM interaction?
Another mechanism that has been suggested to power SLSN-I
light curves, is CSM interaction (e.g., Chatzopoulos et al.
2012; Sorokina et al. 2016; Wheeler et al. 2017), in which the
kinetic energy of the ejecta is converted into radiation (e.g.,
Zel’dovich & Raizer 1967). The exploration of this scenario is
done under two assumptions; the C ii lines are pure emission
and the energy deposition in the outer layers of the ejecta results
from interaction rather than from a magnetar.
The absence of H and He and the presence of strong C ii
emission features in the spectra indicates a possible interaction
A20, page 17 of 26
Gkini, A., et al.: A&A, 685, A20 (2024)
of the SN ejecta with CO CSM (Woosley et al. 2007; Blinnikov
& Sorokina 2010; Chevalier & Irwin 2011; Chatzopoulos et al.
2012; Quataert & Shiode 2012; Sorokina et al. 2016). Chen et al.
(2023b) estimated between 25% and 44% of SLSNe-I are
preferentially fit by the H-poor CSM models. Sorokina et al.
(2016), using hydro-simulations, succeeded to reproduce the
light curves of SN 2010gx (Pastorello et al. 2010) and PTF09cnd
(Quimby et al. 2011) with CSM interaction based on the ejecta
mass, density structure, explosion energy, expansion of the CSM
and C/O ratio. Moreover, Tolstov et al. (2017), using radiation-
hydrodynamics calculations, modeled PTF12dam and found
that interaction with CSM ejected due to the pulsational pair-
instability mechanism (Barkat et al. 1967) can describe the light
curve.
Suzuki et al. (2020) found a systematic relation between
peak bolometric luminosity and rise time depending on the
kinetic energy, CSM mass, CSM radius and ejecta mass.
SN 2020zbf, if powered by interaction, falls in the regime of
1 M ejecta, 3 M CSM and 0.5 × 1051 erg kinetic energy when
considering its peak bolometric luminosity and the rise time.
Khatami & Kasen (2023) define four light-curve classes based
on the parameters of the CSM/ejecta configuration and a qual-
itative comparison of SN 2020zbf shows that it belongs to the
interior breakout–heavy CSM (MCSM > Mej) class in which
the shock breakout occurs within the CSM (Ginzburg & Balberg
2012; Dessart et al. 2015). In addition, the peak light – rise time
relation in Khatami & Kasen (2023) shows that for SN 2020zbf
the CSM mass is estimated between 1 and 10 M toward the
lower limit.
Janka (2012) showed that neutrino-driven SN explosions
cannot explain kinetic energies larger than ∼2× 1051 erg.
Since the kinetic energy of SN 2020zbf is estimated to be
∼0.8× 1051 erg and assuming that the explosion mechanism in
an interaction-dominated SN is driven by neutrinos, it is possible
the main power source of SN 2020zbf could be CSM interaction.
On the other hand, there are no spectral models that can fit
the line shapes in the SLSN-I spectra in the case of interaction
with CO CSM. The C ii line profiles in the spectra of SN 2020zbf
are asymmetric, showing a conspicuous tail extending to the
red, which indicates a multiple electron scattering effect (e.g.,
Jerkstrand 2017). However, it is unknown what CSM configura-
tion could give rise to the sharp peak in the C ii lines.
As discussed above, the C ii lines in SN 2020zbf vanish with
time except for the line at 6580 Å. Pursiainen et al. (2022) stud-
ied the C-rich SLSN-I SN 2018bsz and concluded that at +24
days post-maximum the C ii at 6580 Å is replaced by Hα. Late
hydrogen emission has also been observed in the C-rich SLSN-I
PTF10aacg (Yan et al. 2015) at +75 days post-maximum. Both
objects have been explained with aspherical CSM interaction.
Late-time interaction with spherical CSM resulting in Hα emis-
sion has been seen in other SLSNe-I (e.g. Yan et al. 2015, 2017a;
Fiore et al. 2021) and Yan et al. (2015) estimate that at least 15%
of SLSNe-I interact with previously ejected H-rich material at
late times.
The evolution of the C ii lines of SN 2020zbf shows a similar
pattern as for SN 2018bsz, but there is no obvious absorption
component blueward of 6580 Å in SN 2020zbf. Furthermore,
there is no indication of Hβ and Hγ appearing, but this can be
limited by the low S/N of the spectra. However, due to the weak-
ness of the other C ii lines, we can assume that this feature is
more likely dominated by another ion rather than C ii. Addition-
ally, the epoch of the studied spectrum falls in the period of a
short (∼15 days) plateau in the LCO light curve between +30
and +45 days past peak, which, if real, could be an indication of
interaction. On the other hand, the shape of the SN 2020zbf light
curve could imply a combination of different mechanisms such
as interaction early on followed by a magnetar power source.
However, due to the possible presence of Hα in the spectra,
the plateau is more likely to be associated with interaction with
H-rich material located at .4.5× 1015 cm (vej × 43.7 days; due to
the low coverage this value is set as an upper limit). The multi-
ple CSM layers are consistent with the pulsational pair instability
model (Woosley 2017) that has been suggested as a mechanism
for SLSNe (Woosley et al. 2007).
In reality, SN 2020zbf could be powered by both a magnetar
and CSM interaction. The preferred magnetar 5p11Bx2 model
could well be correct giving a very good match with SN 2020zbf
spectra with a preference to low ejecta mass, but there could
also be some CO CSM contributing mainly to the line profiles
(sharply peaked emission C ii lines) and slightly to the light
curve.
One possible way to break the degeneracy between the
two powering mechanisms would be with radio observations.
Both CSM interaction and magnetars are expected to pro-
duce detectable radio emission, but the timescales should be
very different. Radio emission from CSM interaction should
be produced within the first few years, while with a magne-
tar engine, the emission should not be detectable until around a
decade (Omand et al. 2018). Two SLSNe have been detected in
radio so far, PTF10hgi (Eftekhari et al. 2019; Law et al. 2019;
Mondal et al. 2020; Hatsukade et al. 2021), which is consis-
tent with the magnetar model, and SN 2017ens (Margutti et al.
2023), which is consistent with CSM interaction. Omand et al.
(2018) also show that for the magnetar model, SLSNe with lower
ejecta masses, similar to SN 2020zbf, produce brighter radio
emission at earlier times, which may make SN 2020zbf a good
candidate for follow-up observations.
9.2. Progenitor
Previous studies in the literature have suggested that the host
galaxies of SLSNe-I have low metallicity and high star forma-
tion rate (e.g., Leloudas et al. 2015) while Chen et al. (2017)
found that host galaxies of SLSN-I progenitors have a metal-
licity cut-off at 0.5 solar metallicity (see also Perley et al. 2016;
Schulze et al. 2018).
Given the spectroscopic similarities of SLSNe-I with
stripped-envelope SNe (Pastorello et al. 2010), there are two
main scenarios for the progenitor stars; single WR stars
(e.g., Georgy et al. 2009) and stars in binary systems (e.g.,
Nomoto et al. 1990; Yoon et al. 2010). The possibly identified
progenitors of SN-Ib iPTF13bvn (Cao et al. 2013) and SN-Ic
SN 2017ein (Van Dyk et al. 2018) have been discussed in the
context of both these scenarios (Groh et al. 2013; Bersten et al.
2014; Van Dyk et al. 2018). In addition, Pursiainen et al. (2022)
argue that for SN 2018bsz both a single WR star and a star in
a binary system are possible. Nicholl et al. (2017) show that
SLSNe-I result from CO cores with MCO ≥ 4 M, which cor-
responds to MZAMS ≥ 20 M (Yoon et al. 2010) comparing with
the simulations of Yoon et al. (2006) at comparable metallicities.
Rapid rotation might be the key to magnetar formation since
the angular momentum needed to form a millisecond magne-
tar requires CO cores with initial rotational velocities >200–
300 km s−1 (Yoon et al. 2006). On the other hand, de Mink et al.
(2013) showed that binarity could also supply the necessary
angular momentum to the stars, either through merging or
Roche lobe overflow. Nicholl et al. (2017) argue that the rapid
rotation of SLSN-I progenitors plays a crucial role, with the
A20, page 18 of 26
Gkini, A., et al.: A&A, 685, A20 (2024)
progenitor mass and the low metallicity being consequences of
this.
SN 2020zbf has been characterized as type Ic due to the
absence of H and He in the spectra, which is associated with
H and possibly He envelope losses in the progenitor star by stel-
lar winds or mass transfer in a binary system before the explo-
sion. In the case of a magnetar-powered event, the dominance
of C instead of O in the spectra would potentially imply an
initially high fraction of C in the progenitor star. The best-fit
5p11Bx2 model (see Sect. 6) from Dessart (2019) corresponds
to an initial massive 60 M ZAMS star in a binary system that
has stripped the hydrogen envelope through mass transfer during
core hydrogen burning. It has lost the He envelope, and its sur-
face is within the C-rich part of the CO core (Yoon et al. 2010;
Dessart et al. 2015). This progenitor could be associated with a
carbon-type WC star that ends with a final mass of 5.11 M and
it can create a SN-Ic with 3.63 M ejecta from which 0.89 M
is C and 1.40 M is O. The appearance of strong C ii lines in
the spectra of SN 2020zbf around peak and their vanishing in
the late-time spectra is consistent with the stratification of the
5p11 model, in which the outer ejecta are enriched in C. The
modeling of the light curve with MOSFiT results in a progenitor
pre-SN star of 3.1 M (Mej + MNS). Both the light curve mod-
eling and the spectral comparison show a low-mass progenitor
star at explosion with masses between 3 and 5 M, lower than
the median value of 6.83+4.04−2.45 M in Chen et al. (2023b) but not
unprecedented since there are a few progenitor systems in the
ZTF sample with similar values to SN 2020zbf.
In the case of the CSM interaction, the early photospheric
spectra could point toward an interaction with CO material
expelled before the explosion. The comparison with different
CSM models shows that the progenitor star (Mej + MCSM +
1.4 M, a typical value for a neutron star; Lattimer & Prakash
2007) is between 6 and 10 M (6.4 M; Chatzopoulos et al.
2013, 5.4 M; Suzuki et al. 2020, <10 M; Khatami & Kasen
2023), lower than the value of 17.92+24.11−9.82 M in Chen et al.
(2023b) but still not extraordinary. However, taking into account
the preference of all the models in low ejecta mass (0.22–3 M),
the fast rise time in the SN 2020zbf’s light curve, and the fact that
the CSM might have been expelled well in advance and should
not be considered in the mass of the pre-SN progenitor, we con-
clude that the progenitor star at explosion is estimated to be in
the range of 2–5 M.
Anderson et al. (2018) studied the SLSN-I SN 2018bsz in
particular for the strong C ii and they concluded that these lines
are produced by a magnetar-powered explosion of a massive C-
rich WR progenitor (11.4 M pre-SN mass) in a binary system
by comparing qualitatively with magnetar models from Dessart
(2019). Considering the good match of the 5p11Bx2 model with
the observed spectrum of SN 2020zbf, a possible progenitor sys-
tem for SN 2020zbf could be as well a WR in a binary system.
Binary interaction could also remove the H-envelope and form
a H-rich CSM around SN 2020zbf, with which the ejecta will
interact at later times.
9.3. Implications of C-rich SLSNe-I
In this work, we have compared SN 2020zbf with H-poor SLSNe
that show C ii in emission. The comparison is not with a statisti-
cally complete sample of all C-dominated objects but rather with
a careful selection of published C-rich objects. Anderson et al.
(2018) suggest that the C ii lines in SN 2018bsz stem from the
large fraction of C in the progenitor model. Even though we
reach the same conclusion for SN 2020zbf, our SN presents
stronger lines of C ii, possibly resulting from the lower ejecta
mass. On the other hand, PTF10aagc seems to be the best spec-
tral match for SN 2020zbf, but no study for the powering mech-
anism of that SN has been done so far.
Fiore et al. (2021) fit SN 2017gci photometry to synthetic
light curves and model the nebular spectra and found that
SN 2017gci can be fit well by both a magnetar model and CSM
interaction with an ejecta mass of 10 M, although they do not
comment on the progenitor’s composition. In SN 2020wnt, the
high ejecta mass (∼26 M; Tinyanont et al. 2023) hides the mag-
netar at early times and it was not revealed until the nebular
phase. The progenitor star for that SN was suggested to be a
massive rotating star with high abundances of C in its outer lay-
ers (Tinyanont et al. 2023, their Fig. A.2).
The difference in the properties of the SLSN powering mech-
anisms reflects the large diversity in the light curves and the
spectra of these objects. Carbon is always present in the out-
ermost CO layers of the ejecta, and thus, we expect to see it in
the spectra of SLSNe-I. However, the presence of lines stronger
than for the typical SLSNe-I could imply a distinct mechanism
and/or a specific composition in the ejecta. In any case, the pres-
ence of strong C ii could indicate a higher amount of C in the
ejecta. A detailed analysis of a statistically meaningful C-rich
SLSN-I sample is required in the future to understand the mech-
anisms for the formation of C ii lines and the properties of their
progenitor systems.
10. Conclusions
In this work we have studied the H-poor SLSN SN 2020zbf,
analyzing photometric data from ATLAS, LCO, and UVOT and
spectra taken with NTT+EFOSC and the VLT+X-shooter instru-
ment. Our main conclusions are as follows:
– SN 2020zbf is a fast-rising SLSN-Ic with a rise time of
.26.4 days from the explosion and a peak magnitude of
Mg = −21.2 mag. The rise time is on the faster end of the
distribution for SLSNe-I.
– The early spectra of SN 2020zbf present three strong C ii
lines that are not typically seen in normal SLSNe-I.
– Both the light curve modeling and the comparison with syn-
thetic spectra are consistent with a magnetar-powered SN of
a C-rich star with an ejecta mass of about 1.5–3 M.
– Alternatively and/or additionally, we argue that the shape and
the strength of the C ii lines can also be attributed to an inter-
action between low-mass ejecta and a dense CO CSM.
– Based on the above modeling, the progenitor is estimated to
have a pre-SN mass of between 2 and 5 M.
– A potential late-time Hα emission that is accompanied by a
knee in the LCO light curve could be an indication of inter-
action with H-rich CSM.
– The host galaxy has a mass of log(M/M) = 8.68+0.18−0.22, a
star formation rate of 0.24+0.41−0.12 M yr
−1, and a metallicity of
0.4Z, similar to typical SLSN-I host galaxies.
– There is large variety in the light curves and the spectra of
SLSNe with strong C ii, suggesting that many different con-
figurations can result in such a spectral signature.
This object illustrates the challenges in classifying SLSNe-I by
demonstrating how diverse even the SLSN class can be, as well
as the implications for both progenitor populations and explo-
sion mechanisms. More sophisticated tools for light curve and
spectra modeling are required to explain the peculiarities of sim-
ilar objects.
A20, page 19 of 26
Gkini, A., et al.: A&A, 685, A20 (2024)
Acknowledgements. The authors would like to thank the anonymous referee
for their comments and suggestions. Based on observations collected at the
European Organisation for Astronomical Research in the Southern Hemisphere,
Chile, as part of ePESSTO+ (the advanced Public ESO Spectroscopic Survey for
Transient Objects Survey). ePESSTO+ observations were obtained under ESO
program IDs 106.216 and 108.220C (PI: Inserra). LCO data have been obtained
via an OPTCON proposal (IDs: OPTICON 20A/015, 20B/003, 21B/001, 22B/
002) and as part of the Global Supernova Project. The OPTICON project has
received funding from the European Union’s Horizon 2020 research and inno-
vation programme under grant agreement No 730890. This work makes use of
data from the Las Cumbres Observatory global network of telescopes. The LCO
group is supported by NSF grants AST-1911151 and AST-1911225. This work
has made use of data from the Asteroid Terrestrial-impact Last Alert System
(ATLAS) project. ATLAS is primarily funded to search for near earth asteroids
through NASA grants NN12AR55G, 80NSSC18K0284, and 80NSSC18K1575;
byproducts of the NEO search include images and catalogs from the survey area.
The ATLAS science products have been made possible through the contribu-
tions of the University of Hawaii Institute for Astronomy, the Queen’s Uni-
versity Belfast, the Space Telescope Science Institute, and the South African
Astronomical Observatory. R. Lunnan is supported by the European Research
Council (ERC) under the European Union’s Horizon Europe research and inno-
vation programme (grant agreement No. 10104229 - TransPIre). S. Schulze
acknowledges support from the G.R.E.A.T. research environment, funded by
Vetenskapsrådet, the Swedish Research Council, project number 2016-06012.
M. Nicholl is supported by the European Research Council (ERC) under the
European Union’s Horizon 2020 research and innovation programme (grant
agreement No. 948381) and by UK Space Agency Grant No. ST/Y000692/1.
T.-W. Chen acknowledges the Yushan Young Fellow Program by the Min-
istry of Education, Taiwan for the financial support. M. Pursiainen acknowl-
edges support from a UK Research and Innovation Fellowship (MR/T020784/1).
T. E. Müller-Bravo acknowledges financial support from the Spanish Ministerio
de Ciencia e Innovación (MCIN), the Agencia Estatal de Investigación (AEI)
10.13039/501100011033, and the European Union Next Generation EU/PRTR
funds under the 2021 Juan de la Cierva program FJC2021-047124-I and the
PID2020-115253GA-I00 HOSTFLOWS project, from Centro Superior de Inves-
tigaciones Científicas (CSIC) under the PIE project 20215AT016, and the pro-
gram Unidad de Excelencia María de Maeztu CEX2020-001058-M. This work
was funded by ANID, Millennium Science Initiative, ICN12_009. G. Pignata
acknowledges support from ANID through Millennium Science Initiative Pro-
grams ICN12_009. C. P. Gutiérrez acknowledges financial support from the
Secretary of Universities and Research (Government of Catalonia) and by the
Horizon 2020 Research and Innovation Programme of the European Union under
the Marie Skłodowska-Curie and the Beatriu de Pinós 2021 BP 00168 pro-
gramme, from the Spanish Ministerio de Ciencia e Innovación (MCIN) and
the Agencia Estatal de Investigación (AEI) 10.13039/501100011033 under the
PID2020-115253GA-I00 HOSTFLOWS project, and the program Unidad de
Excelencia María de Maeztu CEX2020-001058-M.
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Appendix A: Photometric data
Table A.1. Photometric data of SN 2020zbf.
MJD Phasea Filter Magnitude Instrumentb
(days) (mag)
59137.5 -22.60 c 20.89 ± 0.35 ATLAS
59145.4 -15.90 c 19.65 ± 0.10 ATLAS
59147.3 -14.23 o 19.73 ± 0.21 ATLAS
59157.4 -5.86 o 18.96 ± 0.18 ATLAS
59161.4 -2.51 o 18.87 ± 0.08 ATLAS
59163.4 -0.84 o 18.93 ± 0.07 ATLAS
59164.8 0 B 18.64 ± 0.20 LCO
59164.8 0 g 18.60 ± 0.07 LCO
59164.8 0 V 18.66 ± 0.06 LCO
59164.8 0 r 18.75 ± 0.10 LCO
59164.8 0 i 18.90 ± 0.07 LCO
59169.3 4.19 c 18.79 ± 0.05 ATLAS
59170.8 5.02 B 18.75 ± 0.14 LCO
59170.8 5.02 g 18.76 ± 0.09 LCO
59170.8 5.02 V 18.85 ± 0.14 LCO
59170.8 5.02 r 18.97 ± 0.04 LCO
59170.8 5.02 i 19.28 ± 0.13 LCO
59176.1 10.04 B 18.87 ± 0.13 LCO
59176.1 10.04 g 18.99 ± 0.06 LCO
59176.1 10.04 V 18.98 ± 0.11 LCO
59176.1 10.04 r 19.04 ± 0.08 LCO
59176.1 10.04 i 19.14 ± 0.08 LCO
59178.01 11.73 UVW1 20.50 ± 0.15 Swift/UVOT
59178.01 11.73 U 19.04 ± 0.11 Swift/UVOT
59178.01 11.73 B 18.99 ± 0.14 Swift/UVOT
59178.01 11.73 UVW2 21.67 ± 0.20 Swift/UVOT
59178.01 11.73 V 18.90 ± 0.24 Swift/UVOT
59178.01 11.73 UVM2 21.27 ± 0.15 Swift/UVOT
59179.66 13.11 UVW1 20.62 ± 0.16 Swift/UVOT
59179.66 13.11 U 19.14 ± 0.11 Swift/UVOT
59179.66 13.11 B 19.12 ± 0.15 Swift/UVOT
59179.67 13.11 UVW2 21.33 ± 0.17 Swift/UVOT
59179.67 13.11 V 18.81 ± 0.22 Swift/UVOT
59179.67 13.11 UVM2 21.36 ± 0.16 Swift/UVOT
59181.0 14.23 o 18.89 ± 0.17 ATLAS
59183.0 15.90 o 19.09 ± 0.22 ATLAS
59185.0 17.58 o 19.16 ± 0.16 ATLAS
59185.38 17.90 UVW1 21.20 ± 0.24 Swift/UVOT
59185.38 17.90 U 19.31 ± 0.13 Swift/UVOT
59185.38 17.90 B 19.27 ± 0.18 Swift/UVOT
59185.38 17.90 UVW2 22.62 ± 0.38 Swift/UVOT
59185.38 17.90 V 18.66 ± 0.22 Swift/UVOT
59185.39 17.90 UVM2 21.52 ± 0.19 Swift/UVOT
59186.1 18.41 B 19.16 ± 0.09 LCO
59186.1 18.41 g 19.16 ± 0.07 LCO
59186.1 18.41 V 19.14 ± 0.07 LCO
59186.1 18.41 r 19.26 ± 0.06 LCO
59186.1 18.41 i 19.38 ± 0.05 LCO
59187.3 19.25 o 19.29 ± 0.14 ATLAS
59191.3 22.60 o 19.65 ± 0.14 ATLAS
59193.9 24.27 B 19.28 ± 0.11 LCO
59193.9 24.27 g 19.30 ± 0.16 LCO
59193.9 24.27 V 19.30 ± 0.13 LCO
59193.9 24.27 r 19.11 ± 0.15 LCO
59193.9 24.27 i 19.34 ± 0.19 LCO
59193.27 24.50 UVW1 21.34 ± 0.38 Swift/UVOT
59193.27 24.50 U 19.58 ± 0.22 Swift/UVOT
Table A.1. continued.
MJD Phasea Filter Magnitude Instrumentb
(days) (mag)
59193.27 24.50 B 19.47 ± 0.30 Swift/UVOT
59193.28 24.50 UVW2 22.46 ± 0.53 Swift/UVOT
59193.28 24.50 V 18.73 ± 0.34 Swift/UVOT
59193.28 24.50 UVM2 21.85 ± 0.33 Swift/UVOT
59193.3 24.27 c 19.27 ± 0.10 ATLAS
59195.3 25.95 c 19.50 ± 0.10 ATLAS
59197.3 27.62 c 19.43 ± 0.12 ATLAS
59201.2 30.97 B 19.57 ± 0.08 LCO
59201.2 30.97 g 19.47 ± 0.04 LCO
59201.2 30.97 V 19.35 ± 0.07 LCO
59201.2 30.97 r 19.39 ± 0.08 LCO
59201.2 30.97 i 19.66 ± 0.07 LCO
59203.3 32.64 o 19.45 ± 0.09 ATLAS
59205.3 34.32 o 19.59 ± 0.17 ATLAS
59207.5 35.99 B 19.65 ± 0.19 LCO
59207.5 35.99 V 19.34 ± 0.22 LCO
59207.5 35.99 o 19.41 ± 0.12 ATLAS
59211.3 39.34 o 19.73 ± 0.27 ATLAS
59213.3 41.01 o 19.49 ± 0.28 ATLAS
59215.2 42.69 g 19.60 ± 0.14 LCO
59215.2 42.69 V 19.49 ± 0.17 LCO
59215.2 42.69 r 19.48 ± 0.12 LCO
59215.2 42.69 i 19.42 ± 0.12 LCO
59219.3 46.04 c 19.52 ± 0.09 ATLAS
59221.1 47.71 B 19.83 ± 0.15 LCO
59221.1 47.71 g 19.65 ± 0.28 LCO
59221.1 47.71 V 19.52 ± 0.09 LCO
59221.1 47.71 r 19.53 ± 0.20 LCO
59221.1 47.71 i 19.60 ± 0.10 LCO
59227.1 52.73 B 19.98 ± 0.14 LCO
59227.1 52.73 g 19.87 ± 0.10 LCO
59227.1 52.73 V 19.66 ± 0.13 LCO
59227.1 52.73 r 19.65 ± 0.15 LCO
59227.1 52.73 i 19.74 ± 0.21 LCO
59227.1 52.73 c 19.87 ± 0.12 ATLAS
59229.2 54.40 c 19.75 ± 0.13 ATLAS
59231.3 56.08 c 19.43 ± 0.10 ATLAS
59233.1 57.76 B 20.12 ± 0.15 LCO
59233.1 57.76 g 20.01 ± 0.11 LCO
59233.1 57.76 V 19.79 ± 0.13 LCO
59233.1 57.76 r 19.77 ± 0.19 LCO
59233.1 57.76 i 19.85 ± 0.25 LCO
59239.1 62.78 g 20.18 ± 0.21 LCO
59239.1 62.78 V 19.96 ± 0.22 LCO
59239.1 62.78 r 19.90 ± 0.18 LCO
59239.1 62.78 i 20.10 ± 0.12 LCO
59245.2 67.80 o 20.07 ± 0.26 ATLAS
59246.8 68.64 B 20.54 ± 0.28 LCO
59246.8 68.64 g 20.38 ± 0.15 LCO
59246.8 68.64 V 20.19 ± 0.22 LCO
59258.8 78.68 V 20.39 ± 0.28 LCO
59258.8 78.68 r 20.29 ± 0.21 LCO
59266.4 85.38 r 20.36 ± 0.20 LCO
59377.4 178.29 r 21.62 ± 0.23 LCO
59377.4 178.29 i 21.79 ± 0.26 LCO
Notes. The photometry is reported on the AB system and is not cor-
rected for reddening. This table is available in machine readable form.
Multiple exposures on any given night are averaged to give the val-
ues presented here. aRest-frame relative to the g-band maximum (MJD
59 164.8). bUVOT photometry is not host galaxy subtracted.
A20, page 22 of 26
Gkini, A., et al.: A&A, 685, A20 (2024)
Appendix B: Spectroscopic data
Table B.1. SN 2020zbf spectroscopic observations.
UT date MJD Phasea Telescope + Exposure Grism Wavelength range
(days) Instrument (s) (Å)
20201109 59162 -2.37 NTT + EFOSC2 900 Gr#13 3650–9250
20201116 59168 2.65 NTT +EFOSC2 1800 Gr#11 + Gr#16 3345–9995
20201118 59170.8 4.33 VLT + X-shooter 2400 – 3000–24800
20201208 59190 21.93 NTT + EFOSC2 2700 Gr#13 3650–9250
20201223 59205 33.62 NTT + EFOSC2 2700 Gr#13 3650–9250
20210104 59217 43.67 NTT + EFOSC2 2700 Gr#13 3650–9250
20210123 59233.1 57.06 NTT + EFOSC2 2700 Gr#13 3650–9250
Notes. aRest-frame relative to the g-band maximum (MJD 59164.8).
A20, page 23 of 26
Gkini, A., et al.: A&A, 685, A20 (2024)
Appendix C: MOSFiT results
log fNi = 0.05+0.030.05
0.0
6
0.0
7
0.0
8
0.0
9
(c
m
2
g
1 )
(cm2 g 1) = 0.06+0.010.00
1.5
0
1.3
5
1.2
0
1.0
5
lo
g
(c
m
2
g
1 )
log (cm2 g 1) = 1.26+0.090.10
0.8
4
0.9
0
0.9
6
1.0
2
1.0
8
lo
g
M
ej
(M
)
log Mej (M ) = 0.90+0.050.03
17
18
19
20
lo
g
n H
,h
os
t
log nH, host = 18.23+1.061.33
99
00
10
20
01
05
00
10
80
0
T m
in
(K
)
Tmin (K) = 10179.98+168.95162.42
0.2
0
0.1
5
0.1
0
0.0
5
log fNi
10
50
01
20
00
13
50
01
50
00
16
50
0
v e
j(k
m
s
1 )
0.0
6
0.0
7
0.0
8
0.0
9
(cm2 g 1)
1.5
0
1.3
5
1.2
0
1.0
5
log (cm2 g 1)
0.8
4
0.9
0
0.9
6
1.0
2
1.0
8
log Mej (M )
17 18 19 20
log nH, host
99
00
10
20
0
10
50
0
10
80
0
Tmin (K)
10
50
0
12
00
0
13
50
0
15
00
0
16
50
0
vej (km s 1)
vej (km s 1) = 13534.96+918.25874.61
Fig. C.1. 1D and 2D posterior distributions of the default 56Ni model parameters from the MOSFiT model. Median and 1σ of the best fit values
are marked and labeled.
A20, page 24 of 26
Gkini, A., et al.: A&A, 685, A20 (2024)
B (1014 G) = 2.18+3.441.01
1.2
1.4
1.6
1.8
M
NS
(M
)
MNS (M ) = 1.62+0.210.27
4.2
4.8
5.4
6.0
P s
pi
n
(m
s)
Pspin (ms) = 5.14+0.500.65
0.0
6
0.0
9
0.1
2
0.1
5
0.1
8
(c
m
2
g
1 )
(cm2 g 1) = 0.12+0.040.04
0
1
2
3
lo
g
(c
m
2
g
1 )
log (cm2 g 1) = 0.18+1.740.26
0.2
0.0
0.2
0.4
lo
g
M
ej
(M
)
log Mej (M ) = 0.18+0.140.13
17
18
19
20
lo
g
n H
,h
os
t
log nH, host = 17.91+1.201.09
10
20
01
08
00
11
40
01
20
00
T m
in
(K
)
Tmin (K) = 10844.27+270.70386.72
2 4 6 8
B (1014 G)
75
00
90
00
10
50
0
12
00
0
13
50
0
v e
j(k
m
s
1 )
1.2 1.4 1.6 1.8
MNS (M )
4.2 4.8 5.4 6.0
Pspin (ms)
0.0
6
0.0
9
0.1
2
0.1
5
0.1
8
(cm2 g 1)
0 1 2 3
log (cm2 g 1)
0.2 0.0 0.2 0.4
log Mej (M )
17 18 19 20
log nH, host
10
20
0
10
80
0
11
40
0
12
00
0
Tmin (K)
75
00
90
00
10
50
0
12
00
0
13
50
0
vej (km s 1)
vej (km s 1) = 9602.77+1303.64952.29
Fig. C.2. 1D and 2D posterior distributions of the slsn magnetar model parameters from the MOSFiT model. Median and 1σ of the best fit values
are marked and labeled.
A20, page 25 of 26
Gkini, A., et al.: A&A, 685, A20 (2024)
log MCSM = 0.68+0.060.05
0.8
0.4
0.0
0.4
0.8
lo
g
M
ej
(M
)
log Mej (M ) = 0.65+0.210.18
17
18
19
20
lo
g
n H
,h
os
t
log nH, host = 17.65+0.750.83
12
.3
12
.0
11
.7
11
.4
11
.1
lo
g
log  = 11.94+0.200.15
0.2
5
0.5
0
0.7
5
1.0
0
s
s = 0.21+0.210.13
97
50
10
00
01
02
50
10
50
0
T m
in
(K
)
Tmin (K) = 10136.91+121.67121.35
0.3
0
0.4
5
0.6
0
0.7
5
log MCSM
90
00
10
50
01
20
00
13
50
0
v e
j(k
m
s
1 )
0.8 0.4 0.0 0.4 0.8
log Mej (M )
17 18 19 20
log nH, host
12
.3
12
.0
11
.7
11
.4
11
.1
log
0.2
5
0.5
0
0.7
5
1.0
0
s
97
50
10
00
0
10
25
0
10
50
0
Tmin (K)
90
00
10
50
0
12
00
0
13
50
0
vej (km s 1)
vej (km s 1) = 12342.65+948.451068.34
Fig. C.3. 1D and 2D posterior distributions of the csm model parameters from the MOSFiT model. Median and 1σ of the best fit values are
marked and labeled.
A20, page 26 of 26