# THE LUMINOUS TYPE IIN SUPERNOVA SN 2017hcc: INFRARED BRIGHT, X-RAY AND RADIO
FAINT

Poonam Chandra,1,2 Roger A. Chevalier,3 Nicholas J. H. James3 and Ori D. Fox4
1National Radio Astronomy Observatory, 520 Edgemont Road, Charlottesville, VA
22903, USA
2National Centre for Radio Astrophysics, Tata Institute of Fundamental
Research, Ganeshkhind , Pune 411007, India
3Department of Astronomy, University of Virginia, P.O. Box 400325,
Charlottesville, VA 22904-4325, USA
4Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD
21218, USA E-mail<EMAIL_ADDRESS>
(Accepted XXX. Received YYY; in original form ZZZ)

###### Abstract

We present multiwavelength observations of supernova (SN) 2017hcc with the
Chandra X-ray telescope and the X-ray telescope onboard Swift (Swift-XRT) in
X-ray bands, with the Spitzer and the TripleSpec spectrometer in near-infrared
(IR) and mid-IR bands and with the Karl G. Jansky Very Large Array (VLA) for
radio bands. The X-ray observations cover a period of 29 to 1310 days, with
the first X-ray detection on day 727 with the Chandra. The SN was subsequently
detected in the VLA radio bands from day 1000 onwards. While the radio data
are sparse, synchrotron-self absorption is clearly ruled out as the radio
absorption mechanism. The near- and the mid-IR observations showed that late
time IR emission dominates the spectral energy distribution. The early
properties of SN 2017hcc are consistent with shock breakout into a dense mass-
loss region, with $\dot{M}\sim 0.1$ $M_{\odot}$ yr-1 for a decade. At few 100
days, the mass-loss rate declined to $\sim 0.02$ $M_{\odot}$ yr-1, as
determined from the dominant IR luminosity. In addition, radio data also
allowed us to calculate a mass-loss rate at around day 1000, which is two
orders of magnitude smaller than the mass-loss rate estimates around the
bolometric peak. These values indicate that the SN progenitor underwent an
enhanced mass-loss event a decade before the explosion. The high ratio of IR
to X-ray luminosity is not expected in simple models and is possible evidence
for an asymmetric circumstellar region.

###### keywords:

supernovae: general — supernovae: individual (SN 2017hcc) — circumstellar
matter — radiation mechanisms: general — radiation mechanisms: non-thermal

††pubyear: 2022††pagerange: THE LUMINOUS TYPE IIN SUPERNOVA SN 2017hcc:
INFRARED BRIGHT, X-RAY AND RADIO FAINT–THE LUMINOUS TYPE IIN SUPERNOVA SN
2017hcc: INFRARED BRIGHT, X-RAY AND RADIO FAINT

##  1 Introduction

The IIn class of supernovae (SNe IIn) was identified by Schlegel (1990) based
on their narrow optical emission lines and hot continua. The narrow lines can
be attributed to slow-moving circumstellar (CS) matter around the supernova
(SN) and the luminosity can be identified with power provided by shock waves.
If the cooling time for the gas is fast compared to the age, the SN power is
efficiently converted to radiation in this class of supernovae (SNe). The
required density of the circumstellar medium (CSM) is high in several SNe IIn
(e.g. Fransson et al., 2014) and are generally not seen in Galactic sources
outside of Luminous Blue Variables (LBVs) during outburst (Smith & Owocki,
2006; van Marle et al., 2008). Radioactivity is not a suitable power source
for these SNe because it predicts a faster decline than observed.

In general, the first electromagnetic signal in SNe occurs when the shock
reaches the SN photosphere and breaks out on time scales of minutes to hours.
However, in dense winds shock breakout can happen in the CSM and last for an
extended duration. There is increasing evidence that in some SNe IIn, the
early light curves can be powered by shock breakouts (Chevalier & Irwin, 2011;
Ofek et al., 2014; Waxman & Katz, 2017).

The early optical spectra of SNe IIn often show emission lines from a wind
that are broadened by electron scattering (Chugai, 2001). Photons from the
inner part of the mass-loss region scatter as they leave the mass-loss region.
An electron column optical depth of about a few, or $10^{24}$ electrons cm-2
column density, is needed to be consistent with the observations. At later
times ($>100$ days), the line profiles are no longer consistent with electron
scattering. Profiles are often fit by multiple Gaussians (e.g., Kiewe et al.,
2012; Szalai et al., 2021). Although reasonable fits can be obtained, there is
not a clear theory that points to Gaussian line shapes. Velocities
corresponding to the line widths are larger than the thermal velocities in the
optically emitting gas.

Considering the fast shock waves moving into a dense CSM, strong X-ray
emission can be expected. Indeed, SNe IIn are well represented among the most
X-ray luminous SNe (Dwarkadas & Gruszko, 2012; Chandra, 2018). However, there
are optically luminous SNe IIn that are weak X-ray emitters; a prime example
is SN 2006gy (Ofek et al., 2007; Smith et al., 2007). In some cases, the
paucity of X-rays can be explained by the effects of high X-ray absorption due
to the large column density through the CS matter (Chevalier & Irwin, 2012;
Svirski et al., 2012). Low X-ray luminosity can also be attributed to a
relatively low density CSM, as in the case of SN 1998S (Pooley et al., 2002).
Chevalier & Irwin (2011) have shown that in some SNe high optical luminosity
and low X-ray luminosity can be explained by shock breakout in dense
surroundings.

SNe IIn can also be radio emitters from synchrotron radiation by shock
accelerated electrons. Many high luminosity SNe IIn are also found to be weak
emitters in radio bands (Chandra et al., 2015). While one expects a high
intrinsic radio luminosity owing to a dense CSM, the same high density CSM
could lead to efficient absorption of radio emission, e.g., the bright Type
IIn SN 2010jl became radio bright only after 500 days (Chandra et al., 2015).
Thus radio emission in SNe IIn is likely to be an interplay between the two.
Internal free-free absorption can also play an important role in these SNe due
to efficient mixing of cool gas between forward and reverse shock into the
synchrotron emitting region (Weiler et al., 1990; Chandra et al., 2012;
Chandra et al., 2020).

Gerardy et al. (2002) first described the late bright infrared (IR) emission
from SNe IIn in which the IR light dominates the spectral energy distribution
(SED). Such emission has been observed in many SNe IIn, including SN 1995N
(Van Dyk, 2013), SN 1998S (Gerardy et al., 2000; Fassia et al., 2000; Mauerhan
& Smith, 2012), SN 2006jd (Fox et al., 2011; Stritzinger et al., 2012), SN
2007rt (Trundle et al., 2009; Szalai et al., 2021), SN 2010jl (Fransson et
al., 2014; Bevan et al., 2020), SN 2005ip (Fox et al., 2009, 2010; Stritzinger
et al., 2012; Bak Nielsen et al., 2018), SN 2013L (Andrews et al., 2017) etc.,
beginning at an age of $\sim 1$ yr. The IR emission in these SNe has an
approximately blackbody distribution, with a temperature $T$ in the range
$\sim 700-1800$ K, which suggests that the emission is from dust.

There are two uncertainties about the dust emission: where is it located and
how is it powered. Gerardy et al. (2002) and Fox et al. (2011) argue that the
IR emission is from pre-existing dust that is beyond the evaporation radius
and the forward shock front. In this case, a decline of the IR emission is
expected once the forward shock has overrun the dust shell, around $3-4$ years
after the explosion. Such a trend has been seen in some SNe IIn (Fox et al.,
2011). On the other hand, efficient dust formation has been considered to be a
viable scenario in interacting SNe, as the dust can form rapidly in the
extremely dense, post-shock cooling layers (Smith, 2014). Gall et al. (2014)
and Sarangi et al. (2018) propose that the dust in SN 2010jl formed in the
cold dense shell (CDS) downstream from the radiative shock wave. They find
that the dust can form at an age of about 1 year, when there is an increase in
the IR emission from the SN. At earlier times, the radiation field is too
strong to allow dust formation. In addition, dust may also form in the inner
SN ejecta. Signatures of the dust formation were reported in SN 2006jc (Smith
et al., 2008), SN 2005ip (Smith et al., 2009) and SN 2010bt (Elias-Rosa et
al., 2018). The main signatures of dust formation are an IR excess, a drop in
optical brightness via dust extinction, and asymmetry of the emission-line
profiles due to dust attenuation (Smith, 2014).

While the most straightforward source for powering the late IR emission in SNe
IIn would be continuing CS interaction, for SN 2010jl the evolution of the
X-ray absorption column implies a drop in the CSM density at an age of about 1
year (Chandra et al., 2015; Sarangi et al., 2018; Dwek et al., 2021). The
density drop would result in a drop in the power produced at the forward
shock. To alleviate this Sarangi et al. (2018) suggested that the observed
power is produced at the reverse shock wave, but the issue is unsettled.

SN 2017hcc (a.k.a ATLAS17lsn) was discovered on 2017 Oct. 2.38 UT (MJD
58028.38) by the Asteroid Terrestrial–impact Last Alert System (ATLAS, Tonry,
2011) and classified as Type IIn SN (Prieto et al., 2017). The SN reached a
peak at 13.7 mag in around $40-45$ days, indicating an absolute peak magnitude
of $-20.7$ mag at a distance $\sim 75$ Mpc (Prieto et al., 2017). The
supernova was not detected on 2017 September 30.4 at a limiting magnitude of
19.04111https://wis-tns.weizmann.ac.il/object/2017hcc/discovery-cert mag,
suggesting it was discovered soon after the explosion. We assume 2017 Oct 1 UT
(MJD 58027) as the date of explosion throughout this paper.

Prieto et al. (2017) obtained a bolometric light curve by using a blackbody
function to fit the SED using the data from ASAS-SN and the ultraviolet
telescope (UVOT) onboard Swift. They obtained a peak luminosity of $L_{\rm
bol,peak}=(1.34\pm 0.14)\times 10^{44}$ erg s-1, making SN 2017hcc one of the
most luminous SNe IIn ever (e.g., Smith et al., 2007; Ofek et al., 2014). They
constrain the peak risetime to be $\sim 27$ days and the total radiated energy
$\sim 4\times{10}^{50}$ erg. Early radio observations at 1.4 GHz with the
Giant Metrewave Radio Telescope resulted in a non-detection (Nayana & Chandra,
2017).

A very high degree of intrinsic polarization at optical wavelengths was
detected from SN 2017hcc ($>4.8$%) (Mauerhan et al., 2017; Kumar et al.,
2019). This is the strongest continuum polarization ever reported for a SN.
Kumar et al. (2019) noted a 3.5% change in the Stokes V polarization in 2
months suggesting a substantial variation in the degree of asymmetry in either
the ejecta and/or the surrounding medium of SN 2017hcc. They also estimated a
mass-loss rate of $\dot{M}=0.12$ $M_{\odot}$ yr-1 (for a wind speed of 20 km
s-1), which is comparable to the mass-loss rate for an LBV in eruption.
However, based on echelle spectra, Smith & Andrews (2020) suggested the CSM
wind to be flowing axi-symmetrically with wind speeds of $v_{w}=40-50$ km s-1
indicating bipolar geometry of the CSM created by losing 10 $M_{\odot}$ of
stellar mass to eruptive mass ejections. Smith & Andrews (2020) modeled the
optical and near IR emission of SN 2017hcc and found indications that the SN
ejecta were hidden behind the photosphere until day 75. They found that the
early time symmetric profiles changed to asymmetric blueshifted profiles at
late times. They suggested the asymmetry to be time and wavelength dependent,
causing a systematic blueshift in the line profiles. This was interpreted as a
signature of the formation of new dust in the SN ejecta as well as in the
post-shock gas within the CDS formed between the forward and the reverse
shocks.

We carried out early to late time observations of SN 2017hcc with the Chandra
and the X-ray telescope (XRT) onboard the Neil Gehrels Swift Observatory
(Swift-XRT) telescopes in X-ray bands, with the Karl G. Jansky Very Large
Array (VLA) in radio bands, and with Spitzer and the TripleSpec in IR bands.
In this paper we report multiwaveband observations of SN 2017hcc and their
interpretation. The observations are described in §2. Our main results are
summarised in §3, and the discussion and comparison with published data are in
§4. Finally, our main conclusions are summarised in §5.

Figure 1: Left panel: Chandra ACIS-S spectrum of SN 2017hcc. The spectrum is
fit with a thermal bremsstrahlung spectrum assuming a temperature of 3 keV.
Right panel: Contours of best fit column density with the best fit
normalization. The red, green and blue lines represent 1-$\sigma$, 2-$\sigma$
and 3-$\sigma$ contours.

##  2 Observations

###  2.1 Chandra Observations

We observed SN 2017hcc with the Chandra ACIS-S on 2019 Sep 27 (MJD 58753) for
an exposure of 40 ks. We extracted the spectrum, response and ancillary
matrices using Chandra Interactive Analysis of Observations software (CIAO;
Fruscione et al., 2006), using task specextractor. The CIAO version 4.6 along
with CALDB version 4.5.9 was used for this purpose. The
HEAsoft222http://heasarc.gsfc.nasa.gov/docs/software/lheasoft/ package Xspec
version 12.1 (Arnaud, 1996) was used to carry out the spectral analysis. Due
to low counts, only 5 channels were averaged and we used maximum likelihood
statistics for a Poisson distribution, i.e. the $c$-statistics (Cash, 1979).

The spectrum was fit with an absorbed thermal bremsstrahlung model. In the
view of Smith & Andrews (2020), the optical measurements implied a SN ejecta
velocity of 4000 km s-1 around the Chandra epoch and a CDS velocity of 1600 km
s-1 . Assuming the CDS velocity to be that of the forward shock, the reverse
shock velocity is 2400 km s-1 (Smith & Andrews, 2020). However, in SNe IIn the
CDS generally absorbs any X-ray emission coming from the reverse shock due to
the high column density of cool gas, and the dominant emission is likely to
come from the forward shock (Margutti et al., 2014; Ofek et al., 2014; Chandra
et al., 2015). The forward shock velocity of 1600 km s-1 estimated by Smith &
Andrews (2020) corresponds to a temperature of $\sim 3$ keV, which is what we
initially assume to fit the Chandra spectrum. Our observation resulted in a
detection with an unabsorbed $0.3-10$ keV luminosity of $1.94\times 10^{40}$
erg s-1 . The best fit column density is $1.13^{+0.78}_{-0.58}\times 10^{22}$
cm-2 (Fig. 1).

We, however, note that the forward shock velocity could be larger than the
above value, as discussed in §4.1. To reflect this possibility, we also ran
the fits by fixing the X-ray emitting shocked plasma temperature to be 20 keV
(corresponding to a forward shock velocity of 4000 km s-1 , an estimate based
on observations of SN 2010jl at an age of 737 days (Ofek et al., 2014; Chandra
et al., 2015)). In this case, our observation resulted in an unabsorbed
$0.3-10$ keV luminosity of $1.74\times 10^{40}$ erg s-1 . The best fit column
density is $(5.7^{+6.2}_{-4.5})\times 10^{21}$ cm-2 . While the X-ray
luminosity is lower by only 10% for 20 keV plasma as compared to that with 3
keV plasma, the best fit column density is lower by a factor of 2, though we
note that the uncertainty in column density is large with a 3-$\sigma$ upper
limit of $\sim 2\times 10^{22}$ cm-2 . Hence this is roughly consistent with
the previous value.

###  2.2 Swift-XRT Observations

We observed SN 2017hcc with the Swift-XRT starting 2017 Oct 28 (MJD 58054).
The observations continued until 2021 May 03 (MJD 59337) at various epochs.
The measurements were taken in the photon counting mode. The xselect program
of the HEAsoft package (version 6.28, CALDB version (XRT(20210915))) was used
to create the spectra and images. Observations closely spaced in time were
combined. None of the Swift-XRT observations resulted in detection with
3-$\sigma$ upper limits ranging around $\sim(1-4)\times 10^{-3}$ counts s-1.
The $0.3-10.0$ keV flux was obtained from the $3\sigma$ upper limits on the
count rates by assuming a thermal plasma of 3 keV and a column density of
$1.1\times 10^{22}$ cm-2 . The details are given in Table 1.

Table 1: X-ray observations of SN 2017hcc

Date of | MJD | Telescope | Exposure | Age | Count rate | Unabs. Flux | Unabs. Luminosity
---|---|---|---|---|---|---|---
Observation (UT) |  |  | (ks) | (day) | (cts s-1) | (erg s-1 cm-2) | (erg s-1)
$\sim$ 2017 Oct 30 | 58056 |  Swift-XRT | 13.47 | 29 | $<1.02\times 10^{-3}$ | $<0.54\times 10^{-13}$ | $<3.64\times 10^{40}$
$\sim$ 2017 Nov 17 | 58074 |  Swift-XRT | 5.20 | 47 | $<2.56\times 10^{-3}$ | $<1.35\times 10^{-13}$ | $<8.97\times 10^{40}$
$\sim$ 2017 Dec 10 | 58097 |  Swift-XRT | 9.58 | 70 | $<1.03\times 10^{-3}$ | $<0.55\times 10^{-13}$ | $<3.65\times 10^{40}$
2018 May 23 | 58261 |  Swift-XRT | 4.94 | 235 | $<3.85\times 10^{-3}$ | $<2.05\times 10^{-13}$ | $<13.7\times 10^{40}$
2019 Sep 27 | 58753 |  Chandra ACIS-S | 40.53 | 727 | $(1.28\pm 0.36)\times 10^{-3}$ | $(2.87\pm 0.53)\times 10^{-14}$ | $(1.94\pm 0.43)\times 10^{40}$
2021 May 03 | 59337 |  Swift-XRT | 4.71 | 1310 | $<2.81\times 10^{-3}$ | $<1.48\times 10^{-13}$ | $<10.0\times 10^{40}$

For Chandra observations, the temperature was fixed to 3 keV (see §4). The
best fit column density is $N_{H}=(1.13^{+0.78}_{-0.58})\times 10^{22}$ cm-1.
The count rates have been converted to fluxes assuming a thermal plasma of 3
keV and column density of $1.1\times 10^{22}$ cm-1 for the Swift-XRT
observations.

###  2.3 Spitzer observations

SN 2017hcc was observed with the Infrared Array Camera (IRAC) (Fazio et al.,
2004) onboard Spitzer (Werner et al., 2004) at three epochs (PI: Fox),
starting 2018 Oct 29 (MJD 58420; Table 2). We used the Spitzer Heritage
Archive (SHA)333https://sha.ipac.caltech.edu to download Post Basic Calibrated
Data (pbcd), which are already fully coadded and calibrated. Standard aperture
photometry was performed, although separate apertures were strategically
placed by eye to best estimate the local background. Figure 2 plots and Table
2 lists the mid-IR photometric data, which are also included in a
comprehensive Spitzer paper by Szalai et al. (2021).

As in the analysis laid out by Fox et al. (2011), the mid-IR photometry can be
fit as a function of the dust temperature, $T_{d}$, and mass, $M_{d}$. Given
only two photometry points per epoch, we assume a simple blackbody model,
consisting of a single graphite grain of size 1 $\mu$m and a single
temperature. We calculate opacities using the optical constants from Draine &
Lee (1984).Table 2 lists the results and Figure 2 shows the Spitzer data.

Figure 2: Spitzer photometry of SN 2017hcc. The data are taken at 3 epochs
(2018 Oct 29 (day 393), 2019 Apr 21 (day 568) and 2019 Nov 9 (day 770)) at 3.6
and 4.5 $\mu$m wavelengths (§2.3). The temporal evolution indices at 3.6$\mu$m
between days 393 and 568, and days 568 and 770 are $-0.43\pm 0.15$, and
$-1.14\pm 0.22$, respectively. The temporal evolution indices at 4.5$\mu$m
between days 393 and 568, and days 568 and 770 are $-0.17\pm 0.16$, and
$-0.76\pm 0.20$, respectively.  Table 2: Spitzer IRAC observations of SN
2017hcc

Quantity | Wavelength | Date of observations
---|---|---
|  | 2018 Oct 29.35 | 2019 Apr 20.58 | 2019 Nov 08.57
|  | (MJD 58420.35) | (MJD 58593.58) | (MJD 58795.57)
AOR id |  | 66120960 | 68799232 | 68799488
Epoch (d) |  | 393 | 568 | 770
Apparent magnitude | 3.6$\mu$m | 12.97$\pm$0.04 | 13.16$\pm$0.05 | 13.52$\pm$0.05
| 4.5$\mu$m | 12.49$\pm$0.04 | 12.56$\pm$0.04 | 12.81$\pm$0.04
Absolute magnitude | 3.6$\mu$m | -21.32$\pm$0.19 | -21.14$\pm$0.19 | -20.77$\pm$0.19
| 4.5$\mu$m | -21.80$\pm$0.19 | -21.73$\pm$0.19 | -21.48$\pm$0.19
Flux density ($\mu$Jy) | 3.6$\mu$m | $1818.72\pm 70.35$ | $1533.79\pm 64.98$ | $1097.13\pm 55.58$
| 4.5$\mu$m | $1813.78\pm 71.83$ | $1703.51\pm 69.03$ | $1354.05\pm 62.36$
$\nu L_{\nu}$ (erg s-1) | 3.6$\mu$m | $(10.42\pm 0.40)\times 10^{41}$ | $(8.90\pm 0.37)\times 10^{41}$ | $(6.29\pm 0.32)\times 10^{41}$
| 4.5$\mu$m | $(8.28\pm 0.33)\times 10^{41}$ | $(7.78\pm 0.32)\times 10^{41}$ | $(6.18\pm 0.28)\times 10^{41}$
Dust massa ($M_{\odot}$) |  | $2.26\times 10^{-3}$ | $3.60\times 10^{-3}$ | $4.57\times 10^{-3}$
Dust Temp (K) |  | $1.28\times 10^{3}$ | $1..07\times 10^{3}$ | $0.93\times 10^{3}$
$L_{\rm bol}$b (erg s-1) |  | $4.80\times 10^{41}$ | $3.73\times 10^{41}$ | $2.70\times 10^{41}$
Blackbody radius (cm) |  | $1.57\times 10^{16}$ | $1.97\times 10^{16}$ | $2.23\times 10^{16}$

a The dust mass estimates are sensitive to the chosen grain parameters. The
listed values are estimated assuming 1 $\mu$m grain size of graphite
composition.

b Bolometric luminosity assuming a single temperature dust blackbody
radiation.

Figure 3:  Triplespec near-IR spectra of SN 2017hcc on 2017 Dec 05 (day 66),
2018 Jan 06 (day 98), 2018 Sep 18 (day 353), 2018 Nov 02 (day 398) and 2019
Jan 17 (day 474). Note that the early spectra are brighter at lower
wavelengths, whereas the late time spectra are brighter at longer wavelengths.
To show this trend better, the top four spectra are offset by $10^{-15}$,
$6\times 10^{-16}$, $5\times 10^{-16}$, and $3\times 10^{-16}$ erg s-1 cm-2
A-1 arcsec-2, respectively.

###  2.4 TripleSpec Observations

We obtained five epochs, 2017 Dec 05 (MJD 58092), 2018 Jan 6 (MJD 58124), Sep
18 (MJD 58379), Nov 02 (MJD 58424) and 2019 Jan 17 (MJD 58500), of near-IR
spectroscopy, with simultaneous continuous wavelength coverage from 0.95-2.46
$\mu$m in five spectral orders, with the TripleSpec spectrograph at the Apache
Point Observatory 3.5-m (Wilson et al., 2004). Our observation sequence
consists of 300-s dithered exposures that could be pair subtracted to allow
for correction of thermal background and night-sky emission lines. We used a
modified version of SpexTool for the spectral extraction (Cushing, 2004).
While the early data are dominated by shorter wavelength emission, the later
data reveal higher flux densities at longer wavelengths (Fig. 3). However, we
note a caveat here that due to calibration uncertainty the TripleSpec flux
should not be considered absolute and ideally should be scaled to near-IR
photometry.

###  2.5 VLA observations

The VLA observed SN 2017hcc between 2020 June 29–30 (MJD 59029–59030) in bands
X (8–12 GHz) and K (18–26 GHz), and later in June 2021 in Ka (26–40 GHz), Ku
(12–18 GHz), X, C (4–8 GHz) and S (2–4 GHz) bands. The data were analysed
using the standard packages within the Common Astronomy Software Applications
package (CASA, McMullin et al., 2007). The details of the observations and the
flux density values can be found in Table 3. We add 10% error in the
quadrature to the flux density values for analysis purposes, a typical
uncertainty in the flux density calibration scale at the observed frequencies
444https://science.nrao.edu/facilities/vla/docs/manuals/oss/performance/fdscale.

Table 3: Radio observations of SN 2017hcc Date of | MJD | Telescope | Age | Central | Flux density | rms | Luminosity | $\nu L_{\nu}$
---|---|---|---|---|---|---|---|---
Obs. (UT) |  |  | (days) | Freq (GHz) |  $\mu$Jy |  $\mu$Jy | (erg s-1 Hz-1) | (erg s-1)
2020 Jun 29 | 59029 | VLA | 1002 | 10.0 | 32.5$\pm$7.6 | 6.7 | $(2.19\pm 0.51)\times 10^{26}$ | $(2.19\pm 0.51)\times 10^{36}$
2020 Jun 30 | 59030 | VLA | 1003 | 22.0 | 77.1$\pm$16.2 | 15.5 | $(5.19\pm 1.09)\times 10^{26}$ | $(11.42\pm 2.40)\times 10^{36}$
2021 Jun 04 | 59369 | VLA | 1342 | 10.0 | $<45.6$ | 15.26 | $<3.07\times 10^{26}$ | $<3.07\times 10^{36}$
2021 Jun 09 | 59374 | VLA | 1347 | 15.4 | 31.0$\pm$11.3 | 10.5 | $(2.09\pm 0.76)\times 10^{26}$ | $(3.21\pm 1.17)\times 10^{36}$
2021 Jun 09 | 59374 | VLA | 1347 | 6.1 | $<26.7$ | 8.9 | $<1.80\times 10^{26}$ | $<1.10\times 10^{36}$
2021 Jun 11 | 59376 | VLA | 1349 | 3.0 | $<182.4$ | 60.7 | $<12.28\times 10^{26}$ | $<3.68\times 10^{36}$
2021 Jun 14 | 59379 | VLA | 1352 | 33.1 | 33.5$\pm$11.8 | 11.6 | $(2.26\pm 0.79)\times 10^{26}$ | $(7.46\pm 2.63)\times 10^{36}$
2021 Jun 16 | 59381 | VLA | 1354 | 1.5 | $<228$ | 76.1 | $<15.34\times 10^{26}$ | $<2.30\times 10^{36}$

##  3 Analysis and Results

###  3.1 X-ray analysis and the column density

Swift-XRT observations of SN 2017hcc started when the SN was around a month
old; however, the first X-ray detection of SN 2017hcc was at an age of $\sim
2$ years with the Chandra telescope. Our Chandra observations did not have
enough counts to determine the shock temperature and the column density
separately. We assumed a temperature of $\sim 3$ keV, corresponding to a CDS
velocity of 1600 km s-1 , as advocated by Smith & Andrews (2020). The best fit
column density is $1.13^{+0.78}_{-0.58}\times 10^{22}$ cm-2 (Fig. 1). The
Milky Way line of sight reddening is $E(B-V)=0.0285$ and that through the host
galaxy is $E(B-V)=0.016$ (Smith & Andrews, 2020). These values correspond to
their respective column densities of $N_{H,\rm MW}=1.64\times 10^{20}$ cm-2
and $N_{H,\rm host}=0.92\times 10^{20}$ cm-2 , using Schlafly & Finkbeiner
(2011) recalibration of the Schlegel et al. (1998) extinction map. To derive
$N_{H}$ we use $N_{H}=E(B-V)\times 5.8\times 10^{21}$ cm-2 . However, this
relation is an empirical relation and has uncertainties (e.g. Gudennavar et
al., 2012; Liszt, 2019). Combining the column density due to the Milky Way and
the host galaxy gives $2.56\times 10^{20}$ cm-2 , which has negligible effect
on the best fit column density which is two orders of magnitude higher. The
excess column density comes from the CSM, i.e. $N_{H,\rm CSM}=\sim 1.10\times
10^{22}$ cm-2 .

The $0.3-10.0$ keV unabsorbed flux is $f_{0.3-10\,\rm keV}=(2.87\pm
0.53)\times 10^{-14}$ erg cm-2 s-1, corresponding to a luminosity of
$L_{0.3-10\,\rm keV}=(1.94\pm 0.43)\times 10^{40}$ erg s-1. This value does
not change significantly even if we use an X-ray emitting plasma temperature
of 20 keV based on the temperature observed in SN 2010jl (Ofek et al., 2014;
Chandra et al., 2015). In Fig. 4, we plot the SN 2017hcc X-ray luminosity
along with other well-observed SNe IIn. Other than SN 1978K and SN 1998S,
detected SNe IIn are brighter than SN 2017hcc. The SN 2017hcc luminosity is
comparable to that of SN 1998S at the same age. We discuss the possible
reasons for low X-ray luminosity in §4.

Figure 4: The 0.3–8 keV luminosities of X-ray detected SNe IIn. The Chandra SN
2017hcc detection is marked with a big magenta circle and the Swift upper
limits with inverted triangles. Reproduced from Chandra (2018).

###  3.2 IR analysis

Figure 5: The spectra on day $\sim 395$ combined with TripleSpec and Spitzer
data. The light purple color shows original spectrum, whereas the purple
spectrum has been shifted upwards to match the blackbody curve passing through
the 3.6 $\mu$m Spitzer data point.

SN 2017hcc is extremely bright at Spitzer IRAC wavelengths. The 4.6 $\mu$m
absolute magnitude reached $-21.7$ mag after a year, which corresponds to a
peak IR luminosity $\sim 10^{42}$ erg s-1. The late time emission is dominated
by the IR emission (Smith & Andrews, 2020), as is frequently the case for SNe
IIn (see §1). We estimate the temporal evolution $\beta$ (where $\beta$ is
defined as $F_{\nu}(t)\propto t^{\beta}$) and between Spitzer flux densities
($F_{\nu}(t)$) in $\mu$Jy (Fig. 2). The temporal evolution between epochs 1
and 2, at 3.6 and 4.5 $\mu$m, are $\beta=-0.46\pm 0.16$ and $\beta=-0.17\pm
0.15$, respectively. The temporal indicies between the epochs 2 and 3 at the
two frequencies are $\beta=-1.10\pm 0.21$ and $\beta=-0.76\pm 0.20$. The
spectral indices $\alpha$ (defined as $F_{\nu}\propto\nu^{\alpha}$) between
the two wavelengths at each epoch are $\alpha=-0.01\pm 0.25$, $\alpha=0.47\pm
0.26$ and $\alpha=0.94\pm 0.31$, respectively. The slow decline of the IR
emission, combined with the fast optical decline (Kumar et al., 2019; Smith &
Andrews, 2020), has important implications for the origin of the dust (§4.5).

The TripleSpec data at five epochs: 2017 Dec 5 (d 65), 2018 Jan 6 (d 97), Sep
18 (d 352), Nov 02 (d 397) and 2019 Jan 17 (d 473) are plotted in Fig. 3.
While the early data are dominated by shorter wavelength emission, the later
data reveal higher flux densities at longer wavelengths. A growing IR excess
over the photospheric emission develops over time. We combine the measurements
of Spitzer and TripleSpec on 2018 Oct 29 and 2018 Nov 2, respectively, and fit
a simple blackbody model to the combined spectrum. The IR spectra using these
data are plotted in Fig. 5. However, the calibration of the TripleSpec data
has large uncertainties and hence the flux should not be considered absolute.
Ideally it should be scaled to near-IR photometry. We have varied the scaling
of the TripleSpec spectrum manually, and tried to shift to the blackbody curve
passing through the Spitzer 3.6 $\mu$m data point. While the blackbody curve
with $T\approx 1600$ K fits the data reasonably well, there seems to be a
slight excess of the 4.5 $\mu$m flux. However, the significance of the fit is
low due to the calibration uncertainties and the fact that the scaled
TripleSpec spectrum does not seem to go through both the Spitzer data points
and thats causes the uncertainty in the scaling factor. Hence we do not use
the parameters derived from these fits.

Figure 6: Comparison of near IR spectra of SN 2017hcc (Borish et al., 2015)
with those of SN 2010jl. Spectra are normalized to the continuum level. The SN
2017hcc spectrum is on day 353 (brown lines, corresponding to Triplesec data
on 2018 Sep 18 (MJD 58379)). SN 2010jn spectra are on day 219 (light blue
lines) and day 403 (dark blue lines) Top panel: The full spectra. Bottom two
panels show zoomed plots for O I, Pa $\beta$, Paschen $\gamma$ and He I lines.

We compare the SN 2017hcc near IR TripleSpec spectra with those of SN 2010jl
in Figure 6. The two SNe have qualitatively similar characteristics but there
are significant differences. In SN 2017hcc, the O I (1.129 $\mu$m) line is
much stronger relative to Pa$\beta$ (1.282 $\mu$m) than in SN 20201jl, while
the He I (1.083 $\mu$m) line is weaker relative to Pa$\beta$. The H lines show
some asymmetry towards the blue side in SN 2017hcc, This is seen visually as a
greater area under the lines to the blue side of the central wavelength vs the
red side. But the asymmetry is less than that in SN 2010jl. This could be due
to dust formation in the post-shock shell and in the SN ejecta, which
progressively obscures the redshifted part. Any associated CS dust was
presumably evaporated by the SN. Narrow lines like those observed in SN 2010jl
are present in SN 2017hcc (Smith & Andrews, 2020), but are not resolved here.
The width of the expected narrow lines is $\sim 50$ km s-1 , whereas the
TripleSpec resolution is $\sim 100$ km s-1 . The estimated wind velocity is
100 km s-1 for SN 2010jl (Fransson et al., 2014), but is $40-50$ km s-1 for SN
2017hcc (Smith & Andrews, 2020)

###  3.3 Radio Analysis

Figure 7: Radio spectra of SN 2017hcc with the VLA at two epochs. The spectra
on day 1000 is in blue color and the one on day 1347 is in red color. The
triangles represent upper limits. A simple FFA model is fit to the data.

The VLA measurements show faint detections ($\leq 5\,\sigma$, as mentioned in
Table 3 and Fig. 7). The peak spectral luminosity is $5.2\times 10^{26}$ erg
s-1 Hz-1, which is towards the lower end of the SNe IIn luminosities (Chandra,
2018). The day 1000 spectral index between these frequencies is $+1.1$. This
indicates that the supernova is still moderately optically thick. Such
behavior has been observed in the early evolution of other SNe IIn such as SN
1986J which showed $\alpha\sim+1.5$ in its rise to maximum (Weiler et al.,
1990) and has been attributed to clumpiness in the emitting region (Weiler et
al., 1990). Our radio data taken a year later show a higher radio luminosity.
The spectral indices between the detected points are flat, indicating that the
radio emission is near its transition from optically thick to thin regime. Due
to the small number of data points, we fit the simple free-free absorption
(FFA) model used in Chandra et al. (2012) assuming an optically thin spectral
index of $-0.6$. The best fits at the two epochs are shown in Fig. 7. We also
attempt to fit the synchrotron self-absorption model (SSA) and from the peak
we estimate the SSA velocity using the formulation in Chevalier (1998). This
velocity is around $\sim 150$ km s-1 , which is an order of magnitude smaller
than the ejecta velocity. This rules out SSA as a significant absorption
mechanism.

##  4 Discussion

###  4.1 Shock breakout in presupernova mass-loss

With a peak bolometric luminosity of $1.3\times 10^{44}$ erg s-1 reached in
$\sim 30$ days and an absolute magnitude of $-20.7$ mag (Prieto et al., 2017),
SN 2017hcc is at the boundary of superluminous supernovae with a two orders of
magnitude larger luminosity than normal core-collapse supernovae. In addition,
strong narrow lines are present (Smith & Andrews, 2020), indicating that dense
outflowing circumstellar gas surrounds the supernova. If this gas is
sufficiently optically thick, the shock breakout radiation from the supernova
will occur after the interaction shell has moved to the mass-loss region.
Quantitatively, if the wind optical depth is $\tau_{\rm w}$ and $v_{\rm sh}$
is the shock speed, then for cases where $\tau_{\rm w}>c/v_{\rm sh}$, the SN
shock breaks out in the wind. Such a breakout through the thick wind may
extend long enough, lasting days to months, and may account for the peak SN
luminosity of SN 2017hcc (Chevalier & Irwin, 2011).

We use the approximate formulation of Chevalier & Irwin (2011) for shock
breakout in a wind. An initial assumption for this is that the mass-loss
density profile is that for a steady wind, $\rho=Dr^{-2}$, where
$D=\dot{M}/(4\pi R^{2}v_{\rm w})$. As per their formulation, there are two
length scales that are important. One, $R_{\rm d}$, is the radius at which the
diffusion time equals the expansion time, so photons can move out through the
wind. We have $R_{\rm d}=\kappa Dv_{\rm sh}/c$, where $\kappa$ is the opacity.
The other important length scale is $R_{\rm w}$ where the wind density cuts
off. The character of the shock breakout is expected to depend on whether
$R_{\rm w}>R_{\rm d}$, when the diffusion is important, or vice versa for
which the diffusion time scales are large (Chevalier & Irwin, 2011). The
indications are that SN 2017hcc is in the case $R_{\rm w}>R_{\rm d}$. This is
indicated by the presence of bright infrared emission to late times (2 years)
implying that the mass loss region extends to large radius. Also, the narrow
line emission in optical bands (Smith & Andrews, 2020) shows a dense mass loss
region. We scale the parameter $D$ to $D_{*}$ such that $D_{*}=D/(5\times
10^{16}$ g cm-1) corresponding to a mass loss rate of $0.1M_{\odot}$ yr-1 and
wind velocity of 100 km s-1. Here $D$ is in g cm-1 and $D_{*}$ is a
dimensionless quantity. As mentioned earlier, in SN 2017hcc the rise time can
be approximated by the rise time of the bolometric light curve, 30 days
(Prieto et al., 2017), so $D_{*}\approx 4.5$. For the wind velocity of 45 km
s-1, we get $\dot{M}=0.1M_{\odot}$ yr-1 which should apply near the time of
maximum light.

At shock breakout, a radiation dominated shock can no longer be maintained and
a cold dense shell (CDS) forms at the shock interface. A simple model for a
supernova density profile, used by Chevalier & Irwin (2011), is an inner
region with a flat density surrounded by a steep $\rho\propto r^{-7}$ profile.
For typical parameters, the interaction region between the ejecta and the CSM
is in the steep power law region. For normal supernova parameters, Equation
(6) in Chevalier & Irwin (2011) shows that the breakout luminosity can be
close to that observed in SN 2017hcc near bolometric peak.

In the above we use 4000 km s-1 for the shock velocity, as indicated in the
X-ray observations of SN 2010jl (Chandra et al., 2015) and SN 2014C (Margutti
et al., 2017). Here we discuss the shock velocity. Smith & Andrews (2020)
determine $v_{\rm sh}$ from the emission line profiles of H$\alpha$ and
H$\beta$. They find that the line profiles can be fit by 2 Gaussian
components, one with a FWHM of 4000 km s-1 and one with 1100 km s-1, which
they identify with emission from the freely expanding ejecta and the CDS,
respectively. After allowance for a geometrical factor, Smith & Andrews (2020)
use 1600 km s-1 for $v_{\rm CDS}$, where $v_{\rm CDS}=v_{\rm sh}$ for a thin
shell. However, there is no clear physical reason for the emission from the
CDS or ejecta to have a Gaussian profile with a width corresponding to a
physical parameter. Dessart et al. (2015) modeled the emission from SN 2010jl
and found reasonable estimates for the line profile of H$\alpha$ considering
emission from the CDS with electron scattering and CDS velocity $\sim 3000$ km
s-1. Also, in a spherically symmetric interaction, a large, unexpected
deceleration would be required between the ejecta and the CDS for the values
mentioned by Smith & Andrews (2020). In the formation of a thin shell between
ejecta with power law index $n$ and CSM with power law $s$, we have $R\propto
t^{(n-3)/(n-s)}$ (Chevalier, 1982) so that $v_{\rm ej}/v_{\rm
CDS}=(n-s)/(n-3)$. For typical values $n=8$ and $s=2$ (e.g., Fransson et al.,
2014, for SN 2010jl), we have $v_{\rm ej}/v_{\rm CDS}=1.2$, as opposed to
$v_{\rm ej}/v_{\rm CDS}=2.50$ from the line observations (Smith & Andrews,
2020). The low value of $v_{\rm CDS}$ is not consistent with the type of
evolution seen in SN 2010jl during the first year if there is spherically
symmetric expansion. Other estimates for $v_{\rm CDS}$ can be considered. In
the case of SN 2010jl, an X-ray observation at 2 yr with NuSTAR showed a
temperature of 20 keV (Ofek et al., 2014; Chandra et al., 2015), which
corresponds to a shock velocity of 4,000 km s-1. A higher velocity would be
expected at earlier times and there is possible evidence for a higher velocity
from a feature that appears in the He I line in SN 2010jl (Borish et al.,
2015; Chandra et al., 2015). We consider both small and large values for
$v_{\rm CDS}$.

For the shock front to break out in the wind region, $\tau_{\rm w}>c/v_{\rm
sh}$ is needed. Taking $R_{\rm b}$ to be the radius of the base of the wind,
we have $\tau_{w}\approx 5\times 10^{16}kR_{\rm b}^{-1}D_{*}$, where $k$ is
the opacity $\kappa$ in units of 0.34 cm2g-1 (Chevalier & Irwin, 2011). As
discussed above, for $D_{*}=4.5$, this gives $\tau_{w}\approx 76(R_{\rm
b}/{10^{15}{\rm cm}})^{-1}$. Thus for reasonable parameters, the shock
breakout does occur in the wind.

One aspect of shock breakout is that the temperature should increase with
rising luminosity, since the initial rise to maximum luminosity should be
primarily due to heating of the photosphere. This has not been seen in SN
2017hcc (Kumar et al., 2019). However, we note that the early measurements for
temperature are not available before the peak in the light curve.

###  4.2 Post-Breakout Interaction

Once the diffusion time is less than the supernova age, the luminosity
corresponds to the power generated at the shock. We have a luminosity $L=4\pi
R^{2}\epsilon\frac{1}{2}\rho v_{\rm sh}^{3}$, where $R$ is the radius of the
CDS and $\epsilon$ is an efficiency factor for the conversion of shock power
to radiation (Chugai, 1990). At early times, the cooling time is less than the
age so the shock power is radiated and $\epsilon\approx 1$. With sufficient
good quality data, the winds of SNe IIn are generally found to be not steady
(Dwarkadas, 2007), but the assumption should lead to a rough estimate for
$\dot{M}$, which, for a steady wind is:

$\dot{M}=2Lv_{\rm w}/\epsilon v_{\rm sh}^{3}$ (1)

The value of $v_{\rm w}$ can be estimated from high spectral resolution
observations. Smith & Andrews (2020) find $v_{\rm w}\approx(40-50)$ km s-1. To
estimate the value of $\dot{M}$, we use the bolometric luminosity of $\sim
10^{42}$ erg s-1 which is the luminosity of infrared dust emission at $1-2$ yr
(Table 2). We now find from Equation 1 $\dot{M}=2\times
10^{-3}\epsilon^{-1}L_{42}(v_{\rm sh}$/ 4000 km s${}^{-1})^{-3}$ $M_{\odot}$
yr-1, where $v_{w}=45$ km s-1 is assumed, and $L_{42}$ is the bolometric
luminosity in the units of $10^{42}$ erg s-1. So we now have $D_{*}\approx
0.024$. We note that the value of $v_{\rm w}$ is determined over the first 100
days (Smith & Andrews, 2020); there could be a change at later times. In
addition, the value of the mass loss is sensitive to the shock velocity.
However, our results indicate that the mass loss density drops more rapidly
with radius than in the steady ($r^{-2}$) case.

The cooling of the postshock gas is important for the value of $\epsilon$ in
the above formula. We now discuss the validity of cooling. Our discussion is
similar to that of Chevalier & Fransson (2017). For electron temperature
$T_{\rm e}\gtrsim 2.6\times 10^{7}$ K, which is expected for the post-forward
shock wave, free-free cooling dominates with a cooling rate $\Lambda=1.0\times
10^{-23}(T_{e}/10^{7}\,{\rm K})^{0.5}$ erg s-1 cm3. Assuming isobaric cooling
behind the shock, the cooling time is $t_{\rm
cool}=5kT_{e}/n_{H}\Lambda(T_{e})$. Using Chevalier & Fransson (2017) eqn.
(18), the transition time at which the age equals the cooling time, for ejecta
velocity 4000 km s-1 (indicated with $v_{\rm ej,4}$ in the formula below) is
(for $n=8$)

$t_{\rm tr}=342D_{*}v_{\rm ej,4}^{-3}{\rm days}.$ (2)

which is 75 days for our preferred parameters. This estimate would imply that
the supernova was past the cooling phase at the time of the X-ray observation.
However, the uncertainties are large so some effect of cooling cannot be ruled
out. The estimates should be taken with caution, since the properties of the
absorbing gas are uncertain because of early Compton heating and radiative
cooling as well as recombination (Lundqvist & Fransson, 1988).

###  4.3 Variable mass-loss rate

The above treatment shows that the mass-loss rate obtained at $\sim 30$ days
from the shock breakout was $\sim 0.1$ $M_{\odot}$ yr-1, whereas at few
hundred days, obtained from the IR measurements, was $\sim 2\times 10^{-3}$
$M_{\odot}$ yr-1. In addition, we can also use the timing of the peak radio
flux density to estimate the mass-loss rate at $\sim 1000$ d. The mass-loss
rate derived from our radio measurement assuming a CSM temperature of $10^{4}$
K is $\sim 6.5\times 10^{-4}$ $M_{\odot}$ yr-1. These epochs correspond to
roughly $\sim 10$, $\sim$100 and $\sim$300 days before explosion and the mass
loss rate changes roughly two to three orders of magnitude. This indicates a
variable mass-loss rate, with enhanced mass-loss event occurring a decade
before the supernova explosion.

There is growing number of evidence that many core-collapse supernovae undergo
similar enhancements during the last years before the explosion. Maeda et al.
(2021) reported evidence of enhanced mass-loss in nearby type Ic supernova SN
2020oi in their Atacama Large Millimeter Array (ALMA) observations, which they
interpreted as the pre-SN activity likely driven by the nuclear burning
activities in the star’s final moments. Similar conclusions were drawn for the
famous supernova SN 2014C undergoing metamorphosis from a stripped envelope
supernova to a Type IIn via their radio and X-ray observations (Anderson et
al., 2017; Margutti et al., 2017; Brethauer et al., 2022). The flash
spectroscopy of various supernovae have also revealed the similar conclusions.
A remarkable example is SN 2013fs, discovered mere 3 hrs after the explosion
by the automated real-time discovery and classification pipeline of the
intermediate Palomar Transient Factory (iPTF) survey, which enabled
spectroscopy measurements within 6 hours of discovery (Yaron et al., 2017).
The early measurements revealed a dense CSM to be confined within $10^{15}$ cm
indicating enhanced mass-loss ($10^{-3}$ $M_{\odot}$ yr-1) in the last one
year before explosion in an otherwise normal core-collpase supernova. This led
authors to suggest that such pre-SN activities may be common in exploding
stars. Shivvers et al. (2015) examined the Keck HIRES spectrum of SN 1998S
taken within a few days after the SN and found convincing evidence for
enhanced mass-loss rate in the last 15 years of star’s life. These
measurements combined with other published measurements indicated much smaller
mass-loss rate at earlier times, arriving at the conclusion of this section.
SN 2020tlf was a normal Type IIP/IIL supernova, where flash spectroscopy soon
after the discovery revealed evidence of photoionization of CSM confined
within $10^{15}$ cm created at a mass-loss rate of $10^{-2}$ $M_{\odot}$ yr-1
in a 10–12 $M_{\odot}$ red supergiant progenitor, whereas 3 orders of
magnitude smaller mass-loss rate at larger distances (Jacobson-Galán et al.,
2022). Similarly Terreran et al. (2022) also found the presence of dense and
confined CSM indicating a phase of enhanced mass-loss in its final moments of
the progenitor of SN 2020pni. Tartaglia et al. (2021) presented the very early
phase to nebular phase observations of Type II supernova SN 2017ahn,
discovered just a day after explosion. Their modeling indicated a complex CSM
surrounding the progenitor star with evidence of enhanced mass-loss in the
last moments of likely massive yellow supergiant progenitor. While for SNe
IIn, LBV scenario has been favored for enhanced mass-loss rates just before
the explosion, above examples favor the idea that many core-collapse
supernovae of different flavors may experience enhanced mass-loss in their
final years, which could be governed by somewhat common physical mechanisms
(Terreran et al., 2022).

###  4.4 Low X-ray luminosity

The lack of bright X-rays in SN 2017hcc is surprising, as the detected SNe IIn
are generally brighter than other core collapse supernovae (Chandra, 2018). A
dense interaction is indicated in SN 2017hcc at early epochs by the optical
luminosity and electron scattering line profiles. X-ray emission is detected
at day 700 but the luminosity of SN 2017hcc is lower than that typically found
for strongly interacting supernovae of Type IIn and is closer to the
luminosity observed from SN 1998S (Fig. 4). In SN 2017hcc during the time of
the Chandra observation, most of the shock power goes into the IR, which is
$\sim 10^{3}$ times the X-ray emission at the same epoch. SN 1998S had late IR
emission, but at a level $\sim 10^{40}$ erg s-1 (Pozzo et al., 2005) at 2
years, which is nearly two orders of magnitude smaller than that of SN
2017hcc; the lower luminosity was consistent with the lower mass-loss rate of
$2\times 10^{-4}M_{\odot}$ yr-1 deduced from the radio and X-ray emission
(Pooley et al., 2002).

When the interaction is strong, the forward shock front is initially radiative
and evolves to a non-radiative phase (Chevalier & Irwin, 2012). In the cooling
phase, the CDS created by the interaction is subject to the nonlinear thin
shell instability (NTSI), giving rise to a corrugated structure (Vishniac,
1994). Numerical simulations of the instability show that the X-ray emission
is suppressed (Steinberg & Metzger, 2018; Kee et al., 2014). The reason for
X-ray suppression is the CDS that forms. The hot post-shock gas cools
radiatively towards lower temperature and pressure in the presence of the CDS.
Much of the cool emitting gas is underpressured relative to the hot
surrounding medium. This pressure difference between the hot shock and the
cool dense shell robs the hot gas of its thermal energy, which is now radiated
with high efficiency at lower temperature. Another factor is that the
corrugation of the forward shock gives oblique shocks that are weaker than
head on shocks. However, as discussed above, estimates of the cooling at the
time of the X-ray detection indicates that the shockwave is probably not
cooling. In this case, the NTSI is not a factor, although the uncertainties
and assumptions may allow the cooling scenario.

###  4.5 Late time excess of IR emission and origin of the dust

The Spitzer IR photometric data corresponds to peak IR luminosities reaching
$\sim 10^{42}$ erg s-1. The origin of the bright IR emission at late times is
not clear. The IR evolution is rather flat, as opposed to a faster decline
seen in the optical band (Kumar et al., 2019), though at earlier times.
Usually the IR excess imply a rising relative contribution of the dust
component to the bolometric luminosity as the optical light-curve fades. The
late-time high IR luminosity and the systematic blue-shift in the line
profiles (Smith & Andrews, 2020) suggest that SN 2017hcc likely has
contributions from two components: 1) dust formation in the ejecta and the
CDS, and 2) pre-existing dust beyond the evaporation radius, giving rise to IR
emission.

Smith & Andrews (2020) find that SN 2017hcc line profiles show a progressively
increasing blueshift. They reject the possibility of this feature being due to
occultation by the SN photosphere, pre-shock acceleration of the CSM, or
asymmetric CSM, and explain it to be arising due to absorption by newly formed
dust in the post-shock shell and then SN ejecta.

Based on our IR fits, we can calculate the blackbody radius, $r_{\rm bb}$,
which is the minimum size of the dust emitting component, using the modified
blackbody expression given in Fox et al. (2010). Using the best fit values in
Table 2, the bolometric luminosities assuming blackbody emission are
$4.79\times 10^{41}$ erg s-1, $3.73\times 10^{41}$ erg s-1 and $2.70\times
10^{41}$ erg s-1 at epochs 393, 568 and 770 days, respectively. These
luminosities imply that the values of $r_{\rm bb}$ are $1.6\times 10^{16}$ cm,
$2\times 10^{16}$ cm and $2.2\times 10^{16}$ cm at these epochs, respectively.
The shock radii at these epochs are $(0.5-1.3)\times 10^{16}$,
$(0.8-2.0)\times 10^{15}$ and $(1.1-2.7)\times 10^{16}$ cm, respectively for
shock velocity ranging between $1600-4000$ km s-1 . We note that these values
have significant uncertainties arising from assumption of single graphite
grain of size 1 $\mu$m and a single temperature, as well as due to
uncertainties in the opacity measurements. These findings give interesting
clues to the dust origin. The shock can typically destroy any pre-existing
dust within the shock radius. The shock radii are smaller (or equal for higher
ejecta velocity) than the blackbody radii at each epoch. This implies that
there could still be pre-existing dust that is beyond the outer shock front
that can give the IR emission (Fox et al., 2011), in addition to newly formed
dust and can cause IR excess.

The dust luminosity of SNe IIn in general is higher than in other SNe,
suggesting that the dense CSM is important for dust formation in strongly
interacting SNe IIn (Tinyanont et al., 2016). However, most of the SNe IIn
with a high IR flux were also associated with a high X-ray luminosity. This is
not the case with SN 2017hcc. Fig. 4 shows the SN 2017hcc X-ray luminosity
with other well observed SNe IIn. Other than SN 1978K and SN 1998S, the
detected SNe IIn are brighter than SN 2017hcc. The luminosity of SN 2017hcc is
comparable to that of SN 1998S at the same age, which is known to be due to
lower CS interaction. This is unlikely to be the case for SN 2017hcc.

###  4.6 Asymmetry in the CSM

Asymmetry has been seen in many SNe IIn. For example, the large X-ray
luminosity of SN 2006jd implied a high density, but a small amount of
photoabsorption in the spectrum implied a low density (Chandra et al., 2012).
These conflicting measurements suggested an asymmetric CSM. This seems to be a
general feature of interacting CSM, and may indirectly support the luminous
blue variable (LBV) progenitor scenario (Smith, 2014).

Direct evidence for asymmetry in SN 2017hcc comes from measurements of high
polarization of the photospheric emission (Mauerhan et al., 2017; Kumar et
al., 2019). The fact that the polarized fraction declined with time indicates
that the asymmetry was in the CSM, not the supernova ejecta.

Also, Smith & Andrews (2020) found, based on blueshifted line profiles, that
the intrinsic emission-line profile from the fast SN ejecta is symmetric,
suggesting that the underlying SN explosion itself was not highly aspherical,
and the polarization is likely related to the CSM. Another line of argument
regarding the highly asymmetric CSM in SN 2017hcc suggests that the
polarization measurements are close to maximum. The emission at maximum light
in SNe IIn is dominated by interaction of the ejecta with the CSM, implying
that the major source of intrinsic polarization at maximum is due to this
interaction, particularly with an asymmetric CSM. The luminosity from the
interaction becomes progressively weaker as the CSM around the SN evolves into
an optically thin state, resulting in a decrease in polarization. In addition,
Smith & Andrews (2020) find that even though optical depths of the CSM are
quite high, the SN ejecta emerge by day 75, which requires a non-spherical
geometry of the CSM.

Smith & Andrews (2020) noted that the line profiles during the first $\sim
100$ days are remarkably symmetric, indicating that the CSM speed remains
similar over a large range of radii regardless of the changing supernova
luminosity. The observed asymmetry at later times is mild and limited to low
radial velocities. Based on these two arguments, they suggested that the
asymmetry is in the form of axisymmetry and our line of sight is close to the
polar region, which when viewed from the earth at an intermediate latitude
will show asymmetry.

As the main power source for SNe IIn is the kinetic energy of the SN ejecta, a
large fraction of the ejecta kinetic energy can be converted into radiation
due to the high density CSM. However, an axisymmetric CSM could lead to an
inefficient conversion of kinetic energy into radiation in some directions,
causing a lower X-ray flux. A low column density can be explained by the
axisymmetry. Our radio measurements may also support this view. The radio flux
is very faint which may indicate absorption of the flux. Thus it is possible
that the bright IR is arising from the same region, which may be more
symmetric and is different from the region where low luminosity radio and
X-rays arise (see Fig. 14 of Smith & Andrews, 2020). In this scenario, the
X-rays are completely absorbed in the region of high luminosity and the
observed X-rays are from a region of weak interaction with relatively lower
column density.

##  5 Summary and conclusions

SN 2017hcc had one of the highest bolometric luminosities for a SN IIn, but
low X-ray and radio luminosities. The peak bolometric luminosity of $\sim
10^{44}$ ergs s-1 (Prieto et al., 2017) is likely to be due to CS interaction.
During the first 100 days, there are only upper limits on the X-ray
luminosity, $<3\times 10^{40}$ ergs s-1, so the ratio of X-ray to total
luminosity is $<0.0003$. These results are comparable to those observed in SN
2006gy (Ofek et al., 2007; Smith et al., 2007). SN 2017hcc has properties that
are consistent with shock breakout in a dense wind (Chevalier & Irwin, 2011).

The optical/IR lines during the first 100 days show electron scattering
profiles that require an electron scattering optical depth of at least a few,
which corresponds to a column density of at least $10^{24}$ cm-2. Our early
time Swift-XRT measurements result in a non-detection; hence, we are unable to
constrain the column depth. It is possible that the early X-rays were absorbed
due to the high column density. Shock breakout in a dense medium also results
in suppression of X-ray emission. In addition, the non-linear thin-shell
instability due to a radiative forward shock results in further suppression of
the X-rays.

There is a late time enhancement in the IR emission with peak IR luminosities
reaching $\sim 10^{42}$ erg s-1. At the same epoch, the X-ray luminosity is
two orders of magnitude lower. We interpret the high IR luminosity as likely
due to contributions from new dust as well as old dust.

We also show that SN 2017hcc likely underwent a phase of enhanced mass-loss
years before explosion. This has been seen in many SNe IIn and may indirectly
support an LBV scenario, though the fine tuning between an LBV undergoing
enhanced mass-loss and SN explosion remains an important issue. Thus SN
2017hcc adds an important input towards our understanding of highly
interacting SNe, whose progenitors remain a mystery.

## Acknowledgements

We thank the referee for useful comments, which helped improve the manuscript.
RAC acknowledges support from NASA through Chandra award number GO0-21054X
issued by the Chandra X-ray Center and from NSF grant AST-1814910. PC
acknowledges support of the Department of Atomic Energy, Government of India,
under project no. 12-R&D-TFR-5.02-0700. The Chandra Center is operated by the
Smithsonian Astrophysical Observatory for and on behalf of NASA under contract
NAS8-03060. The National Radio Astronomy Observatory is a facility of the
National Science Foundation operated under cooperative agreement by Associated
Universities, Inc. This research has made use of data and/or software provided
by the High Energy Astrophysics Science Archive Research Center (HEASARC),
which is a service of the Astrophysics Science Division at NASA/GSFC. This
work is based [in part] on observations made with the Spitzer Space Telescope,
which was operated by the Jet Propulsion Laboratory, California Institute of
Technology under a contract with NASA. Support for this work was provided by
NASA. This work made use of data supplied by the UK Swift Science Data Centre
at the University of Leicester.

## Data Availability

The data underlying this article will be shared upon a reasonable request to
the corresponding author.

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