# Unveiling the Nature of SN 2011fh: a Young and Massive Star Gives Rise to a
Luminous SN 2009ip-like Event

Thallis Pessi E-mail<EMAIL_ADDRESS>Núcleo de Astronomía,
Universidad Diego Portales, Av. Ejército 441, Santiago, Chile  Jose L. Prieto
Núcleo de Astronomía, Universidad Diego Portales, Av. Ejército 441, Santiago,
Chile  Millennium Institute of Astrophysics, Santiago, Chile  Berto Monard
Klein Karoo Observatory, Calitzdorp, South Africa  Christopher S. Kochanek
Department of Astronomy, The Ohio State University, 140 West 18th Avenue,
Columbus, OH 43210, USA  Center for Cosmology and Astroparticle Physics, The
Ohio State University, 191 W. Woodruff Avenue, Columbus, OH 43210, USA  Greg
Bock  Runaway Bay, Gold Coast, Queensland, Australia  Andrew J. Drake
Division of Physics, Mathematics, and Astronomy, California Institute of
Technology, Pasadena, CA 91125, USA  Ori D. Fox Space Telescope Science
Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA  Stuart Parker
Parkdale Observatory, 225 Warren Road, RDl Oxford, Canterbury 7495, New
Zealand  Heloise F. Stevance Department of Physics, The University of
Auckland, Private Bag 92019, Auckland, New Zealand

###### Abstract

In recent years, many Type IIn supernovae have been found to share striking
similarities with the peculiar SN 2009ip, whose true nature is still under
debate. Here, we present 10 years of observations of SN 2011fh, an interacting
transient with spectroscopic and photometric similarities to SN 2009ip. SN
2011fh had a M${}_{r}\sim-16$ mag brightening event, followed by a brighter
M${}_{r}\sim-18$ mag luminous outburst in August 2011. The spectra of SN
2011fh are dominated by narrow to intermediate Balmer emission lines
throughout its evolution, with P Cygni profiles indicating fast-moving
material at $\sim 6400\ \textrm{km s}^{-1}$. HST/WFC3 observations from
October of 2016 revealed a bright source with M${}_{F814W}\approx-13.3$ mag,
indicating that we are seeing the ongoing interaction of the ejecta with the
circumstellar material or that the star might be going through an eruptive
phase five years after the luminous outburst of 2011. Using HST photometry of
the stellar cluster around SN 2011fh, we estimated an age of $\sim 4.5$ Myr
for the progenitor, which implies a stellar mass of $\sim 60$ M⊙, using
single-star evolution models, or a mass range of $35-80\ \textrm{M}_{\odot}$,
considering a binary system. We also show that the progenitor of SN 2011fh
exceeded the classical Eddington limit by a large factor in the months
preceding the luminous outburst of 2011, suggesting strong super-Eddington
winds as a possible mechanism for the observed mass-loss. These findings favor
an energetic outburst in a young and massive star, possibly a luminous blue
variable.

supernovae: general - supernovae: individual (SN 2009ip, SN 2011fh) - stars:
massive - stars: mass-loss

††software:  HOKI (Stevance et al., 2020a, b), BPASS (Eldridge et al., 2017;
Stanway & Eldridge, 2018), FIREHOSE (Simcoe et al., 2013), IRAF (Tody, 1986,
1993), Astropy (Astropy Collaboration et al., 2013, 2018), Photutils (Bradley
et al., 2019), PYPHOT (https://mfouesneau.github.io/docs/pyphot/), DUSTY
(Ivezic & Elitzur, 1997), DOLPHOT (Dolphin, 2016), SUPERBOL (Version 1.7;
Nicholl, 2018), ifuanalysis
(https://ifuanal.readthedocs.io/en/latest/index.html), STARLIGHT (Cid
Fernandes et al., 2005; Mateus et al., 2006; Asari et al., 2007)

##  1 Introduction

Most stars with a zero age main sequence (ZAMS) mass higher than 8 M⊙ are
thought to end their lives as a core-collapse supernova (CCSN). An important
fraction of these explosions happen inside a very dense environment, formed by
the material ejected by the star in its previous mass-loss history. If the
circumstellar material (CSM) is sufficiently dense and dominated by hydrogen,
hydrogen emission lines with multiple velocity components appear, leading to a
Type IIn supernova (SN, Schlegel, 1990; Filippenko, 1997; Ransome et al.,
2021). If, on the other hand, the CSM is helium-rich, a Type Ibn SN is
observed (Pastorello et al., 2007, 2008). Recently, Type Icn SNe have been
discovered, with narrow C and O lines, and no H and He lines (Gal-Yam et al.,
2021; Fraser et al., 2021).

These events are sometimes referred to as interacting SNe, since the
characteristic emission lines in their spectra come from the ionization of the
CSM, by emission from the SN shock (Smith, 2017; Fraser, 2020). Other
observational properties of these transients include a relatively blue
continuum, very diverse and prolonged light curves (Taddia et al., 2015;
Nyholm et al., 2020) and, for a fraction of interacting SNe, enhanced X-ray
(Fox et al., 2000; Katsuda et al., 2014; Chandra et al., 2012a, b) or radio
emission (van Dyk et al., 1993; Pooley et al., 2002; Williams et al., 2002;
Chandra et al., 2009).

The fact that we observe interacting SNe indicates that intense mass loss
activity can occur in the years preceding the collapse of the progenitor star,
and we still do not have a good physical understanding of the mechanisms
involved (see the review by Smith, 2014). Standard stellar evolution models do
not account for the amount of mass-loss needed to explain these transients,
limiting our understanding of their progenitors (Langer, 2012). In fact,
direct observations of the progenitor stars of Type IIn/Ibn SNe have been
elusive, with only a few being confirmed (Gal-Yam et al., 2007; Gal-Yam &
Leonard, 2009; Smith et al., 2010, 2011; Foley et al., 2011).

The progenitors of interacting SNe can be diverse, which is highlighted by the
extreme diversity in their observational properties (Taddia et al., 2015). The
progenitor of the Type IIn SN 2005gl was a massive luminous blue variable
(LBV)-like star (Gal-Yam & Leonard, 2009), so at least some of these events
are generated by LBV-like massive stars. Analysis of the environments of these
explosions, however, show that interacting SNe happen in very similar
locations to Type II SNe (Anderson et al., 2012; Kangas et al., 2017),
indicating that they might have overlapping mass ranges, and are not confined
to very massive stars.

Some massive stars can go through very unstable mass-loss phases, ejecting
large amounts of matter in energetic eruptions and explosions (Smith, 2014).
One well-known example is the massive Galactic LBV Eta Carinae, which went
through a major eruption event (the Great Eruption) in the 19th century
reaching an absolute magnitude of M${}_{V}\sim-14$ and with properties similar
to Type IIn SNe (Rest et al., 2012; Prieto et al., 2014; Smith et al., 2018a,
b). Estimates show that the star produced $\sim 10^{50}$ ergs of kinetic
energy and ejected $10-20$ M⊙ of material (Smith et al., 2003; Smith, 2006).
The main mechanism that caused the Great Eruption of Eta Carinae is still
under debate, although recent observations of its light echoes point to a
possible merger in a triple system (Smith et al., 2018a; Hirai et al., 2021).

Some of these giant stellar eruptions happen inside a very dense medium, and
produce spectroscopically and photometrically similar transients to
interacting SNe. Although being a non-terminal event, in the sense that the
progenitor star is still alive afterwards, these outbursts also interact with
the CSM. Because of these similarities, some of these outbursts are mistakenly
classified as SNe, originating the term ‘supernova impostors’. A famous
impostor associated with a massive star progenitor is SN 1997bs, which had a
peak at M${}_{V}\simeq$ M${}_{R}\simeq-14$, a very slow decline in its light
curve, and a spectrum dominated by strong narrow H$\alpha$ in emission (Van
Dyk et al., 2000), although there is presently no evidence that the star
survived (Adams & Kochanek, 2015). Recently, Pastorello et al. (2019a)
alternatively proposed SN 1997bs to be a stellar merger. Another impostor, SN
2000ch, is thought to have a progenitor similar to Eta Carinae and at its peak
of M${}_{V}\simeq-12.7$ it showed a full width at half maximum (FWHM) of the
Balmer lines up to $\sim 1550$ km s-1 (Filippenko, 2000; Wagner et al., 2004).
It is important to note that SN impostors can also arise from lower mass
stars, such as SN 2008S and similar events (Kochanek et al., 2012), whose
nature and that of other transients like it is still under debate (they might
have been real CCSNe, see e.g., Kochanek, 2011; Adams et al., 2016; Cai et
al., 2018, 2021). These transients have dust-enshrouded progenitors and ZAMS
masses of $\sim 6-10$ M⊙ (Prieto et al., 2008; Thompson et al., 2009).

A remarkable example of the confusion between real CCSNe and impostors is SN
2009ip. When discovered in 2009 (Maza et al., 2009) it was classified as a
Type IIn SN, reaching an absolute magnitude of M${}_{R}\sim-14.5$. One year
later, another bright eruption with similar luminosity was detected
(Pastorello et al., 2013), indicating that the progenitor was still alive and
did not experience a core-collapse event in 2009. These eruptions had, in
fact, many similarities with LBVs, showing velocities of a few hundred km s-1,
although faster material with thousands of km s-1 was also observed. Analysis
of pre-discovery images suggested that the progenitor star of SN 2009ip was a
very massive star with 50–80 M⊙, showing LBV-like variability in the years
before the 2009 eruption (Smith et al., 2010; Foley et al., 2011).

In 2012, SN 2009ip had two bright explosive events, reaching
M${}_{V}\sim-14.5$ in August (usually labeled Event A), and M${}_{R}\sim-18$
in October (called Event B). Several authors have discussed the nature of the
luminous event (Mauerhan et al., 2013; Pastorello et al., 2013; Fraser et al.,
2013). Mauerhan et al. (2013) claimed that the high velocities detected in the
outburst is evidence for a true CCSN during an LBV-like activity. Pastorello
et al. (2013), however, argued that such high velocities were also observed
during the 2011 event, which was clearly not a true SN. This scenario was also
supported by Fraser et al. (2013), who showed that there was no signal of
nebular emission from nucleosynthesis products after the 2012B eruption,
although nebular emission could be hidden by the continuing interaction with
CSM.

In the subsequent years, new analyses and a number of new observations have
contributed to the discussion of whether the 2012 eruptions of SN 2009ip were
a terminal event or not (Prieto et al., 2013; Levesque et al., 2014; Margutti
et al., 2014; Smith et al., 2014; Graham et al., 2014; Mauerhan et al., 2014;
Moriya, 2015; Fraser et al., 2013, 2015; Thoene et al., 2015; Smith et al.,
2016a; Graham et al., 2017). A number of newly detected transients also
presented very similar behavior to SN 2009ip. SN 2010mc, classified as a Type
IIn SN (Ofek et al., 2013), shows a nearly identical light curve to 2009ip,
with precursor variability, a decline after peak consistent with 56Co decay,
and late-time flattening in the light curve due to CSM interaction. Smith et
al. (2014) claimed that SN 2010mc was, in fact, a core-collapse event, arguing
that the missing nebular phase is explained by CSM obscuration, and that a
non-terminal event with energies of $\sim 10^{50}$ erg is implausible. SN
2016bdu (Pastorello et al., 2017) also had spectroscopic features of Type IIn
SNe, with several similarities to SN 2009ip, including the remarkably similar
light curve for its two explosive events. Pre-discovery images also showed
progenitor variability similar to that of SN 2009ip. SN 2018cnf (Pastorello et
al., 2019b) also had a sequence of two bright outbursts, after years of
explosive variability with eruptions reaching M${}_{r}\sim-14.7$. Pastorello
et al. (2019b) proposed a massive hypergiant or an LBV star as a plausible
progenitor for SN 2018cnf. More recently, the event AT2016jbu/Gaia16cfr showed
a variable behavior followed by two bright outbursts (Kilpatrick et al., 2018;
Brennan et al., 2021a, b). The availability of pre-discovery images from the
Hubble Space Telescope (HST) allowed Brennan et al. (2021a, b) to estimate
that the progenitor star had a relatively low mass of $\sim 20$ M⊙. Other
notable SN 2009ip-like transients might include SN 2015bh (Elias-Rosa et al.,
2016; Thöne et al., 2017; Boian & Groh, 2018), SNhunt151 (Elias-Rosa et al.,
2018), LSQ13zm (Tartaglia et al., 2016), and SN 2013gc (Reguitti et al.,
2019).

The growing class of SN 2009ip-like transients show several observational
similarities: i) erratic pre-discovery variability, sometimes including LBV-
like outburts; ii) two luminous eruptive events, with the second peak reaching
$\sim-18$ mag; and iii) the spectroscopic properties of Type IIn SNe,
including multiple velocity components observed in the Balmer lines.

Here, we present observations of SN 2011fh, which share many similarities to
SN 2009ip. Discovered by amateur astronomer Berto Monard on 2011 August 24
from South Africa, SN 2011fh was classified as a Type IIn SN at peak by Prieto
& Seth (2011). The source is located in the outskirts of the nearby galaxy NGC
4806. We describe our 10 years of observations of SN 2011fh in Section 2.
Section 3 describes the host galaxy (Section 3.1), the photometric (Section
3.2) and spectroscopic (Section 3.3) evolution, and the environment where SN
2011fh is located (Section 3.4). We discuss our analysis in Section 4, and
present a final summary of our conclusions in Section 5.

##  2 Observations

SN 2011fh was discovered by amateur astronomer Berto Monard on 2011 August 24
UT111UT dates are used throughout this paper. using unfiltered images from the
Klein Karoo Observatory in South Africa (Monard et al., 2011), at $\sim 14.5$
mag. The source was classified as a Type IIn SN at peak by Prieto & Seth
(2011), using observations obtained with the du Pont 2.5-m telescope and the
Wide Field reimaging CCD camera (WFCCD) at the Las Campanas Observatory (LCO)
on 2011 August 28. A series of spectroscopic observations were obtained after
peak brightness using the Apache Point Observatory (APO) 3.5 m telescope, the
du Pont 2.5 m telescope, and the Magellan Baade and Magellan Clay 6.5 m
telescopes. Photometric measurements were obtained by the Carnegie Supernova
Project-II (CSP2, Hamuy et al., 2006; Phillips et al., 2019; Stritzinger et
al., 2020) using the Swope 1-m and du Pont 2.5-m telescopes, and archival data
were retrieved from the Catalina Real-Time Transient Survey (CRTS, Drake et
al., 2009) and amateur astronomers. We also analyze archival images obtained
by the Spitzer Space Telescope (SST) and the Hubble Space Telescope (HST), as
well as spectroscopic observations made with the Multi-Unit Spectroscopic
Explorer (MUSE) instrument on the Very Large Telescope (VLT).

Figure 1: Unfiltered images of the galaxy NGC 4806 and SN 2011fh obtained by
Berto Monard (left and middle panel) and a r-band CSP2 image at a later phase
(right panel). SN 2011fh reached a peak absolute magnitude of M${}_{r}=-17.88$
on 2011 August 24, as it is shown in the middle panel. The $0.5\arcmin$
horizontal bar corresponds to 4.65 kpc at the distance of NGC 4806. The panels
also show the UT date when the images were taken, and the absolute $r$-band
magnitude of SN 2011fh at that time.

###  2.1 Optical and NIR Photometry

Archival unfiltered images obtained by amateur astronomers and CRTS span from
2007 to 2013 (Figure 1). Berto Monard obtained 27 unfiltered images between
2009 January 06 and 2013 January 22, using Meade 12- and 14-inch f/8
telescopes, and an SBIG ST-8 CCD with a pixel scale of 1.6 arcsec/pixel. Five
images were obtained by Stu Parker using a Celestron 14-inch f6.3 telescope
with a pixel scale of 1.1 arcsec/pixel, and a Celestron 11-inch f6.3 telescope
with a pixel scale of 1.6 arcsec/pixel, both using an SBIG ST-10 3 CCD camera,
between 2011 June 16 and August 09. Three images were also obtained by Greg
Bock using a Meade 14-inch telescope, and an SBIG ST-10 Dual CCD Camera, with
a pixel scale of 1.1 arcsec/pixel, between 2011 June 21 and 2012 February 17.
CRTS unfiltered images were all obtained by the Siding Springs Survey 0.5m
telescope with a pixel scale of 1.8 arcsec/pixel.

The unfiltered images were reduced using standard procedures, including bias
and dark subtraction, and flat-fielding. The photometric measurements were
made using point spread function fitting (PSF) through DOPHOT (Schechter et
al., 1993; Alonso-García et al., 2012). The measured fluxes were calibrated to
CSP2 $r$-band magnitudes using stars in the field. The SDSS $r$-band was
chosen because the emission is strongly dominated by H$\alpha$, as it is usual
in transients arising from the interaction with hydrogen rich CSM, and also
because this band is roughly at the midpoint of the wavelengths spanned by the
unfiltered response function. The photometric measurements are presented in
Table 8 and the $r$-band light curve of SN 2011fh is shown in Figure 2.

After peak, $BVri$ images of SN 2011fh were obtained by the Carnegie Supernova
Project-II (CSP2), with observations spanning from 2013 February 20 to 2016
January 14 using the direct imaging cameras of the du Pont 2.5 m telescope and
the Henrietta Swope 1.0 m telescope (Figure 1) at LCO. Near-infrared (NIR)
$YJH$ images (Krisciunas et al., 2017), were taken between 2013 February 23
and 2017 May 16 with RetroCam on the du Pont telescope.

The CSP2 images were reduced using standard techniques in IRAF (Freedman &
Carnegie Supernova Project, 2005; Hamuy et al., 2006; Stritzinger et al.,
2020). PSF photometry of SN 2011fh was extracted using DOPHOT, and the
calibration used CSP2 photometry of standard stars in the field. The optical
$BVri$ photometry is reported in Table 9 and the NIR $YJH$ photometry in Table
10. The CSP2 light curves of SN 2011fh are shown in Figure 2.

Figure 2: The complete photometry of SN 2011fh, including data from amateur
astronomers, CRTS, CSP2, and the Spitzer Space Telescope. Unfiltered images
from amateur astronomers and the CRTS were calibrated to the SDSS $r$-band
magnitudes and span from 2007 January 25 to 2013 June 29. SN 2011fh was
observed by the CSP2 from 2013 February 20 to 2017 May 16 using the $BVriYJH$
filters. The Spitzer observations include early detections on 2009 August 6
and September 5, and five additional observations from 2013 March 31 to 2015
May 16. Spectra were obtained at the epochs marked as vertical lines at the
top of the figure.

###  2.2 Spectroscopy

SN 2011fh was classified as a Type IIn SN $\sim 3$ days after peak brightness
using a spectrum obtained by the du Pont 2.5 m Telescope with the WFCCD
camera. The spectrum was obtained using a $1.7\arcsec$ slit with a FWHM
resolution of $\sim 8$Å. Three other spectra obtained with the same instrument
and similar configurations were taken on 2012 June 15, 2013 October 7, and on
2013 January 12. Four optical spectra were also obtained with the APO 3.5 m
Telescope on 2012 January 01, 2012 February 29, 2012 April 29, and on 2013
March 18, using the Dual Imaging Spectrograph (DIS), with a $1.5\arcsec$ slit
and a FWHM resolution of $\sim 7$Å. A high-resolution spectrum
($\Delta\lambda=0.2$ Å) was obtained with the Magellan Clay 6.5m Telescope
using the Magellan Inamori Kyocera Echelle (MIKE, Bernstein et al., 2003)
spectrograph with a $1\arcsec$ slit on 2011 August 31. An additional near-
infrared (NIR) spectrum was obtained with the Folded port InfraRed Echellette
(FIRE, Simcoe et al., 2013) spectrograph on the Magellan Baade 6.5m Telescope,
using a $0.6\arcsec$ slit, and the high-throughput prism mode, which gives
continuous coverage from 0.8-2.5 micron and resolutions of R $\sim 500$, 450,
and 300 in the $JHK$ bands, respectively.

The low resolution optical spectra were reduced using standard techniques in
IRAF twodspec and onedspec packages. We also applied a constant flux
calibration correction to the spectra to match the synthetic $r$-band
photometry to the light curve. The MIKE high-resolution spectrum was reduced
using the CarPy reduction pipeline. The NIR FIRE observations were obtained
using an ABBA nodding sequence and a bright A0V star close to the target in
the sky was used as a flux and telluric standard. The spectrum was reduced
using the IDL pipeline FIREHOSE (Simcoe et al., 2013).

The spectra are summarized in Table 1, including the dates, instruments and
configurations used for the observations. The spectral evolution of SN 2011fh
is shown in Figure 3, and the NIR FIRE spectrum is shown in Figure 4.

Table 1: Spectroscopic observations of SN 2011fh

Phase | UT Date | HJD | Telecscope/Instrument | Slit Width | Resolution | Wavelength Range
---|---|---|---|---|---|---
(days) |  |  |  | (arcsec) | (Angstrom) | (Angstrom)
+3 | 2011-08-30 | 2455803.479 | du Pont 2.5m/WFCCD | 1.7 | 8 | 3700 - 9120
+4 | 2011-08-31 | 2455805.479 | Magellan Clay/MIKE | 1.0 | 0.2 | 3500 - 9000
+129 | 2012-01-02 | 2455928.983 | APO 3.5m/DIS | 1.5 | 7 | 3470 - 9520
+187 | 2012-02-29 | 2455986.895 | APO 3.5m/DIS | 1.5 | 7 | 3570 - 9520
+246 | 2012-04-29 | 2456046.740 | APO 3.5m/DIS | 1.5 | 7 | 3470 - 9520
+293 | 2012-06-15 | 2456093.612 | du Pont 2.5m/WFCCD | 1.7 | 8 | 3470 - 9520
+505 | 2013-01-12 | 2456304.846 | du Pont 2.5m/WFCCD | 1.7 | 8 | 3670 - 9225
+570 | 2013-03-18 | 2456369.840 | APO 3.5m/DIS | 1.5 | 7 | 3480 - 9520
+600 | 2013-04-18 | 2456400.507 | Magellan Baade/FIRE | 0.6 | 24/36/73 | 8130 - 24740
+684 | 2013-07-10 | 2456484.465 | du Pont 2.5m/WFCCD | 1.7 | 8 | 3670 - 9120
+1359 | 2015-05-17 | 2457159.037 | VLT/MUSE | 0.8 | 2.3 | 4710 - 9270

Note. — Since the VLT/MUSE spectrum comes from integral field data,
$0.8\arcsec$ is the radius of a circular aperture around the source in the
datacube, not the slit width.

Figure 3: The spectral evolution of SN 2011fh from 3 (top) to 1359 (botton)
days after peak brightness. Wavelength is in the rest frame, and we show
log${}_{10}\textrm{F}_{\lambda}$ plus an offset. The phase is relative to
maximum brightness in $r$-band. Prominent lines of H, He I, Ca II, Na I, [S
II], and [O II] are highlighted with dashed lines. The spectra at phases +3 d,
+293 d, +505 d, and +684 d were obtained by the du Pont 2.5 m + WFCCD, and the
spectra at phases +129 d, +187 d, +246 d, and +570 d by the APO 3.5 m + DIS
instrument (Table 1). The last spectrum, at phase +1359 d, was extracted from
VLT + MUSE observations. All the spectra presented in this figure are
available in WISeREP (https://www.wiserep.org/, Yaron & Gal-Yam, 2012).

Figure 4: Rest frame FIRE NIR spectrum of SN 2011fh obtained on 2013 April 18,
600 days after peak brightness. The spectrum is dominated by emission lines of
hydrogen, with prominent Paschen series (Pa$\alpha$, Pa $\beta$, Pa $\gamma$,
Pa $\delta$, Pa $8-3$, and Pa $9-3$) and Brackett series (Br $\gamma$ and Br
$\delta$) lines highlighted with by the blue dashed lines. Lines of He I at
1.083 and 2.058 $\mu$m are marked by purple lines. The Ca II $\lambda\lambda
8498,8542,8662$ triplet is marked by orange dashed lines. The gray band marks
the region of the strong telluric absorption.

###  2.3 Spitzer Imaging

After SN 2011fh was discovered, Spitzer/IRAC (Fazio et al., 2004) 3.6- and
4.5-micron archival images from 2009 August and September (Program 6007; PI:
K. Sheth) revealed a mid-infrared source at the position of the SN (Monard et
al., 2011). Five more images using the 3.6 and 4.5-micron channels were taken
from 2013 to 2015 (Programs 90137, 10139, 11053, PI: O. Fox). We downloaded
the calibrated PBCD images from the Spitzer Heritage
Archive222https://sha.ipac.caltech.edu/applications/Spitzer/SHA/ and performed
PSF photometry using DOPHOT. The result of our photometry is reported in Table
11. The 3.6- and 4.6-micron evolution of SN 2011fh is also shown in Figure 2.
The Spitzer/IRAC photometry of SN 2011fh 586 days after peak was also
presented by Szalai et al. (2019), and is fairly consistent with our results.

###  2.4 HST Imaging

SN 2011fh was observed by HST on 2016 October 20 (proposal ID 14149; PI: A.V.
Filippenko) using the Wide Field Camera 3 (WFC3) Ultraviolet-Visible (UVIS)
channel, and the F336W and F814W filters with exposure times of 710 s and 780
s, respectively. The data were retrieved from the Mikulski Archive for Space
Telescopes (MAST). PSF photometry on the calibrated flt images was performed
with DOLPHOT (Dolphin, 2016), using the drizzled F814W image as a reference
for point source detection and astrometry. We selected point sources using the
DOLPHOT parameter criteria of Williams et al. (2014): $sharpness^{2}\leq
0.15$, $crowding\leq 1.0$, and $signal-to-noise\geq 5$. The F814W image is
shown in Figure 5 with an inset zooming on the region around SN 2011fh. We
measured the WFC3 Vega magnitudes of m(F336W) $=18.956\pm 0.006$ mag and
m(F814W) $=19.374\pm 0.004$ mag for the SN.

Figure 5: HST F814W image of NGC 4806 from 2016 October 20. The inset shows
the region around SN 2011fh. The $2\arcsec$ horizontal bar at the inset
corresponds to 0.31 kpc at the position of SN 2011fh. SN 2011fh has absolute
magnitudes of M${}_{F814W}=-13.32\pm 0.05$ mag and M${}_{F336W}=-14.25\pm
0.09$ mag.

###  2.5 MUSE Spectroscopy

Integral field spectroscopy of the galaxy NGC 4806 was obtained with MUSE
(Bacon et al., 2014) on the VLT on 2015 May 17, at a phase of $+1359$ days
(program ID:095.D-0172(A); Kuncarayakti et al. 2018). We retrieved the
reduced, wavelength and flux calibrated datacube from the ESO archive333
http://archive.eso.org. The cube has $0.2\arcsec\times 0.2\arcsec$ spaxels, a
seeing of $0.8\arcsec$, and a spectral range of $4750-9350$ Å. We used a
$0.8\arcsec$ radius aperture to extract the spectrum of SN 2011fh from the
MUSE datacube. The spectrum is shown in Figure 3 and in Figure 6, where we
highlight the emission lines observed at very late times.

Figure 6: The restframe VLT/MUSE spectrum of SN 2011fh obtained on 2015 May
17, 1359 days after peak brightness. Lines of H, He I, Ca II, Na I, [O I], [O
III], [S II], and Fe II are highlighted. Forbidden lines of [O III] and [S II]
are most likely due to host galaxy contamination.

##  3 Results

###  3.1 Host Galaxy

SN 2011fh is located at $\alpha=12^{h}56^{m}14.009^{s}$ and
$\delta=-29^{\circ}29^{\prime}54.82"$ (J2000.0), in one of the spiral arms of
the nearby galaxy NGC 4806. The redshift for NGC 4806, obtained from the NASA-
IPAC Extragalactic Database (NED)444https://ned.ipac.caltech.edu, is
$z=0.00803$ (Paturel et al., 2003). Using the heliocentric velocity for the
galaxy of $v=2408\pm 4$ km/s (Paturel et al., 2003), and the Cosmicflows-3
distance-velocity calculator 555 http://edd.ifa.hawaii.edu/NAMcalculator/,
which has a flow model consistent with $\textrm{H}_{0}=75\ \textrm{km}\
\textrm{s}^{-1}\ \textrm{Mpc}^{-1}$ (Kourkchi et al., 2020), we get a distance
of $D=30.66$ Mpc. Since the model has a zero-point error of $3\%$ (Kourkchi et
al., 2020), we obtained an uncertainty in the distance of $0.92\
\textrm{Mpc}$. The distance modulus is therefore $\mu=32.43\pm 0.06$ mag.

We estimate the redshift at the position of SN 2011fh using the MUSE spectrum
of the HII region around the SN. We used Gaussian fitting to estimate the
central position of the emission lines of H$\alpha$, H$\beta$, [N II]
$\lambda\lambda 6548,6584$ Å, [S II] $\lambda\lambda 6716,6731$ Å, and [O III]
$\lambda\lambda 5007,4959$ Å, and measured the local redshift of $z=0.00812$.
This is the value we consider throughout the paper.

The Milky Way (MW) extinction towards NGC 4806 is $A_{V}=0.232$ mag (Schlafly
& Finkbeiner, 2011) and $E(B-V)=0.0748$ mag, with $R_{V}=3.1$. We estimate the
host galaxy extinction using the high resolution MIKE spectrum. The
interstellar Na I doublet (Na I D) absorption lines from the galaxy have
equivalent widths (EWs) of $0.1231$ Å, for Na D1 and $0.1966$ Å, for Na D2. By
scaling these values with the EWs of the MW Na I D absorption features (0.1118
Å and 0.1902Å, for Na D1 and D2, respectively), we obtained $A_{V}=0.24$ mag
or $E(B-V)=0.078$ mag for the host galaxy. Thus, using the $2\%$ photometric
error in the measurements of Schlafly & Finkbeiner (2011), and assuming a
$10\%$ uncertainty in the host extinction, the total estimated V-band
extinction for SN 2011fh is $A_{V}=0.48\pm 0.03\ \textrm{mag}$.

We also estimated the extinction of the cluster of stars near SN 2011fh. We
use a spectrum extracted with a $2-3\arcsec$ annulus aperture from the MUSE
datacube and measured a Balmer decrement of H$\alpha$/H$\beta=3.34$. Using
Levesque et al. (2010), this value implies a total $E(B-V)=0.16\ \textrm{mag}$
or $A_{V}=0.49\ \textrm{mag}$. Since this result is consistent with the total
estimated extinction for SN 2011fh, we decided to use the same value of
$A_{V}=0.48\ \textrm{mag}$ for the stars surrounding the SN.

###  3.2 Photometric Evolution

Figure 2 shows the complete photometric evolution of SN 2011fh from 2007 to
2017. The SN reached peak $r$-band brightness on 2011 August 24 (JD =
2455799), at $r=14.94\pm 0.02$ mag, or M${}_{r}=-17.88\pm 0.04$ mag666For the
uncertainty in the absolute magnitudes, we took into account the estimated
error in the apparent magnitude, the $3\%$ uncertainty in the distance given
by Kourkchi et al. (2020), and the $2\%$ error in color estimated by Schlafly
& Finkbeiner (2011).. Before peak, a five month-long (173 days) brightening
event is observed, where the progenitor goes from $r=17.63\pm 0.15$ on 2011
February 11, to $r=16.26\pm 0.04$ on 2011 July 28. After this slow
brightening, the source shows a rapid increase to $r\sim 14.94$ mag over $\sim
40$ days, culminating in the August peak.

A precursor event followed by a rapid rise to peak brightness is commonly seen
in 2009ip-like transients. For this class of transients, the initial rise is
commonly called “Event A”, and the main peak is called “Event B”. In Figure 7,
we show the $r$-band light curve near peak brightness of SN 2011fh, SN 2009ip
(Mauerhan et al., 2013; Fraser et al., 2013), SN 2016bdu (Pastorello et al.,
2017), and SN 2018cnf (Pastorello et al., 2019b). The Vega $R$-band magnitudes
of SN 2009ip were scaled to AB magnitudes by adding $0.21$ mag (Blanton &
Roweis, 2007). The four transients show a remarkably similar absolute
magnitude at peak, with M${}_{r}=-18.29\pm 0.10$ for SN 2009ip,
M${}_{r}=-18.00\pm 0.16$ for SN 2016bdu, and M${}_{r}=-18.25\pm 0.05$ for SN
2018cnf. SN 2009ip and SN 2016bdu have a very similar Event A, peaking at
M${}_{r}\sim-15.5$, and lasting for $\sim 40$ days. Although different in
brightness and duration, the slow increase in luminosity seen for SN 2011fh
before peak might share a similar origin to the other transients. After the
precursor Event A, all four transients show a rapid brightening to peak, over
$\sim 15$ days for SN 2009ip, $\sim 25$ days for SN 2016bdu, and $\sim 30$
days for SN 2018cnf. It is worth noting that the magnitude difference between
the end of Event A and the peak brightness is considerably larger for the
other transients.

Before the explosive events of 2011, the source of SN 2011fh is seen to be
bright, with an apparently constant luminosity over the preceding three years
(see Figure 2). The images are relatively shallow, so it is hard to tell
whether the variability is real or not. The absolute magnitudes of
M${}_{r}\sim-14.5$ to M${}_{r}\sim-15.5$ are very high in the years preceding
the major brightening event. For comparison, the progenitors of SN 2009ip and
SN 2016bdu had $r$-band absolute magnitudes of $-13$ to $-14$ mag.

The stellar cluster around SN 2011fh might be a source of photometric
contamination in the pre-outburst images, which were all obtained using
cameras with large pixel scales. We estimate the photometric contamination in
Appendix A. Our analysis shows that the contribution of the stellar cluster is
not a serious problem, although it increases for the bluer bands.

Spitzer imaging from 2009 also reveals a bright source at the position of SN
2011fh. Photometry of the 3.6- and 4.5-micron channels show a source with
apparent magnitudes of 15.9 and 15.5, respectively. These imply luminosities
of $3.5\times 10^{6}$ and $2.5\times 10^{6}$ L⊙, respectively. In the
discovery report, Monard et al. (2011) noted that the mid-IR luminosity may
have a contribution from the HII region around the source, due to the 260 pc
spatial resolution of the IRAC instrument.

During the first days after peak, the brightness of SN 2011fh drops
significantly. Unfortunately, between 7 and 110 days after peak in $r$-band,
there were no photometric observations of the source. After 110 days, it is
significantly fainter and slowly fading. The decay, however, is not monotonic,
with two pronounced peaks at 228 (2012 April 11, or JD = 2456029) and 350 days
(2012 August 12, or JD = 2456151) after peak (Figure 7). Bumps and
fluctuations after peak brightness are common in the light curve of SN 2009ip-
like objects, but the amplitudes seen for the other three sources shown in
Figure 7 are significantly smaller. This behavior is usually interpreted as
the collision of the ejecta with circumstellar material showing a complex
density profile. The overall evolution of SN 2011fh in the months following
peak brightness is, however, very similar to SN 2009ip, with a slow fading
probably caused by the ongoing interaction with a dense CSM.

The CSP2 optical $BVri$\- and NIR $YJH$-band images of SN 2011fh, as well as
new IR Spitzer observations, were obtained only starting $\sim 550$ days after
peak (2013 February 20). The behavior in the other filters is very similar to
$r$-band. Initially, all bands show a slow drop in brightness, with an
apparent flattening by the beginning of 2015, 1300 days after peaking in
$r$-band. At this epoch the source is still quite luminous, with
M${}_{r}\sim-13.5$ mag. Surprisingly, an apparent rise in luminosity is
detected by the beginning of 2016 in the optical $BVi$\- and NIR $YJH$-bands,
although not in the $r$-band. The apparent brightening in the NIR bands is
clear in the last point of the light curve where the $Y$-band brightened by
$\Delta Y\approx 0.36$ mag over 418 days.

HST observations from October 2016 show a very bright source at the position
of SN 2011fh (Figure 5), with absolute magnitudes of M${}_{F814W}=-13.32\pm
0.05$ mag and M${}_{F336W}=-14.25\pm 0.09$ mag. This is much more luminous
than even the most luminous quiescent stars known (e.g., R136a1 and BAT99-98
in the Tarantula Nebula of the Large Magellanic Cloud, with absolute
bolometric magnitudes of $-12.2$ and $-12.0$, respectively; Bestenlehner et
al., 2020; Hainich et al., 2014). Another parallel can be made with Eta
Carinae, which reached an absolute magnitude M${}_{V}\sim-14$ mag during its
19th Century Great Eruption (Smith et al., 2018a). This might be simply due to
late-time emission due to the ongoing interaction of the ejecta with the CSM,
as it is commonly seen for interacting transients (e.g., Graham et al., 2017),
or alternatively indicate that we are seeing the star in a LBV-like eruptive
state.

Figure 7: Comparison of the $r$-band light curve near peak brightness of SN
2011fh to SN 2009ip (Mauerhan et al., 2013; Fraser et al., 2013), SN 2016bdu
(Pastorello et al., 2017) and SN 2018cnf (Pastorello et al., 2019b). The phase
is in days relative to peak brightness.

####  3.2.1 Colors

Figure 8 shows the CSP2 $B-V$, $r-i$, $J-H$, and the $[3.6]-[4.5]$ color
evolution of SN 2011fh. For the optical and near-IR bands, unfortunately, we
only have colors at very late-times, starting at 550 days after the maximum
brightness. We are able to estimate the mid-IR colors of the progenitor 750
days before peak brightness.

As commonly seen for Type IIn SNe, the $B-V$ and $r-i$ evolution of SN 2011fh
is very slow (Taddia et al., 2013). The $B-V$ color becomes slightly bluer,
consistent with the temperature increase measured in the spectra and the
overall spectral energy distribution (SED, see Section 3.2.2 and Section 3.3).
The $J-H$ color becomes slightly redder, increasing from 0.22 mag, at phase
$+550$ days, to 0.34 mag, at phase $+1676$ days, although this last value has
large uncertainties.

Type IIn SNe are usually bright in the mid-IR bands, and have very
heterogeneous properties (Szalai et al., 2021). A significant color evolution
is seen for SN 2011fh in the mid-IR, if we compare the values before and after
the luminous outburst of 2011. The source has $[3.6]-[4.5]=0.37$ mag at phase
$-715$ days, and $[3.6]-[4.5]=0.87$ mag, 580 days after peak. After peak, the
mid-IR color becomes a little redder, decreasing slowly with time.

We also show the $B-V$ and $R-I$ evolution of SN 2009ip in Figure 8 from
Fraser et al. (2015). We show the $R-I$ colors, because $r-i$ is not available
for these epochs. The Vega R- and I-band values were scaled to AB magnitudes
by adding, respectively, $0.21$ mag and $0.45$ mag (Blanton & Roweis, 2007).
In $B-V$, SN 2011fh is significantly bluer than SN 2009ip, although both
events become redder with time. This might be due to the fact that SN 2011fh
is located in a blue star-forming region, with the contamination from the
cluster increasing as the SN fades (see Appendix A). The other possible
explanation is that SN 2011fh is intrinsically hotter than SN 2009ip. The
$R-I$ colors of SN 2009ip are redder than the $r-i$ values for SN 2011fh. This
might imply that SN 2009ip has stronger H$\alpha$ emission lines in relation
to the spectral continua throughout its evolution. SN 2009ip has a very
similar mid-IR color evolution to SN 2011fh. Although the mid-IR estimates are
uncertain at late-times, SN 2009ip becomes a little redder after peak, and
then bluer around a phase of $\sim 1000$ days (Szalai et al., 2019, 2021).

Figure 8: $B-V$, $r-i$, $J-H$ and $[3.6]-[4.5]$ color evolution of SN 2011fh.
Phase is in days relative to $r$-band peak. For comparison, we also show the
$B-V$, $R-I$, and $[3.6]-[4.5]$ color evolution of SN 2009ip (Fraser et al.,
2015; Szalai et al., 2019). We show the $R-I$ color, because $r-i$ data is not
available.

####  3.2.2 Spectral Energy Distribution

We use the optical, near-infrared (NIR), and infrared (IR) photometry to
constrain the spectral energy distribution (SED) of SN 2011fh after peak. We
used SUPERBOL (Version 1.7; Nicholl, 2018) to estimate the blackbody
temperature, luminosity, and photosphere radius between phases $+544$ and
$+1602$ days. Since SN 2011fh has very strong H$\alpha$ emission (see Section
3.3), we carried out this first analysis using only the optical $BVi$ bands.
The results of the single-blackbody fits to these three filters are shown in
Figure 9 and reported in Table 3.

SN 2011fh has an unusual temperature evolution for either a SN or a SN
impostor. Transients driven by CSM interaction usually show a smooth decrease
in temperature after peak (Margutti et al., 2014; Kilpatrick et al., 2018),
opposite to the behavior of SN 2011fh. The blackbody temperature at phase
$+544$ days is T${}_{BB}\approx 8920$ K, and it rises to T${}_{BB}\approx
9950$ K at phase $+630$ days. A higher blackbody temperature of $\sim 11000$ K
is measured at 1300 days, but the late-time measurements are uncertain, and
the fluxes may be dominated by line emission at these epochs (see Figure 3 and
Section 3.3). The radius evolution of SN 2011fh also differs from the usual
increase and then smooth decrease after peak. The evolution of the
photospheric radius of SN 2009ip was also unusual and was interpreted by
Margutti et al. (2014) as being due to a more complex CSM structure.

We also model the SED using DUSTY (Ivezic & Elitzur, 1997), a code for solving
radiative transfer through a spherically symmetric dusty medium. We assume a
central stellar source using atmospheric models from Castelli & Kurucz (2003),
a dust shell with a thickness $R_{out}/R_{in}=2$, and silicate dust from
Draine & Lee (1984) with a standard MRN dust size distribution (Mathis et al.,
1977). We find best-fitting models using a MCMC wrapper around DUSTY as in
Adams & Kochanek (2015). In Table 2 we show the resulting photospheric
($T_{phot}$) and dust ($T_{dust}$) temperatures, together with the estimated
photospheric radius ($R_{phot}$) and inner radius of the dust shell
($R_{in,dust}$), the total luminosity ($L$), and the dust optical depth in V
band ($\tau_{V}$). Figure 10 shows the optical, NIR and IR SED, and the DUSTY
fits at phase $+570$ days. At this period, the SED was fit with
$\tau_{V}=0.68$, $T_{phot}\approx 7865$ K, $T_{dust}\approx 950$ K, and
$L\approx 10^{7.74}\ \textrm{L}_{\odot}$. The dust emission increases in
temperature with time, reaching $\sim 1116$ K at phase $+1143$ days.

Figure 9: Temperature (TBB), bolometric luminosity (Lbol), and radius (RBB)
evolution of SN 2011fh, estimated from blackbody fits to the spectral continua
using Astropy models (red squares), and to the spectral energy distributions
(SEDs) using SUPERBOL (green circles). The phase is relative to the maximum
$r$-band luminosity in days.

Figure 10: DUSTY fit to the SED at phase $+570$, using the optical and Spitzer data. The photospheric temperature is $\sim 7865$ K and the dust temperature is $\sim 950$ K. The fit parameters are given in Table 2.  Table 2: DUSTY fits of the SEDs  Phase | $\tau_{V}$ | $T_{phot}$ | $T_{dust}$ | log $L$ | log $R_{phot}$ | log $R_{in,dust}$
---|---|---|---|---|---|---
(days) |  | (K) | ($1\sigma$ error) | (K) | ($1\sigma$ error) | (L⊙) | ($1\sigma$ error) | (R⊙) | ($1\sigma$ error) | (cm) | ($1\sigma$ error)
+570 | 0.68 | 7865 | 511 | 950 | 320 | 7.74 | 0.04 | 3.60 | 0.04 | 16.52 | 0.29
+1113 | 0.91 | 8321 | 537 | 903 | 233 | 7.48 | 0.06 | 3.42 | 0.03 | 16.50 | 0.26
+1143 | 0.61 | 8823 | 779 | 1116 | 240 | 7.48 | 0.05 | 3.37 | 0.05 | 16.38 | 0.18
+1300 | 0.78 | 8915 | 635 | 1006 | 182 | 7.40 | 0.05 | 3.32 | 0.04 | 16.44 | 0.17
Table 3: SUPERBOL fits to the $BVi$-bands  Phase | TBB |  $10^{41}$ LBB |  $10^{14}$ RBB
---|---|---|---
(days) | (K) | (error) | (erg s-1 ) | (error) | (cm) | (error)
+544 | 8910 | 55 | 2.17 | 0.11 | 2.20 | 0.15
+555 | 8950 | 105 | 2.12 | 0.15 | 2.15 | 0.15
+566 | 9090 | 74 | 2.00 | 0.10 | 2.03 | 0.14
+587 | 9380 | 58 | 2.07 | 0.27 | 1.93 | 0.16
+602 | 9010 | 76 | 1.95 | 0.24 | 2.04 | 0.14
+616 | 9500 | 38 | 1.62 | 0.17 | 1.67 | 0.16
+625 | 9940 | 27 | 1.78 | 0.08 | 1.60 | 0.11
+1298 | 10900 | 465 | 1.03 | 0.11 | 1.01 | 0.08
+1300 | 9730 | 246 | 0.86 | 0.05 | 1.16 | 0.08
+1602 | 9820 | 147 | 1.10 | 0.05 | 1.29 | 0.07

###  3.3 Spectroscopic Evolution

In Figure 3, we show the complete spectral evolution of SN 2011fh from 3 to
1359 days after peak brightness. We use vertical lines to highlight prominent
emission features of H, He I, Ca II, Na I, [S II], and [O II]. The spectra are
dominated by Balmer lines in emission during the whole evolution, with
features of He I and Ca II becoming more prominent with time. The forbidden
lines of [S II] and [O II] seen throughout the evolution of SN 2011fh are most
likely due to host galaxy contamination. In Figure 11, we show the evolution
of the emission profiles of H$\beta$, He I $\lambda 5876$, the [Ca II]
$\lambda\lambda 7291,7324$ doublet, and the Ca II $\lambda\lambda
8498,8542,8662$ triplet.

Due to the different emission mechanisms, interacting transients can show very
complex Balmer line profiles (see the review by Smith, 2017). For the spectra
of SN 2011fh, we used Astropy (Astropy Collaboration et al., 2013, 2018) to
fit the H$\alpha$ emission profiles with narrow Gaussian components and
intermediate Lorentzian components, plus a constant continuum, as shown in
Figure 12. The estimated full width at half maximum (FWHM) velocity and the
luminosity of the two components throughout the evolution of SN 2011fh are
shown in Figure 13, and reported in Table 4 and Table 5. Errors are estimated
by the $1\sigma$ parameter uncertainty for the best-fitting Lorentzian and
Gaussian models. A $3\%$ uncertainty is also considered for the error in the
luminosity of the components. Due to its proximity to the spectral resolution,
the FWHM of the narrow components are corrected by the resolution value. At
some epochs, the width of the narrow component falls below that value, and are
indicated as less than the spectral resolution in Table 4 and Figure 13.
Prominent blue absorption features are also seen at certain epochs, especially
at phases +129 and +187 days. The photospheric temperature, radius and
bolometric luminosity evolution of SN 2011fh, obtained from the spectral
continuum fitting, also shown in Figure 9 and reported in Table 6, agree well
with the SED models. Errors in temperature are given by the $1\sigma$
parameter uncertainty for the best-fitting model, while for the error in
luminosity and radius, a $3\%$ uncertainty in the distance is also considered.

The first spectrum in our sequence, obtained 3 days after peak, is dominated
by Balmer lines in emission, and resembles a typical Type IIn SN. It has a
blue continuum with a blackbody temperature of $T_{BB}\sim 8000\ \textrm{K}$.
The H$\alpha$ profile can be fit with an intermediate-width Lorentzian profile
with a FWHM velocity of $v_{FWHM}\sim 2263\ \textrm{km s}^{-1}$, and a narrow
Gaussian component with $v_{FWHM}\sim 204\ \textrm{km s}^{-1}$. The narrow
Gaussian emission is generated by the ionized dense CSM, while the Lorentzian
profile is indicative of broad electron scattering in the unshocked CSM
(Chugai, 2001).

The H$\alpha$ profile becomes significantly broader in the second spectrum at
$+129$ days. The Lorentzian profile now has $v_{FWHM}\sim 3376\ \textrm{km
s}^{-1}$, and the narrow Gaussian has $v_{FWHM}\sim 197\ \textrm{km s}^{-1}$.
The line is well described by a P-Cygni profile with a strong blueshifted
absorption feature. The blueshifted absorption wing extends to $\sim 6430\
\textrm{km s}^{-1}$, and is probably formed by fast-moving material along the
line of sight. At this epoch, the He I features and the Ca II triplet start to
develop, with H$\beta$ blended into the He I $\lambda 4922$ emission line. The
continuum becomes relatively redder, with a blackbody temperature of
$T_{BB}\sim 7400\ \textrm{K}$.

The FWHM of the H$\alpha$ Lorentzian begins to decrease at $+187$ days, with
$v_{FWHM}\sim 3087\ \textrm{km s}^{-1}$, and the narrow Gaussian component is
below the spectral resolution of $\sim 320\ \textrm{km s}^{-1}$. The
blueshifted absorption feature is still present, with a velocity of $\sim
6447\ \textrm{km s}^{-1}$. The blackbody temperature remains almost constant,
with $T_{BB}\sim 7500\ \textrm{K}$.

By $293$ days, the P-Cygni absorption feature is no longer easy to identify,
although the H$\alpha$ line cannot be entirely described by only emission
features, as can be seen in Figure 12. We measure $v_{FWHM}\sim 2568\
\textrm{km s}^{-1}$ for the Lorentzian component. This epoch is marked by a
sudden inversion in temperature to $T_{BB}\sim 8100\ \textrm{K}$. At this
epoch, the emission features of Ca II become more prominent, with the [Ca II]
$\lambda\lambda 7291,7324$ doublet beginning to develop and the three
components of the Ca II $\lambda\lambda 8498,8542,8662$ triplet becoming
easily distinguishable.

The emission features remain very similar during the three next epochs, with
He I and Ca II now very prominent. At phase $+505$ days, the Ca II $\lambda
8662$ line can be fit with a narrow Lorentzian profile with $v_{FWHM}\sim 583\
\textrm{km s}^{-1}$. The Fe II emission line forest is also detected, and
becomes stronger in the last spectrum of our sequence (see Figure 6). At
$+684$ days, the H$\alpha$ profile gets considerably narrower, with
$v_{FWHM}\sim 925\ \textrm{km s}^{-1}$ for the Lorentzian component. At this
time, the blackbody temperature reaches its maximum measured value of
$T_{BB}\sim 10200\ \textrm{K}$.

The unusual temperature increase seen in SN 2011fh, to both SEDs and the
spectral continua, might be generated by the shock interacting with the CSM, a
scenario that might be supported by the bumps seen in the $r$-band light
curve. Another possibility is that the CSM is becoming relatively thin,
revealing a hot stellar photosphere.

The NIR spectrum taken on 2013 April 18, 600 days after peak (Figure 4) shows
prominent emission lines of the Paschen and Brackett series, and features from
He I and Ca II. This was also observed in early NIR spectroscopy of the Event
B in SN 2009ip (Pastorello et al., 2013; Smith et al., 2013; Margutti et al.,
2014). Although our spectrum was taken in a much later time after peak, it is
still strongly dominated by emission features of hydrogen, such as it was
observed in 2012 for SN 2009ip.

The very late-time optical spectrum of SN 2011fh, obtained on 2015 October 7
at phase $+1359$ days, is shown in Figure 6. We use The Atomic Line List v2.04
777 https://www.pa.uky.edu/~peter/atomic/ to identify the emission lines in
the spectrum, and highlight the main features with traced lines. The Fe II
emission features (gray traced lines) become very prominent especially near
5200 Å. At this point, the H$\alpha$ profile is very narrow, being described
by a Lorentzian component with $v_{FWHM}\sim 726\ \textrm{km s}^{-1}$ and a
Gaussian component with $v_{FWHM}<110\ \textrm{km s}^{-1}$. The Ca II emission
lines also get considerably narrower, with $v_{FWHM}\sim 240\ \textrm{km
s}^{-1}$ for Ca II $\lambda 8662$ and the two components of the [Ca II]
$\lambda\lambda 7291,7324$ doublet. The detected [O II], [O III], and [S II]
lines are due most likely to a background H II region, and become stronger
with time as the contribution from SN 2011fh becomes fainter, while the [O I]
lines are probably due to the transient. As shown in Figure 9, we measure a
sudden drop in temperature at this epoch, which is not seen for the SUPERBOL
measurements in photometry. This difference might be due to the last spectrum
being strongly dominated by emission lines, which makes the blackbody fit in
the continuum very uncertain. We also fit the H$\beta$ emission lines with
intermediate- and narrow-witdh Lorentzian profiles, and use the resultant
luminosities to calculate the Balmer decrement evolution. We estimate a
relatively low value of $\textrm{H}\alpha/\textrm{H}\beta=3.01$ near peak
brightness, which might be due to Case B recombination or a high density CSM,
as it was observed for 2009ip (Levesque et al., 2014). For the phase $+246$
days, however, we measure $\textrm{H}\alpha/\textrm{H}\beta\simeq 13$, which
is very similar to the value found for the 2009ip-like SN 2016bdu (Pastorello
et al., 2017) and other Type IIn SNe (e.g., 1996al; Benetti et al., 2016).
Such high ratios might be related to newly formed dust or to Case C
recombination where the CSM is optically thick to the Balmer emission lines.

Table 4: FWHM velocities of the two H$\alpha$ components  Phase | vFWHM,Lorentzian | vFWHM,Gaussian
---|---|---
(days) | (km s-1) | ($1\sigma$ error) | (km s-1) | ($1\sigma$ error)
+3 | 2263.33 | 299.11 | 204.03 | 36.12
+129 | 3376.06 | 99.98 | 197.69 | 45.53
+187 | 3087.12 | 60.19 |  $<$320.00 | 26.19
+246 | 3006.35 | 42.37 |  $<$320.00 | 22.12
+293 | 2568.02 | 40.41 |  $<$365.00 | 13.23
+505 | 1464.20 | 23.71 |  $<$365.00 | 14.67
+570 | 1302.74 | 59.22 |  $<$320.00 | 83.59
+684 | 925.26 | 56.35 |  $<$365.00 | 68.59
+1359 | 726.56 | 14.63 |  $<$110.00 | 12.78
Table 5: Luminosities of the two H$\alpha$ components  Phase |  $10^{39}$ LLorentzian |  $10^{39}$ LGaussian
---|---|---
(days) | (erg s-1) | ($1\sigma$ error) | (erg s-1) | ($1\sigma$ error)
+3 | 61.92 | 11.94 | 8.46 | 1.69
+129 | 47.75 | 3.40 | 0.91 | 0.23
+187 | 24.20 | 1.57 | 0.77 | 0.10
+246 | 30.21 | 1.89 | 0.78 | 0.09
+293 | 18.90 | 1.20 | 1.10 | 0.09
+505 | 11.76 | 0.75 | 0.83 | 0.07
+570 | 12.33 | 1.08 | 0.40 | 0.16
+684 | 8.36 | 0.92 | 0.52 | 0.19
+1359 | 4.41 | 0.29 | 0.13 | 0.02
Table 6: Blackbody fits to the spectral continua  Phase | TBB |  $10^{41}$ LBB |  $10^{14}$ RBB
---|---|---|---
(days) | (K) | ($1\sigma$ error) | (erg s-1 ) | ($1\sigma$ error) | (cm) | ($1\sigma$ error)
+3 | 8052 | 38 | 34.08 | 2.17 | 10.67 | 0.36
+129 | 7432 | 50 | 5.67 | 0.38 | 5.11 | 0.19
+187 | 7528 | 47 | 2.78 | 0.18 | 3.49 | 0.12
+246 | 7543 | 41 | 4.19 | 0.27 | 4.26 | 0.15
+293 | 8140 | 87 | 2.84 | 0.22 | 3.01 | 0.13
+505 | 8815 | 90 | 2.24 | 0.17 | 2.28 | 0.10
+570 | 9012 | 97 | 1.83 | 0.13 | 1.97 | 0.08
+684 | 10170 | 146 | 1.99 | 0.16 | 1.62 | 0.08
+1359 | 6942 | 58 | 0.65 | 0.04 | 1.99 | 0.08

Figure 11: Evolution of H$\beta$, He I $\lambda 5876$, the [Ca II]
$\lambda\lambda 7291,7324$ doublet, and the Ca II $\lambda\lambda
8498,8542,8662$ triplet. The vertical dashed lines mark the rest wavelengths.

Figure 12: Evolution of the H$\alpha$ emission line and a model built with a
narrow-width Gaussian component and an intermediate-width Lorentzian
component. Velocity corresponds to rest-frame wavelength of H$\alpha$, and the
spectral phases are given above each panel. Absorption features are seen in
some epochs, especially at $+129$ and $+187$ days.

Figure 13: Evolution of the H$\alpha$ FWHM velocity (top panel) and luminosity
(bottom panel). The purple squares are the narrow Gaussian component, while
the blue circles are the Lorentzian. The arrows indicate an upper limit for
the narrow component, given by the spectral resolution. Phase is in days
relative to $r$-band maximum.

###  3.4 Environment Analysis

####  3.4.1 MUSE Spectroscopy

We use the MUSE data to investigate the HII regions and the environment around
SN 2011fh. The datacube was corrected for reddening and redshift using the
values from Section 3.1. To avoid contamination from the light of the SN, we
masked a $6\arcsec$ circular region around SN 2011fh. We used ifuanalysis
888https://ifuanal.readthedocs.io/en/latest/index.html to select and estimate
the properties of the HII regions. ifuanalysis uses STARLIGHT (Cid Fernandes
et al., 2005; Mateus et al., 2006; Asari et al., 2007) to fit the stellar
continuum with stellar population synthesis models, and Astropy (Astropy
Collaboration et al., 2013, 2018) to fit the emission lines with Gaussians.

The HII regions were selected to have a minimum peak H$\alpha$ flux twice the
standard deviation of a region outside of the galaxy, and a maximum spatial
size of 5 pixels. In Figure 14, we show the selected regions superimposed on
the H$\alpha$ emission. We also show the estimated oxygen abundances in units
of $12+\textrm{log}_{10}(\textrm{O/H})$, based on the calibration of Dopita et
al. (2016), and the star formation rate (SFR), estimated using the H$\alpha$
luminosity as described in Kennicutt (1998), for each HII region. SN 2011fh is
located in the outskirts of the galaxy, and the oxygen abundance value of its
nearest HII region is $8.39\pm 0.03$, or $0.5\ \textrm{Z}_{\odot}$, assuming
the Solar oxygen abundance of $12+\textrm{log}_{10}(\textrm{O/H})=8.69$
(Asplund et al., 2021). This value is surprisingly similar to the metallicity
estimated for SN 2009ip by Margutti et al. (2014). The SFR is fairly constant
throughout the galaxy, with a median value of $1.8\times 10^{-3}$ M⊙ yr-1,
although there is one HII region with a high value of $4.8\times 10^{-2}$ M⊙
yr-1. SN 2011fh is located near a region with relatively high SFR of $\sim
2.3\times 10^{-3}$ M⊙ yr-1.

Figure 14: The left panel shows the HII regions selected by ifuanalysis
superimposed on the H$\alpha$ emission map. A $6\arcsec$ region around SN
2011fh was masked to avoid contamination from the SN. The middle panel shows
the estimated star formation rate for these regions, in M⊙ yr-1, and the right
panel shows the metallicity in units of $12+\textrm{log}_{10}(\textrm{O/H})$.
In the last two panels, the position of SN 2011fh is marked with a red star.

####  3.4.2 HST Analysis

As can be seen in Figure 5, there are a number of bright sources within $\sim
150$ pc of SN 2011fh, probably forming an association with its progenitor
star. We used DOLPHOT (Dolphin, 2016) to perform PSF photometry on these
sources, using the same selection criteria as described in Section 2.4. We
selected a total of 10 stars within a $1\arcsec$ radius (150 pc) centered in
SN 2011fh. The extinction corrected absolute magnitudes of the stars in the
F336W and F814W filters are shown in Table 7.

We used HOKI (Stevance et al., 2020a), and its modules CMD and AgeWizard
(Stevance et al., 2020b), to estimate the age of the stellar cluster around SN
2011fh, and thus constrain the age of its progenitor. HOKI uses stellar
population models from Binary Population and Spectral Synthesis (BPASS,
Eldridge et al., 2017; Stanway & Eldridge, 2018).

The selected sources are very blue, which suggests a very young progenitor
star for SN 2011fh. In Figure 15, we show the $\textrm{M}_{F336W}$ versus
$\textrm{M}_{F336W}-\textrm{M}_{F814W}$ color-magnitude diagram of the HST
sources, along with the HOKI synthetic CMD. Figure 15 shows that the detected
sources are limited to a magnitude of $\textrm{M}_{F336W}\lesssim-8.4$ mag,
which affects the detection of fainter sources and places an uncertainty in
our analysis. We notice that the crowding of the stars in the region might
also affects the detection of sources. We set the metallicity to a value of
$0.5\ \textrm{Z}_{\odot}$ and found the best-fitting BPASS model to have an
age of $10^{6.7}$ yr. The model was then used in AgeWizard to estimate the
probability distribution of ages based on the observed color-magnitude values.
The resulting distribution is shown in the bottom panel of Figure 15. The
$90\%$ confidence interval spans between $10^{6.4}-10^{6.8}$ yr, with a median
value of $10^{6.6}$ yr. This result indicates that the local stellar
population, and therefore the progenitor of SN 2011fh, is very young.

Table 7: HST photometry of the stars around SN 2011fh  F336W | F814W
---|---
(mag) | (error) | (mag) | (error)
$-10.57$ | $0.10$ | $-9.30$ | $0.07$
$-9.10$ | $0.16$ | $-8.63$ | $0.08$
$-9.23$ | $0.13$ | $-8.23$ | $0.09$
$-9.52$ | $0.12$ | $-7.73$ | $0.11$
$-9.19$ | $0.14$ | $-7.83$ | $0.11$
$-9.13$ | $0.14$ | $-7.73$ | $0.11$
$-8.46$ | $0.19$ | $-7.83$ | $0.10$
$-9.10$ | $0.19$ | $-7.38$ | $0.17$
$-8.33$ | $0.21$ | $-7.49$ | $0.13$
$-8.35$ | $0.21$ | $-7.26$ | $0.15$

Figure 15: The top panel shows the color magnitude diagram (CMD) of the bright
stars within $1\arcsec$ (150 pc) of SN 2011fh, along with a HOKI CMD with an
age of $10^{6.7}$. The bottom panel shows the AgeWizard probability
distribution. The 90 percent confidence interval of the distribution is
highlighted in the red region, and the median is marked as the dashed line.

##  4 Discussion

The members of the family of SN 2009ip-like transients show remarkable
similarities throughout their evolution, which may be indicative of similar
progenitors or explosion mechanisms. In Section 4.1, we describe the
photometric and spectroscopic similarities of SN 2011fh to SN 2009ip and
similar events. We discuss the possibility of dust formation in Section 4.2.
We analyze the late-time spectrum of SN 2011fh in Section 4.3, where we also
make a parallel with LBV eruptions. The implications of the age estimate for
the progenitor star, such as the progenitor mass and possible explosion
mechanisms, are discussed in Section 4.4.

###  4.1 The relation with SN 2009ip-like events

As described in Section 3.2 and shown in Figure 7, SN 2011fh, SN 2009ip
(Mauerhan et al., 2013; Fraser et al., 2013), SN 2016bdu (Pastorello et al.,
2017), and SN 2018cnf (Pastorello et al., 2019b) all show a similar
brightening event (Event A) with M${}_{r}\sim-15$ mag (M${}_{r}\sim-16$ mag
for SN 2011fh), followed by a luminous outburst (Event B) with
M${}_{r}\sim-18$ mag. SN 2009ip had a peak luminosity of $\sim 2\times
10^{43}\ \textrm{erg s}^{-1}$ (Mauerhan et al., 2013; Margutti et al., 2014),
while SN 2016bdu and SN 2018cnf peaked at $\sim 4\times 10^{42}\ \textrm{erg
s}^{-1}$ (Pastorello et al., 2017, 2019b), and SN 2011fh at $\sim 3.4\times
10^{42}\ \textrm{erg s}^{-1}$ (Figure 9).

The interpretation of the precursor event and its connection to the brighter
one is still unclear, with ideas ranging from an explosion (faint core-
collapse or a major outburst in a massive star) followed by ejecta-CSM
interaction (Moriya, 2015; Elias-Rosa et al., 2016; Pastorello et al., 2017),
an outburst preceding a SN (Ofek et al., 2013; Pastorello et al., 2017), or
the result of binary interactions (Kashi et al., 2013; Levesque et al., 2014;
Goranskij et al., 2016). The slow brightening seen for SN 2011fh in the months
before peak might indicate an LBV-like eruptive state shortly before a major
outburst, as it is commonly observed for SN impostors (Smith et al., 2011), or
even Type IIn SNe (Ofek et al., 2014; Strotjohann et al., 2021). The mechanism
behind Event B, in this case, could be either the shock generated by a real
core-collapse of the progenitor star, or some energetic eruption originating
in instabilities that are thought to operate in massive stars, such as
pulsational pair-instability (Woosley et al., 2007), unsteady burning or
interaction with other stars (Smith, 2014; Smith & Arnett, 2014). We cannot
rule out the possibility that the first brightening was in fact the main
explosive event, and that the following event was generated by its shock
interacting with CSM shells. This scenario was invoked by Pastorello et al.
(2017) in the case of SN 2016bdu, and by Elias-Rosa et al. (2016) for SN
2015bh. Moriya (2015) described a similar scenario in which the Event B
observed in SN 2009ip was generated by shell-shell collisions. Unfortunately,
the slow brightening event of SN 2011fh was only observed in $r$-band, and the
lack of other colors or spectroscopic observations during this period makes it
hard to constrain its true nature.

Pre-discovery images of SN 2009ip and SN 2016bdu have shown LBV-like
variability for their progenitors, with observed outbursts reaching
M${}_{r}\sim-14$ mag (Mauerhan et al., 2013; Pastorello et al., 2017). In
Figure 2 we have shown that nearly five years before the luminous outburst of
2011 the source ranged between M${}_{r}\sim-14.5$ mag and M${}_{r}\sim-15.5$
mag. We note, however, that these observations are relatively shallow, and
these oscillations might not be real. For this reason, it is not possible to
conclude if such behavior is related to a physical process in the progenitor
star.

In the days following peak brightness, SN 2009ip, SN 2016bdu, and SN 2018cnf
all show a very similar “shoulder” in their light curve (see Figure 7),
probably caused by a collision between CSM shells (Mauerhan et al., 2013;
Fraser et al., 2013; Pastorello et al., 2017, 2019b). Although photometry of
SN 2011fh is not available at this period, bumps in its $r$-band light curve
are seen at $\sim 228$ and $\sim 350$ days after peak, probably generated by
the same mechanism. The late-time brightness evolution of SN 2009ip and SN
2016bdu is very slow, probably due to ongoing interaction with a dense CSM
(Graham et al., 2017; Pastorello et al., 2017). SN 2011fh also shows a slow
evolution at later times, with an apparent flattening of the light curve after
2014. The 2016 October 20 HST observations revealed a source with an absolute
magnitude of M${}_{F814W}\approx-13.32$ at the position of SN 2011fh. This is
more luminous than the expected for a star in a quiescent stage, which might
suggest that the progenitor star is going through an LBV-like eruptive state.
However, emission due to continuous interaction between the ejecta and the CSM
is a common mechanism used to explain late-time emission of Type IIn SNe and
other interacting transients (Taddia et al., 2013; Fraser, 2020), and might
explain this late-time emission for SN 2011fh.

Besides their photometric correspondences, SN 2009ip-like events also show
remarkable similarities in their spectroscopic evolution. In Figure 16, we
show the comparison of the spectral evolution of SN 2011fh to SN 2009ip
(Fraser et al., 2013; Margutti et al., 2014; Fraser et al., 2015; Graham et
al., 2017), SN 2016bdu (Pastorello et al., 2017), and SN 2018cnf (Pastorello
et al., 2019b). The four objects are very similar near peak brightness, with
blue continua and prominent narrow to intermediate width Balmer emission
features. SN 2011fh and SN 2018cnf, however, have lower blackbody
temperatures, with a measured $\sim 8.3\times 10^{3}$ K for the latter, while
SN 2009ip and SN 2016bdu show temperatures of $\sim 1.2\times 10^{4}$ K and
$\sim 1.7\times 10^{4}$ K, respectively (Graham et al., 2014; Pastorello et
al., 2017).

At $+129$ days, SN 2011fh has a very distinctive P-Cygni absorption feature,
which indicate the presence of fast-moving material at $\sim 6400\ \textrm{km
s}^{-1}$. A similar feature with a velocity of $\sim 5000\ \textrm{km s}^{-1}$
was seen for SN 2009ip at similar epochs (Margutti et al., 2014). At $+78$
days, SN 2016bdu also shows a prominent absorption feature at $\sim 3800\
\textrm{km s}^{-1}$ (Pastorello et al., 2017). We note that higher velocity
components were detected for both transients at earlier epochs, with measured
FWHM values of up to $\sim 13000\ \textrm{km s}^{-1}$ (Margutti et al., 2014;
Pastorello et al., 2017). Another SN 2009ip-like event, LSQ13zm, presented
even higher measured velocities of up to $\sim 22000\ \textrm{km s}^{-1}$
(Tartaglia et al., 2016).

For spherical symmetry, it is possible to estimate the mass-loss rate of the
progenitor star assuming that ejecta-CSM shock interaction powers the light
curve (e.g., Smith, 2013). Using the CSM luminosity, $L_{CSM}$, the velocity
of the cold dense shell, $V_{CDS}$, and the velocity of the unshocked CSM,
$V_{CSM}$, and assuming 100% efficiency of converting shock kinetic energy
into radiation,

$\dot{M}=2\ L_{CSM}\ \frac{V_{CSM}}{V_{CDS}^{3}}.$ (1)

Using $V_{CDS}\approx 3000\ \textrm{km}\ \textrm{s}^{-1}$, $V_{CSM}\approx
100$ km s-1, and $L_{CSM}\simeq L{peak}\approx 3\times 10^{42}\ \textrm{erg
s}^{-1}$, we get $\dot{M}\sim 4\times 10^{-2}\ \textrm{M${}_{\odot}$
yr}^{-1}$. This value is fully consistent with the estimated mass-loss rate
for the progenitor of SN 2009ip of $\sim 10^{-2}-10^{-1}\
\textrm{M${}_{\odot}$ yr}^{-1}$ (Fraser et al., 2013; Margutti et al., 2014).
The observed mass-loss rate of SN 2011fh also falls within the observed range
of mass-loss rates for Type IIn SNe of $10^{-2}-10^{-1}\ \textrm{M${}_{\odot}$
yr}^{-1}$ (Kiewe et al., 2012) and $10^{-4}-10^{-2}\ \textrm{M${}_{\odot}$
yr}^{-1}$ (Taddia et al., 2013).

As noted by Brennan et al. (2021a), 2009ip-like transients often show
asymmetries in their H$\alpha$ emission profile. The line in AT 2016jbu was
very asymmetric, in SN 2009ip it was very symmetric, and in SN 2016bdu it was
somewhat asymmetric. SN 2011fh shows a degree of H$\alpha$ symmetry somewhat
similar to SN 2016bdu, with a clearly asymmetrical shape until phase $+293$
(See Figure 3 and Figure 12). This behavior of the H$\alpha$ profile might be
related to asymmetries in the CSM

At late-times ($\sim 240$ days after peak), SN 2011fh, SN 2009ip, and SN
2016bdu are still dominated by emission lines of H and He, which indicates the
ongoing interaction with a dense CSM. All three objects show Ca II emission at
this phase, and both SN 2009ip and SN 2016bdu show He I $\lambda$7065
emission, which will develop later in SN 2011fh. The Fe emission line forest
is also visible for the three events, and is particularly strong for SN
2009ip. Pastorello et al. (2017) noted that SN 2016bdu shows nebular lines of
[O I] at this epoch, and that those lines are expected for core-collapse SNe.
The same was true for SN 2011fh. Although differing in some details, the
H$\alpha$ profile of all the three events around this epoch can be described
by a narrow feature superimposed on an intermediate width feature.

As noted by Graham et al. (2017) for SN 2009ip, SN 2011fh also shows a
distinctive O I $\lambda$8446 emission line but no O I $\lambda$7774 at very
late-times. This suggests that O I $\lambda$8446 is being created by Bowen
Ly-$\beta$ fluorescence, which is consistent with a very dense and optically
thick medium. Fraser et al. (2015) showed that SN 2009ip had a ratio between
the intensity of the [O I]$\lambda\lambda 6300,6363$ lines of $\sim 1$.
Although this ratio suggests newly synthesized oxygen in the ejecta, Fraser et
al. (2015) argue that the lack of evolution of this value points towards the
presence of primordial oxygen. We note a very similar scenario for SN 2011fh,
which has a ratio of $0.9\pm 0.3$ at the phase of $+684$ days and $1.6\pm 0.3$
at $+1359$ days. As seen in Figure 16, SN 2011fh and SN 2009ip also show other
similarities at phase $\sim+680$ days, such as the presence of strong Ca II
features and the Fe II emission line forest. The
$\textrm{H}\alpha/\textrm{H}\beta$ flux ratio for SN 2009ip of around 40 at
this epoch is very high (Graham et al., 2017), while we estimate
$\textrm{H}\alpha/\textrm{H}\beta\approx 7$ for SN 2011fh.

Figure 16: Spectral comparison of SN 2011fh to SN 2009ip (Fraser et al., 2013;
Margutti et al., 2014; Fraser et al., 2015; Graham et al., 2017), SN 2016bdu
(Pastorello et al., 2017), and SN 2018cnf (Pastorello et al., 2019b) at peak
brightness and near phases $\sim$ 100, 240, and 680 days.

###  4.2 The presence of dust

In Section 3.2.2, we found that the spectral energy distribution (SED) of SN
2011fh has a significant IR excess that can be described by dust emission with
a temperature of $\sim 950$ K at phase $+570$ days, and $\sim 1116$ K at phase
$+1143$ days.

The dust emission is located at $\sim 10^{16.52}\ \textrm{cm}$ at phase
$+570$, and decreases to $\sim 10^{16.44}\ \textrm{cm}$ at later epochs (see
Table 2). Dust emission is quite common for interacting transients (Prieto et
al., 2008; Berger et al., 2009; Foley et al., 2011; Fox et al., 2011), but
distinguishing between pre-existing or newly formed dust is not easy. SN
2011fh shows a similar mid-IR color evolution to several Type IIn SNe and
other interacting transients, such as SN 2009ip (Szalai et al., 2021).
Asymmetries in the H$\alpha$ caused by the suppression of the red wing
relative to the blue wing are often used as an indicator for the detection of
newly formed dust (e.g., Fox et al., 2011). Although SN 2011fh has small
asymmetries, it is hard to conclude whether it is caused by dust formed in the
luminous outburst. The dust inner radius of $\sim 10^{16.5}\ \textrm{cm}$ is
larger than the photospheric radius of $\sim 10^{14}\ \textrm{cm}$. This
indicates that dust is expanding slowly, if at all, and show little optical
depth evolution, all of which suggests that dust is pre-existing.

The dust mass is

$M_{dust}=\frac{4\pi\tau_{V}R_{dust}^{2}}{\kappa_{V}},$ (2)

where $\kappa_{V}\approx 2000\ \textrm{cm}^{2}/\textrm{g}$ is the dust visual
opacity for silicate dust (Fox et al., 2010). Using the properties estimated
at $+570$ days, we get $\sim 2\times 10^{-3}\ \textrm{M}_{\odot}$. For a
typical gas to dust mass ratio of $\sim 100$, the total mass associated with
the dust shell is of $\sim 0.2\ \textrm{M}_{\odot}$. These values are much
larger than were estimated for the LBV outburst UGC 2773 OT2009-1 (Foley et
al., 2011) and for the SN 2009ip-like event AT 2016jbu (Brennan et al., 2021a,
b), but fall well within the range of dust masses estimated for Type IIn SNe
by Fox et al. (2011).

###  4.3 The late-time spectrum

Type IIn SNe often show spectra dominated by Balmer lines in emission until
very late periods, instead of the nebular lines of heavy elements seen in non-
interacting SNe. Some known exceptions are SN 1998S, which showed emission
lines of [O I] and [Ca II] in the years after the explosion (Leonard et al.,
2000), and SN 1995N, where nebular lines of [O I], [O II], [O III] appeared at
very late times (Fransson et al., 2002). As was shown by Taddia et al. (2013),
Type IIn SNe can show H$\alpha$ features with velocities of $\sim 10^{4}\
\textrm{km s}^{-1}$ and luminosities between $10^{40}-10^{41}\ \textrm{erg
s}^{-1}$ up to 180 days after peak. The spectra of SN impostors are also
dominated by Balmer lines at late times, and many show Ca II emission lines
(Smith et al., 2011). One example is SN 2000ch, which reached
M${}_{R}\sim-12.8$ mag (Wagner et al., 2004) and had broad lines of H$\alpha$
with FWHM velocities of $\sim 1500\ \textrm{km s}^{-1}$ (Smith et al., 2011).

As described in Section 3.3, the very late-time spectrum of SN 2011fh shows
prominent lines of Fe II, Ca II, O I, and He I, while still being dominated by
strong Balmer features. The H$\alpha$ profile at this epoch shows an
intermediate-width component with a FWHM velocity of $\sim 726\ \textrm{km
s}^{-1}$, and a luminosity of $\sim 4\times 10^{39}\ \textrm{erg s}^{-1}$.

Although it is hard to determine whether SN 2011fh was a terminal event, we
note several similarities between its late time spectrum and those of SN
impostors. In Figure 17, we compare the spectrum of SN 2011fh to SN 2000ch at
phase $+1454$ days (Smith et al., 2011), and to the 19th century Great
Eruption of the Galactic LBV star Eta Carinae (Smith et al., 2018a). SN 2011fh
and SN 2000ch are both dominated by narrow to intermediate-width Balmer lines
in emission until very late periods, with some He I features. The continuum
slopes of the two spectra are very similar, indicating a similar photospheric
temperature. SN 2000ch also shows a prominent O I $\lambda 8446$ emission
line, which is also present in SN 2011fh. Although SN 2000ch does not show the
strong Ca II features seen in SN 2011fh, these lines are present in the
spectra of Eta Carinae’s Great Eruption (Smith et al., 2018a), and are strong
features in the spectra of some events arising from lower mass progenitors,
such as SN 2008S (Botticella et al., 2009; Smith et al., 2009), NGC 300-OT
(Berger et al., 2009; Bond et al., 2009), and SN 2002bu (Smith et al., 2011).
We note that SN 2002bu and SN 2008S are sometimes referred to as intermediate-
luminosity red transients (ILRTs), and that many studies support the scenario
of SN 2008S-like events being terminal SN explosions (Adams et al., 2016;
Botticella et al., 2009; Cai et al., 2021). The late time spectrum of SN
2011fh also shows strong similarities with the spectra of UGC 2773 OT2009-1
(Smith et al., 2010; Foley et al., 2011), an analogue to Eta Carinae’s Great
Eruption (Smith et al., 2016b).

After peak, SN 2000ch remained in a plateau at $\sim-10.6$ mag, which could
have been the quiescent magnitude of the progenitor star or a S-Dor like
eruptive state (Wagner et al., 2004; Smith et al., 2011). Pastorello et al.
(2010) described in detail other two bright eruptions related to the star, in
2008 and 2009, when it reached $\sim-12.8$ during the brightest one.
Pastorello et al. (2010) noted many observational similarities with the system
HD 5980, which is a binary consisting of an erupting LBV and a Wolf-Rayet
star, and suggested binary interactions as an explanation for the variability
and eruptions detected for SN 2000ch. A similar scenario is often used to
explain the Great Eruption of Eta Carinae.

Figure 17: Spectral comparison of the very late-time spectrum of SN 2011fh,
taken 1359 days after peak, to the spectrum of SN 2000ch at a similar phase
(Smith et al., 2011). We also show the light echo spectrum of the Great
Eruption of Eta Carinae (Smith et al., 2018a).

###  4.4 Implications of the progenitor age and mass

Our estimates of the age of the stars around SN 2011fh indicate that the
progenitor is a very young star with $\sim 4.5$ Myr. Using $0.5\
\textrm{Z}_{\odot}$ and a $4.5$ Myr MIST (Paxton et al., 2011, 2013, 2015,
2018) single-star isochrone, we get a maximum mass of $\sim 60\
\textrm{M}_{\odot}$. Since the fraction of binary systems is very high at
large stellar masses, we also use BPASS (Eldridge et al., 2017; Stanway &
Eldridge, 2018) to estimate the progenitor mass. Using the age distribution of
$2-6$ Myr (see Figure 15), a half solar metallicity, and after selecting
progenitors that could result in neutron star remnants with $<2.5\
\textrm{M}_{\odot}$ (based on Figure 3 of Abbott et al., 2018) and cutting out
pulsational pair-instability candidates that reached a He core mass of $>40\
\textrm{M}_{\odot}$, we obtained the distribution of the number of stars per
million solar masses as a function of their ZAMS mass, as shown in Figure 18.
The discrete number of ZAMS masses presented in Figure 18 is a consequence of
the grid of initial masses used in the BPASS models (see Eldridge et al.,
2017). The distribution shows that there is a large number of stars populating
the $35-80\ \textrm{M}_{\odot}$ mass range, with the three tallest bars being
between $35-50\ \textrm{M}_{\odot}$. We note, however, that an energetic
outburst such as SN 2011fh would be more likely to occur with stars in the
most massive half of the distribution, even though the relative number of
stars is slightly lower than the less massive half. The $\sim 60\
\textrm{M}_{\odot}$, obtained by single-star models and the $35-80\
\textrm{M}_{\odot}$ range, obtained considering binary systems are
surprisingly similar to the estimates of $50-80\ \textrm{M}_{\odot}$ (Smith et
al., 2010) and of $\gtrsim 60\ \textrm{M}_{\odot}$ (Foley et al., 2011) for SN
2009ip. It also agrees well with the results of Pastorello et al. (2019b), who
showed that the progenitor of SN 2018cnf was consistent with a massive
hypergiant or an LBV. We note that our results contrast with the estimates of
Brennan et al. (2021a, b), who showed that the progenitor of AT 2016jbu was a
relatively less massive hypergiant star with $\sim 20\ \textrm{M}_{\odot}$.
This might indicate that this type of event can be originated by progenitors
with a large range of initial masses.

One common factor used to explain instabilities in LBV stars is their
proximity to the classical Eddington limit, which would help to create strong
winds. Eta Carinae is thought to have exceeded the Eddington limit by a factor
of $\sim 5$ during its Great Eruption (Smith et al., 2018a). We estimate a
still larger Eddington factor of $\sim 30$ for SN 2011fh at the beginning of
the brightening event of February of 2011, when it had a magnitude of
M${}_{r}\approx-15.2$. An Eddington factor of $\sim 100$ is reached at the end
of this period, just before the luminous outburst of August 2011.

Super-Eddington winds can arise from energy deposited in the envelope of
massive stars by different mechanisms, including wave heating and unstable
fusion (Quataert & Shiode, 2012; Shiode & Quataert, 2014; Quataert et al.,
2016). Binary interactions may also be a mechanism for injecting energy into
the outer layers of a massive star, since more than half of all massive stars
are in binary systems with short enough periods to generate some kind of
interaction or collision (Sana et al., 2012; Smith, 2014). The interactions
typically begin when the more massive star evolves off the main sequence. This
is a scenario often used to explain the behavior of Eta Carinae (Smith et al.,
2018a; Smith & Frew, 2011).

The pulsational pair instability mechanism is capable of explaining the
extreme mass loss seen in LBV stars (Woosley et al., 2007), and was considered
as an explanation for SN 2009ip by Pastorello et al. (2013), Mauerhan et al.
(2013) and Fraser et al. (2013). Margutti et al. (2014), however, demonstrated
that, for a mass limit of 85 M⊙ and a metallicity of $0.4-0.9$ Z⊙, the
progenitor of SN 2009ip would have a helium core mass lower than the $\sim 40$
M⊙ needed for a pulsational pair instability. Since SN 2011fh is so similar to
SN 2009ip, the pulsational pair instability seems as an unlikely mechanism for
its bright eruptions.

Finally, we cannot fully exclude the possibility of SN 2011fh being a weak SN,
although this appears less likely. Under-energetic explosions arising from
core-collapse are expected to occur in such massive stars (Heger et al., 2003;
Moriya et al., 2010; Sukhbold et al., 2016). Mauerhan et al. (2013) used this
mechanism to explain the double-peaked light curve of SN 2009ip, where Event A
would be a weak core-collapse event, and Event B would be powered by the
ejecta shock interacting with a dense CSM. Smith et al. (2014) noted a
similarity of the Event A of SN 2009ip to SN 1987A, a SN generated by a blue
supergiant star (Arnett et al., 1989). A similar statement was made by Elias-
Rosa et al. (2016) in the case of SN 2015bh. Pastorello et al. (2017) showed
that other SN 2009ip-like events, such as SN 2016bdu, SN 2010mc (Smith et al.,
2014), and LSQ13zm (Tartaglia et al., 2016), have Event A light curves that
are very similar to the faint SN 1987A-like SN 2009E (Pastorello et al.,
2012). SN 2011fh, however, does not appear to present such similarities.

Figure 18: Number of stars per million M⊙ obtained using BPASS. The discrete
number of ZAMS masses come from the grid of initial masses used in the BPASS
models. The highest bars in the distribution extend between 35 and 80 M⊙.

##  5 Conclusion

In this paper we presented optical, near-infrared and infrared observations of
the Type IIn SN 2011fh, spanning from 2007 to 2017. The results can be
summarized as follows:

  * •

SN 2011fh shows several similarities with the peculiar SN 2009ip and similar
transients, such as SN 2016bdu and SN 2018cnf. They include: 1) an initial
spectral resemblance to typical Type IIn SNe, with a blue continuum and narrow
to intermediate width Balmer emission lines; 2) a brightening event that lasts
$\sim 5$ months and reaches M${}_{r}\sim-16$ mag, followed by a luminous
outburst with M${}_{r}\sim-18$ mag at peak; 3) a very similar light curve,
with bumps and a slow fading with time; and 4) a very similar spectroscopic
evolution, with signals of CSM interaction until very late times, distinctive
P-Cygni features indicating a fast-moving material, and the presence of Ca II
and Fe II lines.

  * •

The estimated progenitor mass-loss rate of $\sim 4\times 10^{-2}$ M⊙ yr-1 is
consistent with mass-loss rate estimates for SN 2009ip and Type IIn SNe.

  * •

SN 2011fh shows a significant IR excess that can be described by warm dust
emission with $T_{dust}\approx 1000$ K and $R_{dust}\approx 3\times 10^{16}$
cm. The emission was probably created by pre-existing dust with a mass of
$\textrm{M}_{dust}\sim 2\times 10^{-3}$ M⊙.

  * •

The very late-time spectrum of SN 2011fh shows strong similarities with some
SN impostors observed at a same phase. The detection of narrow Ca II features
and Balmer emission lines is also similar to the Great Eruption of Eta
Carinae.

  * •

SN 2011fh is located in a region with relatively rapid star formation ($\sim
2\times 10^{-3}$ M⊙ yr-1). The H II region oxygen abundance of
$12+\textrm{log(O/H)}=8.39$ ($\sim 0.5\ \textrm{Z}_{\odot}$) is similar to
what was estimated at the location of SN 2009ip.

  * •

The local stellar population has an age of $\sim 4.5$ Myr, which corresponds
to a progenitor with an initial main-sequence mass of $\sim 60$ M$\odot$, if
we consider a single-stellar evolution model, or a range of $35-80\
\textrm{M}_{\odot}$, considering binary systems. These mass estimates are
similar to the estimates for SN 2009ip’s progenitor, and consistent with very
massive stars passing through the LBV phase.

  * •

SN 2011fh exceeded the classical Eddington limit by a large factor in the
months before the luminous outburst of 2011. This suggests that a possible
mechanism behind the bright event is related to strong supper-Eddington winds.

Although SN 2011fh shows striking similarities with both impostors and 2009ip-
like events, it is hard to conclude whether SN 2011fh was a genuine core-
collapse SN, or a non-terminal event where the mass recently ejected collides
with pre-existing CSM and generates a luminous outburst. The lack of
observational signatures related to a core-collapse, such as broad Balmer and
O I lines at late times, makes the scenario of a true SN less likely. Our
results therefore imply that at least a fraction of SN 2009ip-like events
arise from non-terminal eruptions in massive and young stars. The discovery
and follow up observations of new transients similar to SN 2009ip will
certainly help in the understanding of the mechanisms behind this unique class
of events.

We would like to thank Andrea Pastorello and Melissa Graham for kindly sharing
their data on SN 2016bdu, SN 2018cnf, and SN 2009ip. We thank the Carnegie
Supernova Project-II for obtaining several photometric observations presented
in the paper, and in particular Carlos Contreras, Mark Phillips, Nidia Morrell
and Eric Hsiao. TP is supported by CONICYT’s Programa de Astronomía through
the ALMA-CONICYT 2019 grant 31190017. Support for JLP is provided in part by
ANID through the Fondecyt regular grant 1191038 and through the Millennium
Science Initiative grant ICN12009, awarded to The Millennium Institute of
Astrophysics, MAS. CSK is supported by NSF grants AST-1814440 and AST-1907570.
HFS acknowledges the support of the Marsden Fund Council managed through Royal
Society Te Aparangi. This paper includes optical and NIR photometry obtained
by the Carnegie Supernova Project, which was generously supported by NSF
grants AST-1008343, AST-1613426, AST-1613455, and AST-1613472.

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##  Appendix A Contamination

SN 2011fh is located in a relatively crowded region of the NGC 4806 galaxy,
with several bright stars in the proximity (see Figure 5). These stars lie
within 150 pc ($1\arcsec$) around SN 2011fh, with several brighter sources
within 300 pc ($2\arcsec$). Since the pixel scale of the pre SN observations
were limited up to $2\arcsec$, we decided to estimate the possible
contamination of the stellar cluster to our photometric measurements.

To quantify the amount of contamination in the optical bands, we first
performed aperture photometry of the stars around SN 2011fh in the HST F336W
and F814W images, using the Astropy (Astropy Collaboration et al., 2013, 2018)
package Photutils (Bradley et al., 2019) in a $2\arcsec$ region. After
subtracting the flux of SN 2011fh, we obtained a contribution from the stars
in this aperture of 20.67 and 20.05 mag, in the F336W and F814W filters,
respectively. By assuming a flat spectrum for the stellar cluster, we then
interpolated between the F336W and F814W filter fluxes and performed synthetic
photometry in the resulting line with PYPHOT999
https://mfouesneau.github.io/docs/pyphot/, to obtain the estimates of
$V\approx 20.17$, $r\approx 20.18$, and $i\approx 20.18$. If we compare these
values with the PSF photometry of SN 2011fh presented in Tables 8 and 9, the
estimated contribution from the cluster is significantly lower than the flux
of SN 2011fh throughout its photometric evolution. ]

##  Appendix B Tables

Table 8: Amateur astronomers data and CRTS $r$-band photometry of SN 2011fh  UT Date | JD | r
---|---|---
|  | (mag) | (error)
2007-1-25 | 2454126.20298 | 17.879 | 0.12
2007-2-22 | 2454154.20024 | 17.988 | 0.13
2007-3-9 | 2454169.17404 | 18.272 | 0.178
2007-3-27 | 2454187.19841 | 17.961 | 0.125
2007-4-6 | 2454197.17815 | 18.547 | 0.25
2007-4-16 | 2454207.0734 | 18.184 | 0.16
2007-5-7 | 2454228.04447 | 18.278 | 0.177
2007-5-14 | 2454235.05647 | 18.097 | 0.143
2007-6-16 | 2454267.97322 | 17.943 | 0.131
2007-7-4 | 2454285.85462 | 17.999 | 0.134
2007-7-20 | 2454301.85781 | 18.232 | 0.166
2007-8-12 | 2454324.88361 | 17.985 | 0.129
2007-12-28 | 2454463.18063 | 18.09 | 0.136
2008-1-22 | 2454488.1978 | 18.273 | 0.182
2008-2-7 | 2454504.25914 | 18.057 | 0.135
2008-2-8 | 2454505.25765 | 18.121 | 0.148
2008-2-29 | 2454526.17647 | 17.988 | 0.135
2008-3-11 | 2454537.18612 | 18.286 | 0.179
2008-5-6 | 2454593.06685 | 17.937 | 0.127
2008-5-11 | 2454598.10991 | 17.912 | 0.131
2008-5-28 | 2454615.0704 | 18.239 | 0.179
2008-6-27 | 2454644.9809 | 17.992 | 0.14
2008-12-31 | 2454832.20995 | 18.118 | 0.15
2009-1-5 | 2454837.49042 | 18.121 | 0.243
2009-2-25 | 2454888.21891 | 18.207 | 0.165
2009-3-15 | 2454906.27249 | 17.734 | 0.106
2009-3-27 | 2454918.21621 | 17.877 | 0.127
2009-4-15 | 2454937.1613 | 18.235 | 0.172
2009-4-27 | 2454949.13822 | 17.857 | 0.11
2009-4-29 | 2454951.09237 | 17.945 | 0.138
2009-5-11 | 2454963.0813 | 18.071 | 0.156
2009-5-12 | 2454964.30533 | 17.799 | 0.163
2009-5-29 | 2454981.11436 | 17.79 | 0.126
2009-6-19 | 2455002.00408 | 18.2 | 0.164
2009-7-18 | 2455030.93214 | 18.241 | 0.166
2010-1-21 | 2455218.17694 | 17.777 | 0.101
2010-2-21 | 2455249.13263 | 18.289 | 0.171
2010-3-9 | 2455265.12761 | 17.893 | 0.118
2010-4-8 | 2455295.17376 | 18.055 | 0.131
2010-4-17 | 2455304.11246 | 18.077 | 0.144
2010-5-10 | 2455327.01388 | 17.666 | 0.091
2010-5-22 | 2455339.01684 | 17.92 | 0.118
2010-6-13 | 2455361.0242 | 17.762 | 0.107
2010-6-30 | 2455378.24321 | 17.948 | 0.189
2010-7-16 | 2455393.9443 | 17.752 | 0.118
2010-7-22 | 2455399.93913 | 17.83 | 0.112
2010-7-28 | 2455406.20502 | 17.969 | 0.2
2011-2-11 | 2455603.5911 | 17.636 | 0.146
2011-2-24 | 2455617.25406 | 17.216 | 0.061
2011-2-24 | 2455617.4946 | 17.549 | 0.102
2011-4-13 | 2455665.05019 | 17.037 | 0.054
2011-4-24 | 2455676.10863 | 17.176 | 0.063
2011-4-25 | 2455677.4096 | 17.271 | 0.098
2011-5-9 | 2455691.0465 | 16.845 | 0.052
2011-5-26 | 2455708.02969 | 16.936 | 0.053
2011-6-2 | 2455715.3151 | 16.716 | 0.075
2011-6-3 | 2455716.22604 | 16.883 | 0.072
2011-6-16 | 2455728.86432 | 16.855 | 0.089
2011-6-21 | 2455734.03201 | 16.671 | 0.073
2011-6-27 | 2455739.96366 | 16.694 | 0.045
2011-6-28 | 2455741.29443 | 16.799 | 0.067
2011-6-30 | 2455742.91587 | 16.861 | 0.057
2011-6-30 | 2455742.91587 | 16.837 | 0.058
2011-7-17 | 2455759.84951 | 16.614 | 0.048
2011-7-28 | 2455770.90473 | 16.264 | 0.036
2011-8-2 | 2455776.20691 | 16.342 | 0.068
2011-8-9 | 2455782.84013 | 15.726 | 0.037
2011-8-24 | 2455798.24181 | 14.942 | 0.028
2011-8-25 | 2455799.24933 | 14.962 | 0.029
2011-9-4 | 2455809.21792 | 15.362 | 0.031
2011-9-5 | 2455810.22498 | 15.321 | 0.033
2011-12-16 | 2455911.57505 | 16.731 | 0.046
2011-12-25 | 2455920.58029 | 16.704 | 0.041
2012-1-6 | 2455932.58682 | 16.748 | 0.041
2012-1-25 | 2455951.5711 | 17.092 | 0.051
2012-1-28 | 2455954.56738 | 17.373 | 0.06
2012-2-16 | 2455973.54651 | 17.275 | 0.061
2012-2-17 | 2455975.08906 | 17.701 | 0.151
2012-2-24 | 2455981.55593 | 17.543 | 0.078
2012-3-25 | 2456012.45263 | 17.469 | 0.068
2012-3-30 | 2456017.15598 | 17.379 | 0.073
2012-4-10 | 2456028.43419 | 16.886 | 0.047
2012-4-11 | 2456029.14365 | 16.912 | 0.055
2012-4-19 | 2456037.0854 | 17.121 | 0.062
2012-5-15 | 2456063.026 | 17.33 | 0.073
2012-7-8 | 2456116.95661 | 17.984 | 0.21
2012-8-12 | 2456151.87688 | 17.671 | 0.104
2012-8-12 | 2456152.21301 | 17.248 | 0.063
2012-12-17 | 2456278.54791 | 18.305 | 0.173
2012-12-21 | 2456282.56931 | 18.225 | 0.167
2013-1-1 | 2456294.20582 | 18.136 | 0.154
2013-1-22 | 2456314.55775 | 18.103 | 0.183
2013-2-7 | 2456330.57937 | 17.952 | 0.106
2013-2-18 | 2456341.58446 | 18.258 | 0.161
2013-2-21 | 2456344.5748 | 18.055 | 0.124
2013-2-24 | 2456348.40002 | 18.19 | 0.236
2013-3-2 | 2456354.49434 | 18.059 | 0.203
2013-3-9 | 2456361.45919 | 17.904 | 0.164
2013-3-17 | 2456368.53775 | 18.08 | 0.164
2013-4-11 | 2456394.39371 | 18.169 | 0.142
2013-4-19 | 2456402.41645 | 17.842 | 0.106
2013-4-30 | 2456413.32136 | 18.037 | 0.134
2013-5-1 | 2456414.3878 | 18.168 | 0.167
2013-5-2 | 2456415.27285 | 18.079 | 0.152
2013-5-4 | 2456417.32889 | 18.428 | 0.221
2013-5-5 | 2456418.36387 | 18.204 | 0.146
2013-5-10 | 2456423.27335 | 17.848 | 0.117
2013-5-11 | 2456424.21595 | 18.38 | 0.201
2013-5-12 | 2456425.26083 | 17.842 | 0.098
2013-5-14 | 2456427.23059 | 18.231 | 0.178
2013-5-16 | 2456429.43702 | 18.163 | 0.146
2013-5-31 | 2456444.31843 | 18.282 | 0.171
2013-6-3 | 2456447.31712 | 18.319 | 0.173
2013-6-5 | 2456449.28426 | 18.329 | 0.169
2013-6-15 | 2456459.21472 | 18.617 | 0.289
2013-6-16 | 2456460.28568 | 18.334 | 0.193
2013-6-26 | 2456470.25602 | 18.18 | 0.15
2013-6-28 | 2456472.28668 | 18.422 | 0.194
2013-6-29 | 2456473.2751 | 17.967 | 0.108
Table 9: $BVri$-CSP2 photometry of SN 2011fh  UT Date | JD | B | V | r | i
---|---|---|---|---|---
|  | (mag) | (error) | (mag) | (error) | (mag) | (error) | (mag) | (error)
2013-2-20 | 2456343.85 | 18.843 | 0.0215 | 18.493 | 0.0173 | 18.063 | 0.0132 | 18.4135 | 0.0151
2013-3-3 | 2456354.78 | 18.8695 | 0.0226 | 18.5055 | 0.0193 | 18.0805 | 0.0142 | 18.41 | 0.0185
2013-3-15 | 2456366.71 | 18.9345 | 0.0244 | 18.576 | 0.0193 | 18.1375 | 0.0154 | 18.5 | 0.0202
2013-3-18 | 2456369.91 | - | - | - | - | 18.15 | 0.0155 | 18.407 | 0.0324
2013-3-25 | 2456376.76 | - | - | - | - | 18.155 | 0.0146 | 18.529 | 0.0213
2013-4-5 | 2456387.68 | 18.913 | 0.0193 | 18.573 | 0.0195 | 18.1565 | 0.0132 | 18.525 | 0.0171
2013-4-20 | 2456402.73 | 18.962 | 0.0265 | 18.5725 | 0.0244 | 18.193 | 0.0182 | 18.522 | 0.0324
2013-5-4 | 2456416.68 | 19.1855 | 0.0273 | 18.8505 | 0.0245 | 18.1565 | 0.0132 | 18.816 | 0.0232
2013-5-12 | 2456424.75 | 19.103 | 0.0233 | 18.795 | 0.0243 | 18.3765 | 0.0151 | 18.795 | 0.0251
2013-6-1 | 2456444.68 | 19.046 | 0.0254 | - | - | 18.341 | 0.0192 | 18.717 | 0.0291
2015-3-16 | 2457097.76 | 19.7675 | 0.0205 | 19.559 | 0.0165 | 18.341 | 0.0192 | 19.558 | 0.0177
2015-3-18 | 2457099.80 | 19.883 | 0.0266 | 19.5225 | 0.0167 | - | - | - | -
2016-1-14 | 2457401.80 | 19.625 | 0.0341 | 19.3705 | 0.0306 | 19.229 | 0.0342 | 19.475 | 0.03815
Table 10: $YJH$-CSP2 photometry of SN 2011fh  UT Date | JD | Y | J | H
---|---|---|---|---
|  | (mag) | (error) | (mag) | (error) | (mag) | (error)
2013-2-23 | 2456346.86 | 17.62 | 0.031 | - | - | - | -
2013-2-27 | 2456350.81 | 17.58 | 0.03 | 17.457 | 0.026 | 17.19 | 0.037
2013-3-26 | 2456377.75 | 17.75 | 0.023 | 17.649 | 0.021 | - | -
2015-4-30 | 2457142.58 | 18.865 | 0.044 | 18.833 | 0.051 | 18.491 | 0.0788
2016-3-27 | 2457474.70 | 18.575 | 0.075 | 18.64 | 0.085 | 18.25 | 0.113
2017-5-16 | 2457889.50 | 18.208 | 0.043 | - | - | - | -
Table 11: Spitzer photometry of SN 2011fh UT Date | JD | Channel 1 (3.6 $\mu$m) | Channel 2 (4.5 $\mu$m)
---|---|---|---
|  | (mag) | (error) | (mag) | (error)
2009-8-6 | 2455050.094 | 16.019 | 0.047 | 15.686 | 0.045
2009-9-5 | 2455080.09 | 16.142 | 0.043 | 15.77 | 0.047
2013-3-31 | 2456383.113 | 15.899 | 0.045 | 15.024 | 0.04
2014-4-4 | 2456752.49 | 16.27 | 0.048 | 15.572 | 0.04
2014-9-12 | 2456913.073 | 16.228 | 0.059 | 15.492 | 0.045
2014-10-12 | 2456943.206 | 16.24 | 0.052 | 15.529 | 0.044
2015-5-16 | 2457158.625 | 16.265 | 0.052 | 15.64 | 0.045