# SN 2017hpa: A Nearby Carbon-Rich Type Ia Supernova with a Large Velocity
Gradient

Xiangyun Zeng  Xinjiang Astronomical Observatory, Chinese Academy of Sciences, Urumqi, Xinjiang 830011, People’s Republic of China  School of Astronomy and Space Science, University of Chinese Academy of Sciences, Beijing 100049, People’s Republic of China  Xiaofeng Wang<EMAIL_ADDRESS>Physics Department and Tsinghua Center for Astrophysics (THCA), Tsinghua University, Beijing, 100084, People’s Republic of China  Beijing Planetarium, Beijing Academy of Science of Technology, Beijing 100044, People’s Republic of China  Ali Esamdin<EMAIL_ADDRESS>Xinjiang Astronomical Observatory, Chinese Academy of Sciences, Urumqi, Xinjiang 830011, People’s Republic of China  Craig Pellegrino  Department of Physics, University of California, Santa Barbara, CA 93106-9530, USA  Las Cumbres Observatory, 6740 Cortona Drive Suite 102, Goleta, CA 93117-5575, USA  WeiKang Zheng  Department of Astronomy, University of California, Berkeley, CA 94720-3411, USA  Jujia Zhang  Yunnan Observatories (YNAO), Chinese Academy of Sciences, Kunming 650216, People’s Republic of China  Key Laboratory for the Structure and Evolution of Celestial Objects, Chinese Academy of Sciences, Kunming 650216, People’s Republic of China  Center for Astronomical Mega-Science, Chinese Academy of Sciences, 20A Datun Road, Chaoyang District, Beijing, 100012, People’s Republic of China  Jun Mo  Physics Department and Tsinghua Center for Astrophysics (THCA), Tsinghua University, Beijing, 100084, People’s Republic of China  Wenxiong Li  Physics Department and Tsinghua Center for Astrophysics (THCA), Tsinghua University, Beijing, 100084, People’s Republic of China  The School of Physics and Astronomy, Tel Aviv University, Tel Aviv 69978, Israel  D. Andrew Howell  Las Cumbres Observatory, 6740 Cortona Drive Suite 102, Goleta, CA 93117-5575, USA  Department of Physics, University of California, Santa Barbara, CA 93106-9530, USA  Alexei V. Filippenko  Department of Astronomy, University of California, Berkeley, CA 94720-3411, USA  Miller Senior Fellow, Miller Institute for Basic Research in Science, University of California, Berkeley, CA 94720, USA  Han Lin  Physics Department and Tsinghua Center for Astrophysics (THCA), Tsinghua University, Beijing, 100084, People’s Republic of China  Thomas G. Brink  Department of Astronomy, University of California, Berkeley, CA 94720-3411, USA  Edward A. Baron  Homer L. Dodge Department of Physics and Astronomy, University of Oklahoma, USA  Jamison Burke  Las Cumbres Observatory, 6740 Cortona Drive Suite 102, Goleta, CA 93117-5575, USA  Department of Physics, University of California, Santa Barbara, CA 93106-9530, USA  James M. DerKacy  Homer L. Dodge Department of Physics and Astronomy, University of Oklahoma, USA  Curtis McCully  Las Cumbres Observatory, 6740 Cortona Drive Suite 102, Goleta, CA 93117-5575, USA  Department of Physics, University of California, Santa Barbara, CA 93106-9530, USA  Daichi Hiramatsu  Las Cumbres Observatory, 6740 Cortona Drive Suite 102, Goleta, CA 93117-5575, USA  Department of Physics, University of California, Santa Barbara, CA 93106-9530, USA  Griffin Hosseinzadeh  Center for Astrophysics | Harvard & Smithsonian, 60 Garden Street, Cambridge, MA 02138-1516, USA  Las Cumbres Observatory, 6740 Cortona Drive Suite 102, Goleta, CA 9311-5575, USA  Department of Physics, University of California, Santa Barbara, CA 93106-9530, USA  Benjamin T. Jeffers  Department of Astronomy, University of California, Berkeley, CA 94720-3411, USA  Timothy W. Ross  Department of Astronomy, University of California, Berkeley, CA 94720-3411, USA  Benjamin E. Stahl  Department of Astronomy, University of California, Berkeley, CA 94720-3411, USA  Department of Physics, University of California, Berkeley, CA 94720-7300, USA  Marc J. Staley Graduate Fellow  Samantha Stegman  Department of Astronomy, University of California, Berkeley, CA 94720-3411, USA  Department of Chemistry, University of Wisconsin, Madison, WI 53706, USA  Stefano Valenti  Department of Physics, University of California, Davis, CA 95616, USA  Lifan Wang  George P. and Cynthia Woods Mitchell Institute for Fundamental Physics & Astronomy, Texas A&M University, Department of Physics and Astronomy, 4242 TAMU, College Station, TX 77843, USA  Danfeng Xiang  Physics Department and Tsinghua Center for Astrophysics (THCA), Tsinghua University, Beijing, 100084, People’s Republic of China  Jicheng Zhang  Physics Department and Tsinghua Center for Astrophysics (THCA), Tsinghua University, Beijing, 100084, People’s Republic of China  Tianmeng Zhang  Key Laboratory of Optical Astronomy, National Astronomical Observatories, Chinese Academy of Sciences, Beijing 100012, People’s Republic of China

###### Abstract

We present extensive, well-sampled optical and ultraviolet photometry and
optical spectra of the Type Ia supernova (SN Ia) 2017hpa. The light curves
indicate that SN 2017hpa is a normal SN Ia with an absolute peak magnitude of
$M_{\rm max}^{B}\approx$ $-19.12\pm 0.11$ mag and a post-peak decline rate
${\Delta}m_{15}(B)$ = $1.02\pm 0.07$ mag. According to the quasibolometric
light curve, we derive a peak luminosity of $1.25\times 10^{43}$ erg s-1and a
56Ni mass of $0.63\pm 0.02\,M_{\odot}$. The spectral evolution of SN 2017hpa
is similar to that of normal SNe Ia, while it exhibits unusually rapid
velocity evolution resembling that of SN 1991bg-like SNe Ia or the high-
velocity subclass of SNe Ia, with a post-peak velocity gradient of $\sim$
$130\pm 7$ km s-1 d-1. Moreover, its early spectra ($t<-7.9$ d) show prominent
C ii $\lambda$6580 absorption feature, which disappeared in near-maximum-light
spectra but reemerged at phases from $t\,\sim\,+8.7$ d to $t\,\sim\,+11.7$ d
after maximum light. This implies that some unburned carbon may mix deep into
the inner layer, and is supported by the low C ii $\lambda$6580 to Si ii
$\lambda$6355 velocity ratio ($\sim 0.81$) observed in SN 2017hpa. The O i
$\lambda$7774 line shows a velocity distribution like that of carbon. The
prominent carbon feature, low velocity seen in carbon and oxygen, and large
velocity gradient make SN 2017hpa stand out from other normal SNe Ia, and are
more consistent with predictions from a violent merger of two white dwarfs.
Detailed modelling is still needed to reveal the nature of SN 2017hpa.

supernovae: individual: SN 2017hpa — supernovae: general: high velocity
gradient

††software: SNooPy (Burns et al., 2011, 2014), SN-Spectral Evolution
(https://github.com/mwvgroup/SN-Spectral-Evolution), Minim Code (Chatzopoulos
et al., 2013), IRAF (Tody, 1986, 1993), DAOPHOT (Stetson, 1987), PyZOGY
(Zackay et al., 2016; Guevel & Hosseinzadeh, 2017), lcogtsnpipe (Valenti et
al., 2016), LOSS data-reduction pipeline (Ganeshalingam et al., 2010; Stahl et
al., 2019, 2020), Astropy (Astropy Collaboration et al., 2013), Matplotlib
(Hunter et al., 2007), Scipy (https://www.scipy.org/), Numpy
(https://numpy.org/)

##  1 Introduction

Type Ia supernovae (SNe Ia) are widely believed to arise from explosions of
carbon-oxygen (CO) white dwarfs (WDs) in a binary system, which have a typical
absolute $V$-band peak magnitude of $\sim-19$ mag (Phillips, 1993; Perlmutter
et al., 1999; Wang et al., 2006). The relatively uniform stellar explosions of
SNe Ia make them useful as standardizable candles in measuring extragalactic
distances (Phillips, 1993; Riess et al., 1996; Wang et al., 2005; Guy et al.,
2005; Howell, 2011; Burns et al., 2018), leading to the discovery of the
accelerating of the Universe (Riess et al., 1998; Perlmutter et al., 1999). In
recent years, larger samples of SNe Ia have been used to further constrain the
nature of dark energy driving the acceleration (e.g., Betoule et al., 2014;
Abbott et al., 2019).

However, the progenitor systems and explosion mechanism of SNe Ia still remain
controversial (e.g., Maoz et al., 2014). Two popular scenarios are the violent
merger-triggered explosion of two WDs, known as the double-degenerate (DD)
scenario (Webbink, 1984; Iben & Tutukov, 1984), and the accretion-triggered
explosion of a WD with a nondegenerate companion, known as the single-
degenerate (SD) scenario (Whelan & Iben, 1973; Nomoto, 1982; Nomoto et al.,
1997). In general, the detection of signatures of circumstellar material (CSM)
around some SNe Ia supports the SD scenario (Hamuy et al., 2003; Aldering et
al., 2006; Patat et al., 2007; Sternberg et al., 2011; Dilday et al., 2012;
Maguire et al., 2013; Silverman et al., 2013; Wang et al., 2019), though some
theoretical studies show that the CSM could be also produced in the DD
scenario (Shen et al., 2013; Raskin & Kasen, 2013). On the other hand, there
is also evidence for nondetections of companion signatures for some SNe Ia,
thus favoring the DD scenario (González Hernández et al., 2012; Schaefer &
Pagnotta, 2012; Olling et al., 2015; Tucker et al., 2019).

Popular explosion models of SNe Ia include the following cases. (1) The CO WD
accretes material from the companion star until its mass nearly reaches the
Chandrasekhar mass limit ($M_{\rm Ch}$, $\sim$1.4 $M_{\odot}$, Chandrasekhar,
1957) and compressional heating at the center causes the explosion (Piersanti
et al., 2004). (2) The detonation of a thin layer of He on the surface of a WD
(Kromer et al., 2010; Shen & Moore, 2014) triggers a second detonation in the
WD center and hence the explosion of a sub-$M_{\rm Ch}$ mass C-O WD. (3) The
violent merger or secular merger of two WDs, accompanied by radiation of
gravitational waves (Röpke et al., 2012; García-Berro & Lorén-Aguilar, 2017).
(4) In triple systems, oscillations of the third star cause a direct collision
of two WDs and trigger the SN explosion (Thompson, 2011; Mazzali et al.,
2018). In view of these explosion mechanisms, the delayed detonation model is
one of the most suitable ones to account for the observed properties of SNe
Ia, which initially involves a deflagration of a $M_{\rm Ch}$ CO WD and later
a supersonic detonation (Khokhlov, 1991; Höflich et al., 2017). Nevertheless,
the double detonation model of sub-$M_{\rm Ch}$ CO WDs shows many striking
features, and can also explain the observed properties of some SNe Ia (Shen et
al., 2018).

Observationally, there is increasing evidence for spectroscopic and
photometric diversity of SNe Ia. Most SNe Ia can be classified as
spectroscopically normal ones, while a small fraction exhibit peculiar
properties in some respects (e.g., Branch et al., 1993; Filippenko, 1997),
such as the SN 1991T-like overluminous SNe (Filippenko et al., 1992a; Ruiz-
Lapuente et al., 1992; Phillips, 1993), the SN 1991bg-like subluminous SNe
(Filippenko et al., 1992b; Leibundgut et al., 1993), or the SN 2002cx-like
subclasses (Filippenko, 2003; Li et al., 2003). Based on differences in Si ii
velocity evolution, Benetti et al. (2005) divided SNe Ia into three
subclasses: high velocity gradients (HVG), low velocity gradients (LVG), and
FAINT. According to the equivalent width (EW) of Si ii $\lambda$6355 and Si ii
$\lambda$5972 absorption lines near maximum brightness, Branch et al. (2006)
divided SNe Ia into core normal (CN), broad line (BL), cool (CL), and shallow
silicon (SS) subgroups. The sublumious SN 1991bg-like and overluminous SN
1991T-like SNe Ia have large overlap with the CL and SS subclasses,
respectively (Branch et al., 2006). Based on the Si ii $\lambda$6355 velocity
measured near the time of $B$-band maximum, Wang et al. (2009) classified SNe
Ia into normal-velocity (NV) and high-velocity (HV) subsets. The HV subclass
is found to share some common properties such as red $B-V$ color, slow decay
in blue bands starting at $t\approx 40$ d from the peak, and abundant
surrounding CSM (Wang et al., 2008, 2009, 2019; Foley et al., 2011; Mandel et
al., 2014). Although asymmetric explosions have been proposed to explain the
occurrence of HV and NV subclasses of SNe Ia (Maeda et al., 2010), it is
difficult to account for the fact that these two subgroups have different
birth environments (Wang et al., 2013).

Early-time observations can place important constraints on the explosion
physics of SNe Ia, including the size of the primary WD (Bloom et al., 2012),
the radius of the companion star (Hosseinzadeh et al., 2017), the distribution
of 56Ni in the ejecta, and the possible existence of CSM (Piro & Morozova,
2016). Therefore, clarifying the progenitor systems and explosion mechanisms
affects our understanding of stellar evolution and precision cosmology. The
unburned carbon detected in early-time spectra can provide important clues to
the progenitor system and explosion mechanism of SNe Ia (Yamanaka et al.,
2009; Silverman et al., 2011; Taubenberger et al., 2011; Thomas et al., 2011;
Silverman & Filippenko, 2012; Hsiao et al., 2013; Li et al., 2019a).

Previous studies show that nearly 30% of SNe Ia show signatures of C ii
$\lambda$6580 absorption at $t\approx-4$ d (relative to the time of maximum
light), while this fraction is over 40% when examining the $t\approx-10$ d
spectra (Parrent et al., 2011; Thomas et al., 2011; Folatelli et al., 2012;
Silverman & Filippenko, 2012; Maguire et al., 2014). These studies show that
carbon-positive SNe Ia tend to be LVG subtypes (Folatelli et al., 2012) and
have bluer optical colors around maximum light (Thomas et al., 2011; Silverman
& Filippenko, 2012). Among those carbon-positive SNe Ia, there are two events
which show carbon absorption lasting until 1–3 weeks after maximum light. One
is SN 2002fk, which has detecable carbon absorption lines in the $t\approx 10$
d spectrum Cartier et al. (2014). Another example is SN 2018oh studied by Li
et al. (2019a), the first Kepler-discovered SN Ia with a spectroscopic
classification; the carbon feature can be detected even in the $t\approx 20.5$
d spectrum, representing the latest detection of carbon in SNe Ia. The origin
of these carbon detections in post-maximum spectra still remains unclear. SN
2017hpa is the third SN Ia with persistent carbon feature; it exploded in the
spiral galaxy UGC 3122 (see Fig. 1) at a distance of $\sim$ 65.6 Mpc (redshift
$z\approx 0.0156$). The prominent carbon features and small distance of SN
2017hpa provide us with another excellent chance to study the observed
diversity of SNe Ia.

In this paper, the optical observations and data reduction are presented in
Section 2. Section 3 discusses the light and color curves, while Section 4
shows the spectroscopic evolution. The quasibolometric light curve and origin
of prominent carbon feature of SN 2017hpa are discussed in Section 5. We
summarize in Section 6.

##  2 Observations and Data Reduction

###  2.1 Discovery and Host Galaxy

SN 2017hpa was discovered at $\alpha=04^{h}39^{m}50^{s}.750$,
$\delta=07^{\circ}03\arcmin 54\arcsec.90$ (J2000) on 2017 Oct. 25.35 (UT dates
are adopted throughout this paper) during the Puckett Observatory World
Supernova Search (POSS) Gagliano et al. (2017). Figure 1 shows a color image
of SN 2017hpa. A spectrum taken $\sim 0.65$ d after the discovery classified
it as a normal SN Ia (Floers et al., 2017).

Figure 1: The left panel shows a color image synthesized from $gri$-band
observations from PanSTARRS, and the faint star right beside the position of
SN 2017hpa is totally covered in following observations. The right panel shows
a color image synthesized from $gri$-band observations from TNT; SN 2017hpa is
marked with a red circle while the reference stars are numbered.

The host galaxy of SN 2017hpa is UGC 3122, which is classified as SAB(rc)s at
$z=0.015631\pm 0.000005$ (Paturel et al., 2002; Springob et al., 2005). This
redshift corresponds to a distance modulus $\mu=34.05\pm 0.38$ mag with a
velocity uncertainty of 500 km s-1 (Willick et al., 1997), assuming a Hubble
constant of 73.5 km s-1 Mpc-1 (Riess et al., 2018).

###  2.2 Photometry

After the discovery of SN 2017hpa, we triggered follow-up photometric
observations on several telescopes, including the 0.8 m Tsinghua-NAOC
telescope (TNT; Huang et al., 2012; Zhang et al., 2015), the Las Cumbres
Obervatory (LCO) Telescope network (Shporer et al., 2011; Brown et al., 2013),
the 0.76 m Katzman Automatic Imaging Telescope (KAIT) at Lick Observatory (Li
et al., 2001; Filippenko et al., 2001; Filippenko, 2005), and the 1 m Nickel
reflector at Lick Observatory. The TNT, Nickel, and KAIT monitored SN 2017hpa
in the $BVRI$ bands, while the LCO 1 m telescope sampled its light curves in
the $BVgri$ bands.

For photometric observations obtained from the LCO during the Global Supernova
Project, PyZOGY (Zackay et al., 2016; Guevel & Hosseinzadeh, 2017) is employed
for image subtraction while the $lcogtsnpipe$ (Valenti et al., 2016) is
applied for measuring the SN flux. An ad hoc pipeline (based on the IRAF
DAOPHOT package; Stetson, 1987) is applied to reduce images from the TNT and
extract instrumental magnitudes of the SN. Photometric images from the Lick
Observatory are reduced using the Lick Observatory Supernova Search (LOSS;
Filippenko et al., 2001) data-reduction pipeline (Ganeshalingam et al., 2010;
Stahl et al., 2019, 2020), while DAOPHOT (Stetson, 1987) is applied to
implement the point-spread-function (PSF) photometry.

For the TNT instrumental magnitudes, the $BV$-band images are calibrated using
the APASS catalog (Henden et al., 2016) and the $gri$-band magnitudes are
calibrated using the PanSTARRS catalog (Chambers et al., 2016; Waters et al.,
2016; Flewelling et al., 2016; Magnier et al., 2016). The local standard stars
with photometric magnitudes from APASS and PanSTARRS are listed in Table 1.
The unfiltered instrumental magnitudes from LOSS are calibrated to the
standard Landolt $R$-band magnitudes based on the transformed local standards
of SDSS (Li et al., 2003; Zheng et al., 2017). The LOSS $BVRI$ instrumental
magnitudes are calibrated to the Johnson system using a series of Landolt
(Landolt, 1992) standard stars taken on a number of photometric nights.

Table 1: Photometric Standards in the SN 2017hpa Field 1aaStandard stars used for calibration of instrumental magnitudes. Star |  $\alpha$(J2000) |  $\delta$(J2000) |  $B$ (mag) |  $V$ (mag) |  $g$ (mag) |  $r$ (mag) |  $i$ (mag)
---|---|---|---|---|---|---|---
1 | 04:40:00.331 | +07:02:17.167 | 14.125(020) | 13.368(015) | 13.689(014) | 13.114(034) | 12.915(040)
2 | 04:40:03.904 | +07:01:18.451 | 16.618(065) | 15.697(037) | 16.160(089) | 15.398(107) | 15.110(070)
3 | 04:39:40.185 | +07:02:54.100 | 15.836(029) | 15.046(049) | 15.403(030) | 14.779(034) | 14.593(045)
4 | 04:40:00.268 | +07:06:44.968 | 14.140(034) | 13.228(016) | 13.625(016) | 12.887(036) | 12.632(035)
5 | 04:40:02.771 | +07:00:42.491 | 16.812(121) | 16.013(077) | 16.389(088) | 15.752(053) | 15.491(174)
6 | 04:39:52.514 | +07:01:50.567 | 17.515(169) | 16.242(070) | 16.645(040) | 15.957(026) | 15.637(040)
7 | 04:39:49.255 | +07:04:03.612 | 17.125(095) | 16.130(077) | 16.594(044) | 15.802(069) | 15.477(068)
8 | 04:39:57.276 | +07:04:14.560 | 15.626(052) | 14.823(023) | 15.191(018) | 14.556(044) | 14.326(043)
9 | 04:39:54.303 | +07:01:52.794 | 17.553(151) | 16.262(008) | 16.918(114) | 15.686(037) | 15.101(034)
10 | 04:39:58.351 | +07:00:56.540 | 17.480(238) | 16.547(158) | 16.996(066) | 16.230(041) | 15.995(040)
11 | 04:39:41.224 | +07:05:43.976 | 17.430(026) | 16.468(078) | 17.033(093) | 16.288(066) | 16.102(040)

Optical and ultraviolet (UV) observations of SN 2017hpa were also obtained
with the Neil Gehrels Swift Observatory (Gehrels et al., 2004). The Swift/UVOT
observations started at relatively early phases in six bands including $uvw2$,
$uvm2$, $uvw1$, $u$, $b$, and $v$ (Roming et al., 2005). The filters in lower
case are used throughout this paper for photometric observations in UVOT
bands. Using zeropoints extracted by Breeveld et al. (2011) in the Vega
system, the data-reduction pipeline of the Swift Optical/Ultraviolet Supernova
Archive (SOUSA; Brown et al., 2014) is applied to obtain the Swift optical/UV
light curves of SN 2017hpa. The source counts are measured using a 3$\arcsec$
aperture and corrections are based on the average PSF. The template-
subtraction technique has also been applied to the Swift images and the final
uncertainty in the photometry is the combination of statistical uncertainties
in galaxy subtraction count rates and a 2% systematic fluctuation at each
pixel caused by differences in response sensitivity across the photon
detector. The final observed Swift and ground-based light curves are shown in
Figure 2, and the corresponding magnitudes are tabulated in Table 2 and Table
3.

Table 2: Photometric Observations of SN 2017hpa by Ground-Based Telescopes  MJD | EpochaaRelative to the epoch of $B$-band maximum brightness (MJD = 58,066.6). Magnitudes are calibrated to AB magnitude system. |  $B$ (mag) |  $V$ (mag) |  $R$ (mag) |  $I$ (mag) |  $g$ (mag) |  $r$ (mag) |  $i$ (mag) |  $Clear$ (mag) | Telescope
---|---|---|---|---|---|---|---|---|---|---
58053.34 | -13.30 | 17.235(035) | 16.967(022) |  $\cdots$ |  $\cdots$ | 17.054(081) | 17.060(097) | 17.311(143) |  $\cdots$ | TNT
58053.54 | -13.10 | 16.971(031) | 16.869(027) | 16.741(023) | 16.668(039) |  $\cdots$ |  $\cdots$ |  $\cdots$ | 16.623(022) | KAIT4
58053.82 | -12.83 | 16.946(041) | 16.848(030) |  $\cdots$ |  $\cdots$ | 16.981(025) | 16.863(065) | 17.078(043) |  $\cdots$ | LCO
58054.35 | -12.29 | 16.897(033) | 16.684(020) |  $\cdots$ |  $\cdots$ | 16.896(167) | 16.806(102) | 17.122(136) |  $\cdots$ | TNT
58054.36 | -12.28 | 16.741(012) | 16.634(008) | 16.539(010) | 16.463(015) |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | Nickel
58054.36 | -12.28 | 16.744(012) | 16.642(011) | 16.546(011) | 16.463(016) |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | KAIT4
58054.54 | -12.11 | 16.717(027) | 16.614(021) | 16.505(019) | 16.442(190) |  $\cdots$ |  $\cdots$ |  $\cdots$ | 16.355(024) | KAIT4
58054.84 | -11.80 | 16.705(043) | 16.659(042) |  $\cdots$ |  $\cdots$ | 16.692(022) | 16.614(058) | 16.856(045) |  $\cdots$ | LCO
58055.19 | -11.45 | 16.674(031) | 16.484(017) |  $\cdots$ |  $\cdots$ | 16.469(074) | 16.507(089) | 16.846(130) |  $\cdots$ | TNT
58055.54 | -11.10 | 16.460(021) | 16.409(017) | 16.276(017) | 16.217(030) |  $\cdots$ |  $\cdots$ |  $\cdots$ | 16.147(016) | KAIT4
$\vdots$ | $\vdots$ | $\vdots$ | $\vdots$ | $\vdots$ | $\vdots$ | $\vdots$ | $\vdots$ | $\vdots$ | $\vdots$ | $\vdots$
58203.00 | +136.35 | 19.821(478) | 18.920(192) |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | TNT
58207.17 | +140.52 | 20.258(264) | 19.730(330) | 20.210(455) | 19.657(309) |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | Nickel
58207.17 | +140.52 | 20.281(285) | 19.851(363) | 20.501(477) | 19.562(302) |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | KAIT4
58210.16 | +143.52 | 20.056(159) | 19.796(416) | 19.774(295) |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | Nickel
58210.16 | +143.52 | 20.021(159) | 19.960(435) | 19.991(337) | 19.606(277) |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ | KAIT4
Table 3: Swift UVOT Photometry of SN 2017hpa MJD | EpochaaRelative to the epoch of $B$-band maximum (MJD = 58,066.6). Magnitudes are calibrated to Vega magnitude system. |  $uvw2$ (mag) |  $uvm2$ (mag) |  $uvw1$ (mag) |  $u$ (mag) |  $b$ (mag) |  $v$ (mag)
---|---|---|---|---|---|---|---
58053.19 | -13.45 | 19.73(16) | 20.18(33) | 18.87(13) | 17.51(06) | 17.06(04) | 16.66(05)
58054.59 | -12.06 | 20.18(50) |  $\cdots$ | 18.98(26) | 16.92(11) | 16.73(07) | 16.47(11)
58055.71 | -10.93 | 19.41(17) | 20.40(33) | 18.25(12) | 16.35(05) | 16.29(04) | 16.22(06)
58060.76 | -5.89 | 18.70(09) | 19.72(24) | 17.24(06) | 15.43(03) | 15.62(03) | 15.59(04)
58067.35 | 0.70 | 18.30(13) | 19.52(28) | 17.24(12) | 15.36(05) | 15.42(04) | 15.22(06)
58069.00 | 2.35 | 18.61(11) | 19.63(20) | 17.45(09) | 15.42(03) | 15.42(03) | 15.25(04)
58071.00 | 4.34 | 18.68(10) | 20.23(30) | 17.69(10) | 15.63(04) | 15.54(03) | 15.31(04)
58073.05 | 6.41 | 18.88(12) | 19.59(20) | 17.67(10) | 15.87(04) | 15.66(03) | 15.34(04)
58074.59 | 7.95 | 18.97(12) | 19.67(19) | 17.90(11) | 15.98(04) | 15.79(03) | 15.39(04)
58080.95 | 14.31 | 19.17(13) | 20.00(24) | 18.55(13) | 16.85(06) | 16.40(04) | 15.77(04)
58082.35 | 15.71 | 19.12(19) | 20.52(51) | 18.62(18) | 16.97(09) | 16.56(05) | 15.80(07)
58085.14 | 18.49 | 19.70(15) | 19.94(18) | 19.14(14) | 17.36(06) | 16.92(04) | 16.00(04)
58088.60 | 21.96 | 19.96(20) | 20.11(22) | 19.50(17) | 17.63(08) | 17.21(05) | 16.27(05)
58095.17 | 28.52 | 20.02(19) | 20.29(23) | 19.96(22) | 18.38(12) | 17.75(06) | 16.46(06)
Figure 2: The observed UV and optical light curves of SN 2017hpa.

###  2.3 Spectroscopy

A total of 26 low-resolution optical spectra of SN 2017hpa have been obtained
using different telescopes and equipment, including the AFOSC mounted on the
Asiago Ekar telescope, the BFOSC mounted on the Xinglong 2.16 m telescope
(XLT; Jiang et al., 1999; Zhang et al., 2016; Fan et al., 2016), the YFOSC on
the Lijiang 2.4 m telescope (LJT; Chen et al., 2001; Zhang et al., 2012; Wang
et al., 2019) of Yunnan Astronomical Observatories, the Kast spectrograph on
the Lick 3 m Shane telescope (Miller & Stone, 1993; Stahl et al., 2020), and
the LCO 2 m Faulkes Telescope North (FTN; Brown et al., 2013). The journal of
spectroscopic observations is presented in Table 4, including one spectrum
from the Transient Name Server111https://wis-tns.weizmann.ac.il/ and six
spectra from Stahl et al. (2020). When no data-reduction pipeline for was
available, we applied standard IRAF routines to reduce the spectra.
Spectrophotometric standard stars observed at an airmass comparable to the
target on the same night were used to calibrate the flux density of SN
2017hpa. The extinction curves of the various observatories are utilized to
correct for atmospheric extinction, and spectra of the standard stars are used
to eliminate the telluric absorption lines.

Table 4: Spectroscopic Observations of SN 2017hpa MJD | EpochaaRelative to the epoch of $B$-band maximum (MJD = 58,066.6). | $\lambda_{\rm Start}$ | $\lambda_{\rm End}$ | Instrument
---|---|---|---|---
58052.5 | -14.1 | 3387 | 8249 | Asiago (public)
58053.2 | -13.5 | 3496 | 9173 | LJT
58054.2 | -12.4 | 3502 | 9171 | LJT
58056.5 | -10.1 | 3622 | 10400 | Lick 3 m
58058.7 | -7.9 | 3399 | 9999 | LCO
58064.5 | -2.1 | 3249 | 9999 | LCO
58068.7 | 2.0 | 3249 | 9999 | LCO
58071.3 | 4.7 | 3746 | 8840 | XLT
58075.4 | 8.7 | 3250 | 10000 | LCO
58076.1 | 9.5 | 3744 | 8839 | XLT
58078.4 | 11.7 | 3630 | 10400 | Lick 3 m
58091.7 | 25.1 | 3500 | 9100 | LJT
58094.3 | 27.7 | 3299 | 9999 | LCO
58099.3 | 32.6 | 3299 | 9999 | LCO
58099.3 | 32.6 | 3630 | 10400 | Lick 3 m
58105.4 | 38.7 | 3632 | 10400 | Lick 3 m
58107.7 | 41.1 | 3500 | 9100 | LJT
58110.3 | 43.7 | 5896 | 8182 | Lick 3 m
58111.3 | 44.7 | 3249 | 10000 | LCO
58118.7 | 52.1 | 3500 | 9100 | LJT
58119.4 | 52.8 | 3300 | 10000 | LCO
58121.3 | 54.7 | 3299 | 10000 | LCO
58127.8 | 61.1 | 3500 | 9100 | LJT
58131.4 | 64.7 | 3632 | 10400 | Lick 3 m
58132.3 | 65.7 | 3400 | 9300 | LCO
58153.3 | 86.6 | 3300 | 9299 | LCO

##  3 Light Curves

###  3.1 Optical and Ultraviolet Light Curves

The multiband UV/optical light curves of SN 2017hpa are shown in Figure 2; one
can see that the observations in optical bands have nearly daily sampling,
ranging from about 2 weeks before to over 100 d after $B$-band maximum light.
The light curves of SN 2017hpa are similar to those of normal SNe Ia, reaching
maximum slightly earlier in the $I/i$ and $UV$ bands than the $B$ band, and
having a prominent shoulder in $R/r$ as well as a secondary maximum in $I/i$.
The slight deviations between the $BVgri$ light curves of different telescopes
are primarily due to different filter transmission functions, as shown in
Figure 3. The transmission differences at the red edge of the $I$-band filters
may cause the $I$-band discrepancies between LCO and TNT. Applying a
polynomial fit to the $B$-band light curves around maximum light yields a peak
of $15.48\pm 0.03$ mag on MJD = 58066.6 (UT 2017 November 9.6). The $V$-band
light curve reached its peak of $15.35\pm 0.2$ mag on MJD = 58068.4, $\sim
1.8$ d after the $B$-band peak.

Figure 3: The transmission curves of the TNT and LCO filters; each curve is
normalized to the peak.

Figures 4 and 5 compare the multiband UV/optical light curves of SN 2017hpa
with those of several well-observed normal SNe Ia which have comparable
${\Delta}m_{15}(B)$, including SN 2003du(Stanishev et al., 2007), SN 2005cf
(Wang et al., 2009), SN 2011fe (Maguire et al., 2013), SN 2012cg (Munari et
al., 2013; Brown et al., 2014), SN 2013dy (Pan et al., 2015; Zhai et al.,
2016), and SN 2018oh (Li et al., 2019a). The UV/optical light curves of the
comparison SNe Ia have been normalized to SN 2017hpa. As can be seen from
Figure 4, SN 2017hpa and other normal comparision SNe Ia have similar light-
curve shapes near $B$-band maximum. Although the UV light curves of SN 2017hpa
are similar to those of other comparison SNe Ia, they seem to show excess
emission at early phases, especially the first two data points. This may
suggest additional energy beyond the radioactive decay of centrally-located
nickel, such as surface nickel mixing (Piro & Nakar, 2013) or interaction of
SN ejecta with a companion star or with CSM (Kasen, 2010). The post-peak
decline rate ${\Delta}m_{15}(B)$ of the $B$-band light curve is measured to be
$1.02\pm 0.07$ mag, and the color stretch (Burns et al., 2014) is determined
to be $s_{BV}$= $0.94\pm 0.03$.

Figure 4: Comparison of the optical light curves of SN 2017hpa with other
well-observed SNe Ia having similar decline rates. The light curves of the
comparison SNe Ia have been normalized to match the observed peak magnitudes
of SN 2017hpa.  Figure 5: Comparison of the UV light curves of SN 2017hpa with
other well-observed SNe Ia having similar decline rates. The light curves of
the comparison SNe Ia have been normalized to match the peak magnitudes of SN
2017hpa.

###  3.2 Reddening and Color Curves

Assuming $R_{V}=3.1$ (Cardelli et al., 1989), we obtain the line-of-sight
Galactic extinction for SN 2017hpa to be $A_{V}=0.485$ mag (Schlegel et al.,
1998; Schlafly & Finkbeiner, 2011), corresponding to a color excess of
$E(B-V)_{\rm gal}=0.156$ mag. After removing the Galactic reddening, the $B-V$
color is found to be $0.002\pm 0.05$ mag at $t=0$ d and $1.08\pm 0.06$ mag at
$t=35$ d relative to $B$ maximum, consistent with typical values of normal SNe
Ia (Phillips et al., 1999; Wang et al., 2009).

We applied SuperNovae in object-oriented Python (SNooPy; Burns et al., 2011,
2014) to fit the multiband light curves of SN 2017hpa, as shown in Figure 6.
Both $EBV$ and $st$ models in SNooPy2 are adopted to estimate the host-galaxy
extinction, and an average host reddening is derived to be $E(B-V)_{\rm host}$
= $0.06\pm 0.06$ mag. The relatively low host-galaxy reddening is consistent
with the fact that the SN is located far away from the center of the host
galaxy. Moreover, the spectra of SN 2017hpa show no detectable absorption
feature of Na i D due to host galaxy.

Figure 6: Best-fit light-curve model from SNooPy2. The light curves are
shifted vertically for clarity. The dashed lines represent the 1$\sigma$
uncertainty of the best-fit light-curve templates.

The optical intrinsic color evolution of SN 2017hpa is shown in Figure 7. At
$t\gtrsim-10$ d, both the $B-V$ and $g-r$ color curves evolve toward the red
until reaching the reddest color at 4–5 weeks after $B$ maximum. Both the
$V-I$ and $g-i$ color curves show a short-term evolution from red to blue
until $t\approx-10$ d; then they evolve redward and reach the red peak at
$t\approx 35$ d. After that, the $V-I$ and $g-i$ color curves became
progressively bluer.

Figure 7: The $B-V$, $V-I$, $g-r$, and $g-i$ color curves of SN 2017hpa
compared with those of SNe 2003du, 2005cf, 2011fe, 2012cg, 2013dy, and 2018oh.
All light curves including those of SN 2017hpa have been dereddened using
SNooPy2.

Overall, the color-curve evolution of SN 2017hpa is similar to that of SN
2005cf and SN 2018oh, except it has a bluer color at very early phases
(especially the $g-r$ color). Based on the near-UV (NUV) colors, SNe Ia can be
classified into NUV-red and NUV-blue subgroups (Milne et al., 2013). Figure 8
shows the observed $uvw1-V$ color evolution of SN 2017hpa together with that
of SNe 2005cf, 2011fe, and 2018oh. One can see that SN 2017hpa can be put into
the NUV-red group.

Figure 8: The $uvw1-v$ color of SN 2017hpa compared to the group of NUV-blue
and NUV-red SNe. The pink shaded region represents the regions covered by SNe
Ia that are classified as NUV-red and the blue shaded region represents the
regions covered by SNe Ia that are classified as NUV-blue. The overplotted
curve is the unreddened color evolution of SN 2017hpa (see Milne et al.,
2013).

###  3.3 First-Light Time

The rise time and first-light time can put additional constraints on the
radius of the exploding star itself (Bloom et al., 2012; Piro & Nakar, 2013).
The observation on Oct. 13, 2017 by Gagliano et al. (2017) provides a non-
detection limit magnitude of 20.5 mag. However, this observation is about 12
days before the discovery and can not provide very useful constraint for the
explosion date and hence the rise time of the light curves. We thus only
utilized the discovery magnitude, i.e., 17.9 mag $\pm$ 0.3 mag in clear band
(close to broadband $R$), obtained at $\sim$2.0 days before our multi-color
observations when performing the rise time fitting. The ideal expanding
fireball (Riess et al., 1999) model and broken-power-law (Zheng & Filippenko,
2017) model are both adopted to fit the $R$-band light curve of SN 2017hpa (as
shown in Figure 9), and the first light time of the light curve is estimated
as 58047.08 $\pm$ 0.73 days and 58049.65 $\pm$ 0.24 days, respectively. The
mean fitted first-light time (FFLT) is adopted as 58048.37 $\pm$ 0.97 days.

Figure 9: Fit to the observed $R$-band light curves using the analytic
function from Zheng & Filippenko (2017) and the ideal fireball model (Riess et
al., 1999). The black triangle represents the discovery magnitude, i.e., 17.9
mag $\pm$ 0.3 mag in clear band (close to broadband R), obtained at $\sim$2.0
days before our multi-color observations. The bottom panel shows the residual
relative to the best-fit curves. The horizontal dashed line in the bottom
panel represents zero residual.

With the time of maximum light in $B$ and the derived FFLT, the rise time of
SN 2017hpa is estimated to be $18.26\pm 0.97$ d, comparable to that of typical
SNe Ia (Zheng et al., 2017). The first multi-color observation of SN 2017hpa
is thus estimated to be $\sim 5$ d after the FFLT, $\sim 13$ d prior to $B$
maximum.

##  4 Optical Spectra

###  4.1 Temporal Evolution of the Spectra

The evolution of the optical spectra of SN 2017hpa is displayed in Figure 10.
The early-time spectra are characterized by prominent absorption lines of
intermediate-mass elements (IMEs), such as Fe ii $\lambda\lambda$4404,5018, Mg
ii $\lambda$4481, Si ii $\lambda$6355, S ii $\lambda\lambda$5468,5654, Ca ii
NIR triplet and Ca ii H&K. At $t\sim$ 2 weeks before the $B$-band maximum, the
absorption troughs near 4300 $\rm\AA$ and 4800 $\rm\AA$ could be attributed to
Fe ii/Fe iii/Mg ii, while the distinct absorption notches near 6300 $\rm\AA$
and 7000 $\rm\AA$ could be due to C ii $\lambda$6580 and C ii $\lambda$7234,
respectively. The C ii $\lambda$6580 absorption is relatively strong while the
C ii $\lambda$7234 absorption is weaker. The Si ii $\lambda$6355 absorption
lines at this phase display a perfect gaussian profile without invoking the
high-velocity feature (HVF), while a Ca ii NIR HVF could be detected through
multi-gaussian fitting (Zhao et al., 2015, 2016). After $t\sim$ 1 week before
maximum light, both C ii $\lambda$6580 and C ii $\lambda$7234 absorptions are
still prominent in the spectra of SN 2017hpa, and the absorption lines of
“W”-shaped S ii and Si ii $\lambda$5972 start to emerge in the spectra of SN
2017hpa. With the decreasing of the expansion velocity of the photosphere, the
absorption minimum of Si ii $\lambda$6355 line gradually shifted redward, and
the absorption lines of iron group elements and sulfur gradually increase in
its strength. At around $B$-band maximum, the spectra are primarily dominated
by “W”-shaped S ii absorption features near 5400 $\rm\AA$, the blended
absorption lines of Fe ii and Si ii/Si iii near 4500 $\rm\AA$ and Si ii
$\lambda$6355, while the C II features become invisible at this phase. By
$t\sim$ 0 days, the HVFs of Ca ii NIR triplet become relatively weak and the
photospheric component started to emerge in the spectra. At $t\sim+10$ days
after the $B$ maximum, the photospheric components of the Ca ii NIR continute
to gain the strength and start to dominate the spectra features.
Interestingly, the C ii $\lambda$6580 absorption feature seems to reemerge in
the spectra of SN 2017hpa around this phase, which is rarely seen in other SNe
Ia. At about one month after the $B$-band maximum, the Ca ii H&K lines and NIR
triplet are the main spectral features. Meanwhile, the features of iron group
element begin to dominate in the spectra when the SN enter the early nebular
phase. Figure 11 compares spectra of SN 2017hpa at several epochs with those
of well-observed SNe Ia with similar ${\Delta}m_{15}(B)$.

Figure 10: Optical spectral evolution of SN 2017hpa. All of the spectra have
been corrected for the redshift of the host galaxy and reddening. The epochs
shown on the right side represent the phases in days relative to $B$-band
maximum light. The dashed line marks the center of Si ii $\lambda$6355 line
profile at +2.04 d from $B$-band maximum. The color of the spectrum stands for
different instruments. The spectra have been shifted vertically for clarity.
Figure 11: Spectra of SN 2017hpa at $t\approx-14$, $-8$, -2, and +32 d
relative to $B$-band maximum, compared with spectra of SNe 2005cf (Garavini et
al., 2007; Wang et al., 2009), 2011fe (Mazzali et al., 2014; Zhang et al.,
2016), and 2013dy (Zheng et al., 2013; Pan et al., 2015; Zhai et al., 2016) at
comparable phases. Correction of reddening and redshift of the host galaxy had
been done for all of the given spectra. The spectra have been shifted
vertically for clarity.

For SN 2017hpa, the C ii $\lambda$6580 and C ii $\lambda$7234 absorptions
appear stronger than those of the comparison SNe Ia in the early-phase
spectra, as shown in Figure 11(a). Moreover, the O i $\lambda$7774 absorption
line of SN 2017hpa seems to also be stronger than in the comparison SNe Ia
except for SN 2011fe and SN 2012cg. The pseudo equivalent-width (pEW) of C ii
$\lambda$6580 is measured to be 14.0$\pm$1.0 $\rm\AA$ for SN 2017hpa at
$t\approx-13$ d, while those measured for SNe 2005cf, 2011fe and 2012cg are
8.0$\pm$1.1 $\rm\AA$, 0.8$\pm$0.2 $\rm\AA$ and 1.0$\pm$0.6 $\rm\AA$,
respectively. No C ii absorption feature is detected in the spectra of SN
2013dy at similar phase. The corresponding pEW of O i $\lambda$7774 is
measured as 52.0$\pm$6.3 $\rm\AA$ for SN 2017hpa at this epoch, comparable to
that of SN 2011fe (48.2$\pm$0.4 $\rm\AA$), while those measured for SNe 2012cg
and 2013dy are 16.2$\pm$1.4 $\rm\AA$ and 16.3$\pm$3.4 $\rm\AA$, respectively.
No prominent O i $\lambda$7774 is detected in the spectra of SN 2005cf at
similar epoch, which is consistent with the findings by Wang et al. (2009).
Following the discovery by Zhao et al. (2016) that the velocity of O i
$\lambda$7774 line shows a positive correlation with that of C ii
$\lambda$6580\. We propose that more unburned carbon and oxygen may be kept in
the explosion ejecta of SNe 2017hpa, 2011fe, and 2012cg, although the stronger
O i $\lambda$7774 absorption observed in SN 2017hpa could result from higher
oxygen abundance of the exploding white dwarf (Cui et al., 2020). A detached
Ca ii NIR HVF could be detected through multi-gaussian fitting proposed by
Zhao et al. (2015, 2016).

Figure 11(b) shows the comparison at $\sim 1$ week before maximum light. All
spectra show an increase in absorption strength of IMEs. The Si ii
$\lambda$6355 velocity of SN 2017hpa derived from absorption minimum is
$12,500\pm 180$ km s-1, which is comparable to that of the comparison sample.
The C ii $\lambda$6580 absorption line remained visible in the red wing of Si
ii $\lambda$6355 at this epoch.

The spectra near maximum light are displayed in Figure 11(c). The absorption
features due to IMEs such as Si ii at 4130 $\rm\AA$, Si iii at 4560 $\rm\AA$,
and S ii at 5468, 5612, and 5654 $\rm\AA$ become prominent at this phase. The
C ii absorption features at around 6300 and 7000 $\rm\AA$ remain noticeable in
SN 2017hpa but they are barely seen in the comparison SNe Ia. The $R$(Si ii),
defined as the line-strength ratio of Si ii $\lambda$5972 to Si ii
$\lambda$6355 (Nugent et al., 1995), can be used as indicator of the
photospheric temperature. A lower value of $R$(Si ii) corresponds to a higher
photospheric temperature for the SNe Ia. At around maximum light, $R$(Si ii)
is measured to be $0.18\pm 0.03$, comparable to that of SN 2018oh ($R$(Si ii)
$=0.15\pm 0.04$), suggesting that these two SNe have similar photoshperic
temperature around maximum light. The relatively larger ratio indicates a
lower photospheric temperature for SN 2017hpa. The pseudo-equivalent widths
(pEWs) of Si ii $\lambda$5972 and Si ii $\lambda$6355 near maximum light are
measured to be $15.5\pm 0.6$ $\rm\AA$ and $83.9\pm 2.2$ $\rm\AA$,
respectively, putting SN 2017hpa into the CN subtype of Branch et al. (2006)
classification.

Figure 11(d) shows the spectral evolution at $t\approx 30$ d. The C ii
$\lambda$6580 absorption line disappeared in this phase in all of our objects.
With the receding of the photosphere, the Fe ii features gain strength and
gradually dominate at wavelengths between 4700 and 5000 $\rm\AA$. The
absorption profiles of SN 2017hpa and the comparison sample are well developed
and tend to have uniform morphologies.

###  4.2 Carbon Features

The presence of C ii absorption can be easily identified in the early-time
spectra of SN 2017hpa around 6300 and 7000 $\rm\AA$. The C ii absorption
features of SN 2017hpa are stronger than in the comparison SNe Ia. The left
panel of Figure 12 shows that the C ii $\lambda$6580 absorption lines persist
in the spectra from $t\approx-14.1$ to $-7.9$ d. This absorption feature
disappeared in the spectra approaching maximum light and then reemerged at
$t\approx 9.5$ d. As a possible explanation, we propose that the C ii will be
highly excited when the detonation front or the deflagration front propagates
outward through the ejecta of SNe Ia (Ciaraldi-Schoolmann et al., 2013;
Seitenzahl et al., 2013) and this will make the C ii absorption features
disappear temporarily. With the receding and cooling of the photosphere of SNe
Ia, the C ii absorption trough will reemerge in the spectra. The right panel
of Figure 12 shows the relatively weak C ii $\lambda$7234 absorption; it is
noticeable in the earliest four spectra, and it then became flattened.
Inspection of the spectra does not reveal significant absorption of C ii
$\lambda$4267 in SN 2017hpa. Both C ii $\lambda$6580 and C ii $\lambda$7234
became barely visible in spectra taken $\sim 10$ d after maximum light.

Figure 12: The left panel shows the C ii $\lambda$6580 temporal evolution
while the right panel shows that of C ii $\lambda$7234\. The dashed lines mark
the Doppler velocity range from $-18,000$ km s-1 to $-5000$ km s-1.

The SN Spectroscopic Evolution222https://mwvg-spec-
evolve.readthedocs.io/en/latest/ package is employed to fit the absorption
components of Si ii $\lambda$6355 and C ii $\lambda$6580\. For SN 2017hpa, the
C ii $\lambda$6580 velocity is found to range from $\sim 13,000$ km s-1 at
$t\approx-14.1$ d to $\sim 9,300$ km s-1 at $t\approx-7.9$ d. According to
Silverman & Filippenko (2012), the average velocity ratio between C ii
$\lambda$6580 and Si ii $\lambda$6355 is $\sim 1.05$ for SNe Ia with
observations at least four days prior to $B$-band maximum light. However, the
mean C ii $\lambda$6580 to Si ii $\lambda$6355 velocity ratio (hereafter $R$(C
ii/Si ii)) measured for SN 2017hpa is only $\sim 0.81$ (as shown in Figure
13), suggesting that significant unburned carbon may have mixed deep into the
ejecta.

Figure 13: Top panel: temporal evolution of the C ii $\lambda$ expansion
velocity for SNe Ia with carbon detections. Bottom panel: temporal evolution
of the velocity ratio C ii $\lambda$6580 to Si ii $\lambda$6355\. The
comparison data are taken from Silverman & Filippenko (2012).

###  4.3 Ejecta Velocity

The ejecta velocities measured from the absorption lines, such as S ii
$\lambda\lambda$5468,5640, Si ii $\lambda$6355, C ii $\lambda$6580, C ii
$\lambda$7234, and O i $\lambda$7774, are shown in Figure 15. The photospheric
velocity measured from Si ii $\lambda$6355 at $t\approx-13.5$ d is $\sim
16,000$ km s-1, which is comparable to that of the Ca ii NIR triplet ($\sim
16,200$ km s-1), but it is faster than the C ii velocity ($\sim 12,000$ km s-1
for both C ii $\lambda$6580 and C ii $\lambda$7234). The velocity of the C ii
$\lambda$6580 absorption is roughly within the typical expansion velocity of
normal SNe Ia (Silverman & Filippenko, 2012). At the time of $B$-band maximum,
the velocity of Si ii $\lambda$6355 is estimated to be $\sim$ $9550\pm 170$ km
s-1, which can put SN 2017hpa into the NV subclass in the Wang et al. (2009)
classification scheme (the basic paramaters of SN 2017hpa are listed in Table
5), as shown in Figure 14. However, the Si ii $\lambda$6355 velocity of SN
2017hpa seems to have a large gradient, $\sim$ $130\pm 7$ km s-1 d-1, measured
within about 10 d after maximum light.

Table 5: Parameters of SN 2017hpa Parameter | Value
---|---
Photometric
$B_{\rm max}$ |  $14.88\pm 0.02$ mag
$B_{\rm max}-V_{\rm max}$ |  $0.005\pm 0.007$ mag
$M_{\rm max}(B)$ |  $-19.12\pm 0.11$ mag
$E(B-V)_{\rm host}$ |  $0.06\pm 0.06$ mag
$\Delta m_{15}(B)$ |  $1.02\pm 0.07$ mag
$s_{BV}$ | $0.94\pm 0.03$
$t_{\rm max}(B)$ |  $58,066.64\pm 0.36$ d
$t_{0}$ |  $58,050.22\pm 1.20$ d
$\tau_{\rm rise}$ |  $18.26\pm 0.97$ d
$L_{\rm bol}^{\rm max}$ |  $1.25\times 10^{43}$ erg s-1
$M_{{}^{56}\rm Ni}$ | $0.63\pm 0.02\,M_{\odot}$
Spectroscopic
$v\rm_{0}$(Si ii) |  $9550\pm 170$ km s-1
$\dot{v}$(Si ii) |  $130\pm 7$ km s-1 d-1
$R$(Si ii) | $0.18\pm 0.03$
Figure 14: Velocity evolution of SN 2017hpa as derived from the absorption
minimum of Si ii $\lambda$6355, compared with SNe 2005cf, 2011fe, 2013dy, and
2018oh. The average velocity curves obtained for SN 1991T-like and SN 1991bg-
like SNe are overplotted in red and blue dashed lines, respectively. The
normal subclass of SNe Ia is plotted with a black solid line. The shaded
region represents the 1$\sigma$ uncertainty for the mean velocity curve of
normal SNe Ia. Data for the comparison SNe and the region of normal SNe Ia are
extracted from (Li et al., 2019a).  Figure 15: Velocity evolution of different
elements measured from spectra of SN 2017hpa.

###  4.4 High-Velocity Feature

At early phases, the HVFs of the Ca ii NIR triplet can be clearly recognized
from the coresponding absorption-line profiles in the spectra. We utilize a
multi-Gaussian function to fit the absorption profile of the Ca ii NIR triplet
following the method described by Zhao et al. (2015, 2016). For SN 2017hpa,
the HVF of the Ca ii NIR triplet seen in the $t\approx-13.5$ d spectrum has a
velocity of $\sim 24,000$ km s-1, comparable to that of SN 2011fe ($\sim\
22,000$ km s-1; Zhang et al., 2016). The Ca ii NIR HVFs exhibit a velocity
plateau of 20,000 km s-1 from $t\approx-10$ to $-2$ d, which is similarly seen
in SN 2018oh but at different epochs (Li et al., 2019a). Note that there are
no obvious Si ii HVFs in the early-phase spectra of SN 2017hpa. It is
suggested that HVFs are more commonly detected in line profiles of the Ca ii
NIR triplet than in Si ii (Maguire et al., 2012; Childress et al., 2014;
Maguire et al., 2014; Pan et al., 2015; Silverman et al., 2015). Most SNe Ia
are found to have strong Ca ii NIR HVFs in their spectra at $t<7$ d while no
more than 30% of them have strong Si ii HVFs (Zhao et al., 2015).

##  5 Discussion

###  5.1 Distance and Quasibolometric Light Curve

Applying a standard cosmological model and assuming H${}_{0}=73.5$ km s-1
Mpc-1, $\Omega_{M}=0.3$, and $\Omega_{\Lambda}=0.7$ (Riess et al., 2018), a
distance modulus of $\sim 34.05$ mag can be obtained for the host galaxy of SN
2017hpa. We also utilize the latest $EBV$ model of SNooPy2 to fit the light
curves of SN 2017hpa in several optical bands, and the best-fit result gives
an average distance modulus of $34.00\pm 0.09$ mag. These two distance moduli
agree well with each other within the uncertainties. Adopting the distance
modulus as $34.00\pm 0.09$ mag and assuming $R_{V}=3.1$, we derive the
absolute $B$-band peak magnitude to be $M_{\rm max}(B)$ = $-19.12\pm 0.11$ mag
after correcting for both Galactic and host-galaxy extinction. This value
agrees well with the typical value of normal SNe Ia ($M_{\rm
max}(B)\approx-19.3$ mag; Phillips et al., 1999; Wang et al., 2009).

Our extensive photometric observations are used to establish the
quasibolometric light curve of SN 2017hpa. The spectral energy distribution
(SED) includes flux contributions from the following bands: $uvw2$, $uvm2$,
$uvw1$, $B$, $g$, $V$, $R$, $r$, $I$, and $i$. We adopt the procedure used for
SN 2018oh to establish the SED at several epochs (Li et al., 2019a). The
observed magnitudes are dereddened and converted into flux density. The flux
densities are then integrated using Simpson’s rule (Rogers, 1920; Syam, 2003)
through the effective wavelengths.

To get better knowledge of the peak luminosity, we use the UV and optical
observations to construct the quasibolometric light curves by assuming the NIR
contribution to be 5% at maximum light (Leloudas et al., 2009; Wang et al.,
2009; Zhang et al., 2016; Zhai et al., 2016). Applying a polynomial fitting,
the maximum luminosity is estimated to be $L_{\rm peak}$ = $1.25\times
10^{43}$ erg s-1 at about 0.85 d prior to $B$-band maximum. This peak
luminosity is comparable to that of SN 2011fe ($\sim 1.13\times 10^{43}$ erg
s-1; Zhang et al., 2016) but lower than that of SN 2018oh ($\sim 1.49\times
10^{43}$ erg s-1; Li et al., 2019a).

The modified radiation diffusion model of Arnett (Arnett, 1982; Chatzopoulos
et al., 2012; Li et al., 2019a) is applied to evaluate the initial nickel mass
together with other physical parameters of the SN ejecta. The Minim Code
(Chatzopoulos et al., 2013) is used to fit the quasibolometric light curve
with a constant opacity approximation. The model input parameters are the
first-light time (FLT) $t_{0}$, the radioactive 56Ni ejecta mass $M_{\rm Ni}$,
the timescale $t_{\rm lc}$ of the light curve, and the leaking timescale of
gamma rays $t_{\gamma}$ (see, e.g., Chatzopoulos et al., 2012, 2013). We set
all of these parameters free when performing the model fitting. The final
best-fit result of the quasibolometric luminosity evolution of SN 2017hpa is
shown in Figure 16. Based on $\chi^{2}$ minimization, a set of parameters is
found: $t_{0}=-0.94\pm 1.06$ d, $t_{\rm lc}=\,$$15.86\pm 0.76$ d, $M_{\rm
Ni}$= $0.63\pm 0.02\,M_{\odot}$, and $t_{\gamma}=\,$$28.37\pm 3.44$ d. The
initial nickel mass is comparable to the estimates of $M_{\rm Ni}$$\approx
0.57$ $M_{\odot}$ for SN 2011fe (Zhang et al., 2016) and $0.55\pm 0.04$
$M_{\odot}$ for SN 2018oh (Li et al., 2019a), but smaller than that of
$0.77\pm 0.11$ $M_{\odot}$ for SN 2005cf (Wang et al., 2009) and $0.68\pm
0.14$ $M_{\odot}$ for SN 2003du (Stanishev et al., 2007).

Figure 16: The quasibolometric light curve (dots) with an Arnett (1982)
radiation diffusion model (blue curve).

Adopting the method used by Li et al. (2019a), the average opacity $\kappa$ is
estimated to be $0.36\pm 0.15$ cm2 g-1. With the best-fit $t_{\rm lc}$ and
$t_{\gamma}$, we then obtain the ejecta mass and kinetic energy as $M_{\rm
ej}$$=0.70\pm 0.22$ $M_{\odot}$ and $E_{\rm kin}$$=(0.70\pm 0.50)\times
10^{51}$ erg. These values are within the range of typical SNe Ia as suggested
by Scalzo et al. (2019).

###  5.2 High Velocity Gradient

The ejecta velocity (i.e., $9550\pm 170$ km s-1) measured for SN 2017hpa near
maximum light is comparable to that of normal SNe Ia. According to the
velocity gradient of the Si ii, SNe Ia can be divided into LVG, HVG, and FAINT
subtypes (Benetti et al., 2005). Most normal velocity (as opposed to HV) SNe
Ia tend to be LVG or FAINT objects (Silverman et al., 2012). The left panel of
Figure 17 shows the scatter plot of the ${\Delta}m_{15}(B)$ versus velocity
gradient of SNe Ia, and the right panel displays that of the velocity gradient
vs. velocity measured around the time of maximum light. It can be seen that SN
2017hpa should be classified in the HVG subcategory, contradicting the trend
that SNe Ia showing prominent carbon features tend to be LVG objects (Parrent
et al., 2011). According to previous studies, HV SNe Ia tend to have larger
velocity gradients and vice versa, while this tendency seems to be broken by
SN 2017hpa.

Figure 17: Spectroscopic subclassification of SN 2017hpa (as marked with black
dot) based on the scheme of Benetti et al. (2005). Left panel:
${\Delta}m_{15}(B)$ is plotted with respect to the velocity gradient which is
measured from Si ii $\lambda$6355\. The SNe from different subtypes are taken
from Benetti et al. (2005) and Chakradhari et al. (2018), the four
transitional objects are from Pastorello et al. (2007) and Sahu et al. (2013),
SN 2005cf is from Wang et al. (2009), and SN 2018oh is from Li et al. (2019a).
Right panel: the scatter plot of the velocity measured from Si ii
$\lambda$6355 near maximum light versus the velocity gradient. The velocities
are taken from Silverman et al. (2012) and Wang et al. (2019). The horizontal
dashed line in the left pannel marks the boundary between HVG and LVG, which
is 70 km s-1 d-1 Benetti et al. (2005).

Previous studies have shown that for the HVG and LVG subclasses, the
difference in velocity gradient may be due to the different nature of the
explosion or the mixing degree of heavy elements (Sahu et al., 2013). An off-
center ignition will trigger SNe to explode asymmetrically. In this case,
different viewing angles will cause the observed velocity gradient to vary
greatly (Maeda et al., 2010). It is suggested that varying the criterion for
deflagration-to-detonation transition (DDT; Woosley et al., 2009) in
explosions also can result in a wide range of velocity gradients (Blondin et
al., 2011). However, these two scenarios are not suited to SN 2017hpa, which
has a low velocity but high velocity gradient. Alternatively, the effective
mixing of heavy elements in the SN ejecta may lead to the high velocity
gradient of the HVG subclass, while the inefficient mixing may cause the low
velocity gradient of the LVG subclass (Blondin et al., 2012; Sahu et al.,
2013).

###  5.3 Prominent Carbon Features

Detection of unburned carbon is important for constraining the explosion
mechanisms or progenitor systems of SNe Ia, and different explosion models
predict the presence of unburned material in different regions of the ejecta
(Fink et al., 2010; Pakmor et al., 2012; Sim et al., 2012; Seitenzahl et al.,
2013; Shen et al., 2018; Li et al., 2021). The velocities of IMEs in DDT will
increase with the explosion strength, leading to the diffusion of unburned
materials farther outward (Fink et al., 2010; Blondin et al., 2011). The
carbon and oxygen have distinct velocity distributions in the delayed-
detonation model while a similar velocity distribution is predicted in the
violent merger (Röpke et al., 2012). For those carbon-positive SNe Ia such as
SNe 2005di, 2005el, 2005ki, and SNF20080514-002, the absorption notch due to C
ii $\lambda$6580 usually disappears about one week before maximum light
(Thomas et al., 2011). Recent studies suggest that some SNe Ia or peculiar SNe
Ia exhibit unusually persistent carbon features in their spectra, such as SN
2002fk (Cartier et al., 2014), iPTF14atg (Cao et al., 2015), and SN 2018oh (Li
et al., 2019a). SN 2017hpa is another example showing such a prominent C ii
$\lambda$6580 feature at early epoches ($t\leq\ -7.9$ d), and this carbon
feature disappeared in near-maximum-light spectra. However, unlike other SNe
Ia, the carbon feature seems to reemerge at phases from $t\,\sim\,$+8.7 d to
$\sim\,$+11.7 d after maximum light.

The velocity ratio of C ii $\lambda$6580 to Si ii $\lambda$6355 is an
important parameter for setting constraints on the explosion models of SNe Ia
(Parrent et al., 2011; Folatelli et al., 2012). As shown in Figure 13, the
typical value of such a velocity ratio is $\sim 1.05$ for SNe Ia (Silverman &
Filippenko, 2012). However, the mean $R$(C ii/Si ii) is measured to be 0.81
for SN 2017hpa, much lower than the typical value. As noted by Silverman et
al. (2012), for a given object, the $R$(C ii/Si ii) usually increases somewhat
with time. This conclusion may be also supported by the data presented by
Parrent et al. (2011). Scalzo et al. (2010) suggested that the prominent C ii
$\lambda$6580 feature, concurrent with low velocities, could be associated
with a pre-explosion envelope of progenitor material originated from the
merger of two white dwarfs. With the receding of the photosphere of SNe Ia,
the ejecta of the inner layer with more uniform velocity distribution begin to
show up, which leads to the observed phenomenon that the $R$(C ii/Si ii)
increases slowly with time. The abundance distribution inferred from the
violent merger model indicates that both carbon and oxygen can be mixed deep
into the inner layer of the ejecta (Röpke et al., 2012). The prominent carbon
features and high velocity gradient may suggest that SN 2017hpa is reminiscent
of a low-luminosity subclass like SN 1991bg. However, the light curves and
color curves of SN 2017hpa are quite different from those of low-luminosity
objects like SN 2005bl (Taubenberger et al., 2008) or even transitional
objects like SN 2004eo (Pastorello et al., 2007).

To investigate the abnormal behavior of SN 2017hpa, we perform a further
comparison of C ii $\lambda$6580 absorption in Figure 18, where the sample
includes SNe 2002fk, 2005el, 2005cf, 2009dc, 2011fe, 2012cg, 2013dy, and
2018oh. The spectra of the comparison SNe are taken from the (Wang et al.,
2009; Yaron & Gal-Yam, 2012; Silverman et al., 2012; Zhai et al., 2016; Zhang
et al., 2016; Guillochon et al., 2017; Li et al., 2019a; Stahl et al., 2020).
In spectra of SN 2013dy, a strong C ii $\lambda$6580 absorption feature can be
found with a velocity up to $\sim 16,300$ km s-1 at early epochs, but this
absorption feature quickly fades away at $t\approx-12.9$ d, $\sim 3$ d after
explosion (Zheng et al., 2013; Pan et al., 2015). SN 2012cg shows moderately
strong C ii $\lambda$6580 at the same phase with respect to SN 2017hpa, and
the C ii absorption feature lasts until $\sim 8$ d before maximum light
(Silverman et al., 2012). The spectra of SN 2009dc exhibit very prominent C ii
$\lambda$6580 absorption that lasts for a long time, and this SN is proposed
to result from a super-Chandrasekhar mass progenitor system (Howell et al.,
2006; Silverman et al., 2011; Taubenberger et al., 2011; Tanaka et al., 2010).
SNe 2005cf, 2005el, and 2011fe show moderate C ii features along their spectra
evolution, while SNe 2002fk and 2018oh have C ii absorption comparable to that
of SN 2017hpa. All three of these normal SNe Ia show prominent C ii absorption
feature. The C ii absorption lines could be detected even in the spectra at
$\sim$7 days after $B$-band maximum light.

Figure 18: The C ii $\lambda$6580 evolution of SN 2017hpa compared to some
well-observed SNe Ia, including SN 2002fk, 2005el, 2005cf, 2009dc, 2011fe,
2012cg, 2013dy, and 2018oh.

Previous studies suggest that carbon-positive SNe Ia tend to have bluer
optical or UV colors (Thomas et al., 2011; Silverman et al., 2012; Milne et
al., 2013). Swift/UVOT observations also suggest that SNe Ia with prominent
carbon features are NUV-blue objects (Roming et al., 2005; Milne et al., 2010;
Thomas et al., 2011). The only exception is SN 2005cf, which belongs to the
NUV-red subgroup but with signatures of carbon features (Silverman et al.,
2012; Milne et al., 2013). SN 2017hpa is also a NUV-red SN Ia and has carbon
in its early-time spectra. Based on model comparisons, Brown et al. (2019)
suggested that the physical origin of the NUV-blue and NUV-red subclasses are
likely related to metal abundance. As suggested by Heringer et al. (2017), the
C ii absorption features could be hidden by the emission from iron, leading to
a lower metallicity within the outer layers of SNe Ia with carbon signatures.
However, if metallicity is the dominant origin of the NUV differences and the
presence of C ii features, a continuous distribution is expected in each of
them (Brown et al., 2019). A large sample of SNe Ia with positive C ii
features is needed for modeling metallicity effects on C ii absorption in the
spectra.

Based on the above discussions, SN 2017hpa shows prominent carbon features
with distinct evolution, a low C ii $\lambda$6580 to Si ii $\lambda$6355
velocity ratio, and normal ejecta velocity but a high velocity gradient, all
of which are unusual for known subtypes of normal SNe Ia. We suggest that SN
2017hpa could result from a violent merger explosion of two carbon-oxygen
white dwarfs, which brings up the prominent and distinct C ii features in its
spectra. The deep mixing of the SN ejecta may result in the high velocity
gradient of SN 2017hpa.

##  6 Conclusion

In this paper, we present extensive optical photometry and spectroscopy of the
Type Ia SN 2017hpa, which was discovered at a relatively young phase. This
object can be put into the category of normal and NUV-red SNe Ia, with
${\Delta}m_{15}(B)$= $1.02\pm 0.07$ mag and an absolute $B$-band magnitude
$M_{\rm max}(B)$ = $-19.12\pm 0.11$ mag.

The quasibolometric light curve of SN 2017hpa is established by using
extensive UV/optical photometric observations. Arnett’s 56Ni and 56Co
radioactive-decay-driven radiation diffusion model is utilized to fit the
quasibolometric light curve, deriving a peak luminosity of SN 2017hpa as
$L_{\rm peak}$ = $1.25\times 10^{43}$ erg s-1. The mass of nickel synthesized
during the explosion is estimated to be $M_{\rm Ni}$= $0.63\pm
0.02\,M_{\odot}$, and the ejecta mass $M_{\rm ej}$$=0.70\pm 0.22$ $M_{\odot}$.

The spectral evolution of SN 2017hpa is roughly the same as that of normal SN
Ia such as SN 2018oh. However, prominent C ii absorption and abnormal velocity
evolution distinguish it from other normal SNe Ia. The carbon and oxygen
features appear stronger than in normal SNe Ia and lasted until about 10 d
after maximum light, and both the carbon and oxygen have a lower velocity than
intermediate-mass elements such as Si ii and Ca ii. Although SN 2017hpa has a
typical ejecta velocity, $\sim$ $9550\pm 170$ km s-1 as measured near the
maximum light, it has an unsually large velocity gradient ($\sim$ $130\pm 7$
km s-1 d-1) in comparison with other normal SNe Ia. The significant amount of
unburned C and O in the ejecta, lower velocity relative to IMEs, and large
velocity gradient are more consistent with the merger model. More observations
and detailed modeling are needed to reveal the exact explosion physics in
objects like SN 2017hpa.

## Acknowledgments

We thank the anonymous referee for the suggestive comments, which improved the
manuscript. Funding for this work was provided by the National Natural Science
Foundation of China (NSFC, grants 11873081, U2031209, 12033002, 11633002, and
11761141001) and the National Program on Key Research and Development Project
(grant 2016YFA0400803), the High Level Talent-Heaven Lake Program of Xinjiang
Uygur Autonomous Region of China. This work is partially supported by the
Scholar Program of Beijing Academy of Science and Technology (DZ:BS202002). We
acknowledge the staffs of the Lijiang 2.4 m telescope (LJT), the Xinglong 2.16
m telescope (XLT), and Lick Observatory for their support. The Chinese Academy
of Sciences and the People’s Government of Yunnan Province provide support for
the LJT, which is corporately run and maintained by Yunnan Observatories and
the Center for Astronomical Mega-Science (CAS). JuJia Zhang is supported by
the National Natural Science Foundation of China (NSFC; grants 11773067,
11403096), the Youth Innovation Promotion Association of the CAS (grant
2018081), and the Ten Thousand Talents Program of Yunnan for Top-notch Young
Talents. Support for A.V.F.’s group at U.C. Berkeley was provided by the
TABASGO Foundation, the Christopher R. Redlich Fund, and the Miller Institute
for Basic Research in Science (U.C. Berkeley).

This work makes use of data from the Las Cumbres Observatory network. JB, DH,
DAH, and CP were supported by NSF grant AST-1911225. The Swift/UVOT data were
reduced by P.J. Brown and released in the Swift Optical/Ultraviolet Supernova
Archive (SOUSA), which is supported by NASA’s Astrophysics Data Analysis
Program (grant NNX13AF35G). Some of the observations with the Lick Observatory
1 m Nickel telescope were conducted by U.C. Berkeley undergraduate students
Sanyum Channa, Edward Falcon, Nachiket Girish, Romain Hardy, Julia Hestenes,
Andrew Hoffman, Evelyn Liu, Shaunak Modak, Costas Soler, Kevin Tang, Sameen
Yunus, and Keto Zhang; we thank them for their excellent work. Lick/KAIT and
its ongoing operation were made possible by donations from Sun Microsystems,
Inc., the Hewlett-Packard Company, AutoScope Corporation, Lick Observatory,
the U.S. National Science Foundation, the University of California, the Sylvia
& Jim Katzman Foundation, and the TABASGO Foundation. A major upgrade of the
Kast spectrograph on the Shane 3 m telescope at Lick Observatory was made
possible through generous gifts from the Heising-Simons Foundation as well as
William and Marina Kast. Research at Lick Observatory is partially supported
by a generous gift from Google.

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