arXiv:1803.01050v2 [astro-ph.SR] 9 May 2018
Draft version May 11, 2018
Typeset using L
A
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SN 2017ein and the Possible First Identification of a Type Ic Superno
va Progenitor
Schuyler D. Van Dyk,
1
WeiKang Zheng,
2
Thomas G. Brink,
2
Alexei V. Filippenko,
3, 4
Dan Milisavljevic,
5
Jennifer E. Andrews,
6
Nathan Smith,
6
Michele Cignoni,
7, 8
Ori D. Fox,
9
Patrick L. Kelly,
2,10
Angela Adamo,
11
Sameen Yunus,
2
Keto Zhang,
2
and Sahana Kumar
2,12
1
Caltech/IPAC, Mailcode 100-22, Pasadena, CA 91125, USA 000
0-0001-9038-9950
2
Department of Astronomy, University of California, Berkel
ey, CA 94720-3411, USA
3
Department of Astronomy, University of California, Berkel
ey, CA 94720-3411, USA 0000-0003-3460-0103
4
Miller Senior Fellow, Miller Institute for Basic Research i
n Science, University of California, Berkeley, CA 94720, US
A
5
Department of Physics and Astronomy, Purdue University, 52
5 Northwestern Avenue, West Lafayette, IN 47907, USA
6
Steward Observatory, University of Arizona, 933 N. Cherry A
venue, Tucson, AZ 85721, USA
7
Dipartimento di Fisica ’Enrico Fermi’, Universit ́a di Pisa
, largo Pontecorvo 3, I-56127 Pisa, Italy 0000-0001-6291-6
813
8
INFN, Largo B. Pontecorvo 3, I-56127 Pisa, Italy
9
Space Telescope Science Institute, 3700 San Martin Drive, B
altimore, MD 21218, USA 0000-0002-4924-444X
10
College of Science & Engineering, Minnesota Institute for A
strophysics, University of Minnesota, 115 Union St. SE, Min
neapolis, MN
55455 USA 0000-0003-3142-997X
11
Department of Astronomy, Oskar Klein Centre, Stockholm Uni
versity, AlbaNova University Centre, SE-106 91 Stockholm,
Sweden
0000-0002-8192-8091
12
Department of Physics, Florida State University, 77 Chieft
ain Way, Tallahassee, Florida 32306, USA
ABSTRACT
We have identified a progenitor candidate in archival
Hubble Space Telescope
(
HST
) images for the
Type Ic SN 2017ein in NGC 3938, pinpointing the candidate’s location via
HST
Target-of-Opportunity
imaging of the SN itself. This would be the first identification of a stellar
-like object as a progenitor
candidate for any Type Ic supernova to date. We also present obs
ervations of SN 2017ein during
the first
∼
49 days since explosion. We find that SN 2017ein most resembles the w
ell-studied Type
Ic SN 2007gr. We infer that SN 2017ein experienced a total visual e
xtinction of
A
V
≈
1
.
0–1.9 mag,
predominantly because of dust within the host galaxy. Although the
distance is not well known, if
this object is the progenitor, it was likely of high initial mass,
∼
47–48
M
⊙
if a single star, or
∼
60–80
M
⊙
if in a binary system. However, we also find that the progenitor cand
idate could be a very blue
and young compact cluster, further implying a very massive (
>
65
M
⊙
) progenitor. Furthermore, the
actual progenitor might not be associated with the candidate at all
and could be far less massive. From
the immediate stellar environment, we find possible evidence for thre
e different populations; if the SN
progenitor was a member of the youngest population, this would be c
onsistent with an initial mass of
∼
57
M
⊙
. After it has faded, the SN should be reobserved at high spatial re
solution and sensitivity, to
determine whether the candidate is indeed the progenitor.
Keywords:
supernovae: individual (SN 2017ein), stars: massive, binaries: ge
neral, galaxies: stellar
content, galaxies: individual (NGC 3938)
1.
INTRODUCTION
Supernovae (SNe) are among the most powerful ex-
plosions in the Universe and highly influential within
their host galaxies throughout cosmic time. Under-
standing their origins, as the catastrophic endpoints
of stellar evolution, is therefore an important line of
inquiry in modern astrophysics. Stars more massive
than
∼
8–10
M
⊙
perish as core-collapse SNe (CCSNe).
The most common CCSNe in the local Universe are
the Type II-Plateau (SNe II-P), and we now have solid
evidence that these arise from stars in the red super-
giant phase (e.g.,
Smartt et al. 2009
). The situation is
less clear for SNe whose progenitor has had its H-rich
envelope partially or entirely stripped away before ex-
plosion, the so-called stripped-envelope SNe (SESNe;
e.g.,
Filippenko 1997
). We now have accumulating
evidence that the progenitors of Type IIb SNe (SNe
IIb), stripped but still retaining
.
0
.
1
M
⊙
of hydro-
gen, appear to be yellow to somewhat blue supergiants
(e.g.,
Maund et al. 2011
;
Van Dyk et al. 2011
,
2014
;
2
Folatelli et al. 2015
;
Bersten et al. 2018
), likely in inter-
acting binary systems of moderate mass (
∼
12–15
M
⊙
;
e.g.,
Podsiadlowski et al. 1993
;
Stancliffe & Eldridge
2009
;
Benvenuto et al. 2013
).
For CCSNe which are hydrogen-free, there is only
one identified progenitor of a SN Ib, for iPTF13bvn
(
Cao et al. 2013
), although the actual nature of the
star is still to be determined (
Folatelli et al. 2016
;
Eldridge & Maund 2016
). However, for the H-free SNe
with little or no He as well, the SNe Ic, the progeni-
tors to date have been elusive. Here we are considering
normal SNe Ic; more extreme broad-lined examples of
SNe Ic have been found associated with long-duration
gamma-ray bursts (e.g.,
Hjorth & Bloom 2012
). SN pro-
genitor stars with binary companions that remain bound
after explosion can evolve into many exotic configura-
tions, including neutron star–neutron star mergers that
are associated with short-duration gamma-ray bursts,
kilonovae, and gravitational waves (e.g.,
Abbott et al.
2017
).
To remove most or all of the He from the progen-
itor, some mechanism for extensive mass loss is re-
quired.
Hachinger et al.
(
2012
) showed that it is in-
deed hard to hide H and He in such SNe, so the lack
of such spectral features implies the presence of efficient
mass loss. One way is through a strong stellar wind
from a single, massive star; for example,
Georgy et al.
(
2009
) concluded that SNe Ic would result from the ex-
plosion of a highly massive, but stripped, star in the
WC or WO Wolf-Rayet (WR) phase.
Dessart et al.
(
2012
) suggested that the lack of observed He
i
lines
in SN Ic spectra could arise from a high-mass, possi-
bly single, progenitor. The other mass-loss mechanism
is via envelope stripping in a mass-transfer binary sys-
tem.
Nomoto et al.
(
1990
) modeled the SN Ic 1987M
progenitor as a low-mass (
∼
3–3.5
M
⊙
) He star, with
initially 12–15
M
⊙
, in a close interacting binary sys-
tem. A low-mass binary model was also invoked to ex-
plain the SN Ic 1994I (
Nomoto et al. 1994
).
Yoon et al.
(
2010
) found a bimodality in their model SN Ic pro-
genitors at solar metallicity, with lower-luminosity SNe
(e.g., SN 1994I) from lower-mass (
M
ZAMS
≈
12–13
M
⊙
)
binaries and a majority of SNe Ic from systems with
higher-mass (
M
ZAMS
&
33
M
⊙
) primaries with WR
wind mass loss. A similar progenitor system dichotomy
had been found earlier by
Wellstein & Langer
(
1999
)
and
Pols & Dewi
(
2002
). Many of the progenitors of
SESNe must be lower-mass binaries, since the fraction
of such SNe is locally too high to have single-star pro-
genitors (
Smith et al. 2011
; see also
Graur et al. 2017
,
e.g., their Figure 10). See also
Zapartas et al.
(
2017
) for
a discussion of pathways to SESNe.
The lack of detected progenitors for SNe Ic is not
for a lack of attempting to locate them. Over the
last two decades, a valiant effort has been expended
by a number of investigators toward detecting SN
Ic progenitors. Their identification has been partic-
ularly thwarted by their proximity to luminous star
clusters (e.g., SNe 2004gt and 2013dk, both in the
Antennae;
Gal-Yam et al. 2005
;
Maund et al. 2005
;
Elias-Rosa et al. 2013
) or high extinction (
Eldridge et al.
2013
, who considered a number of SNe Ic), or both ef-
fects. Other influences on our ability to detect SN Ic
progenitors include the available archival pre-SN imag-
ing (e.g.,
Eldridge et al. 2013
); moreover, massive WR
stars, although bolometrically highly luminous, become
optically less luminous toward the end of their lives
(
Yoon et al. 2012
). Additionally, archival images are of-
ten not obtained in filters sensitive to the strong, broad
emission lines of WR stars, particularly He
ii
λ
4686, at
which WR stars are brighter than their continua by up
to 3 mag (
Massey & Johnson 1998
; see also Figure 7 of
Shara et al. 2013
).
The existing observational limits on SN Ic progenitors
provide some informative constraints. Based on
HST
data,
Gal-Yam et al.
(
2005
) placed limits of
M
V
>
−
5
.
5
and
M
B
>
−
6
.
5 mag on the SN 2004gt progenitor
and eliminated analogs of more than half of the known
Galactic WR stars as possible progenitors. Similarly,
Maund et al.
(
2005
), from the same dataset, limited
the SN 2004gt progenitor to a low-luminosity, high-
temperature star, such as a single high initial mass
(
>
40
M
⊙
), carbon-rich WC star.
Elias-Rosa et al.
(
2013
) set a limit of
M
F555W
&
−
5
.
7 mag on the SN
2013dk progenitor luminosity and could not rule out
WR stars. The observational limits in the case of either
SN did not provide constraints on lower-mass binary
systems. However, more recently,
Johnson et al.
(
2017
)
placed astoundingly low 1
σ
upper limits of
M
U
>
−
3
.
8,
M
B
>
−
3
.
1,
M
V
>
−
3
.
8, and
M
R
>
−
4
.
0 mag on the
presence of the SN Ic 2012fh progenitor in deep ground-
based images of the host galaxy, NGC 3344, ruling out
essentially all single-star models.
In this paper we describe what could be the first-ever
identification and characterization of the progenitor of
a SN Ic, the nearby SN 2017ein in NGC 3938. We first
present early-time data on the SN itself and compare
its properties to those of other well-studied SNe Ic, in
particular SN 2007gr. We then show our detection of
a candidate for the progenitor in archival
Hubble Space
Telescope
(
HST
) images, which we isolated with high-
spatial-resolution
HST
images of the SN; we initially
reported this identification in
Van Dyk et al.
(
2017
). Fi-
nally, we attempt to constrain the nature of the detected
3
E
N
6
5
4
3
2
SN
1
SN 2017ein
Figure 1.
KAIT unfiltered image, obtained on 2017 June
26.207, of SN 2017ein and the surrounding 6
.
′
4
×
6
.
′
4 field.
The SN position is indicated, along with several stars that
were used as local calibrators (see Table
1
) for the KAIT and
Nickel SN photometry.
Table 1.
Photometric Sequence for SN 2017ein
a
Star
B
V
R
I
(mag)
(mag)
(mag)
(mag)
1 14.383(034) 13.842(014) 13.519(016) 13.151(018)
2 15.844(034) 15.065(015) 14.613(017) 14.195(019)
3 17.721(034) 16.936(014) 16.481(016) 16.044(018)
4 18.403(035) 17.961(017) 17.691(019) 17.385(021)
5 18.207(035) 17.477(016) 17.052(018) 16.608(020)
6 18.093(033) 17.003(013) 16.380(016) 15.818(018)
a
Uncertainties are provided in parentheses in thousandths o
f a
magnitude.
progenitor and compare these properties to predictions
from the recent models of SESN progenitors.
SN 2017ein was discovered optically by
Arbour
(
2017
)
on May 25.99 (UT dates are used throughout this paper)
at
∼
17
.
6 mag, and confirmed by
Gaia
on May 29.71 at
G
≈
16
.
9 mag and given the designation Gaia17bjw
1
.
The SN was classified from a spectrum obtained on
May 26.6 as Type Ic, although initially as broad-lined
1
https://wis-tns.weizmann.ac.il/object/2017ein.
within one week of maximum light (
Xiang et al. 2017
).
Im et al.
(
2017
), through their regular monitoring of the
host galaxy, were able to determine that the discovery
by Arbour must have been made shortly after explo-
sion, since the SN was detectable in their images from
May 25, but not from May 24.
Im et al.
found that the
SN was continuing to rise in brightness in early June;
thus, the initial classification spectrum was most likely
not obtained near maximum light. We note that NGC
3938 was also host to the SN II-L 1961U (
Bertola 1963
),
the SN Ic 1964L (
Bertola et al. 1965
), and the SN II-P
2005ay (
Tsvetkov et al. 2006
;
Gal-Yam et al. 2008
).
2.
OBSERVATIONS
2.1.
Ground-Based Imaging
SN 2017ein was constantly monitored photometrically
in the bands
BV R
C
I
C
with both the 0.76-m Katzman
Automatic Imaging Telescope (KAIT;
Filippenko et al.
2001
) and the 1-m Nickel telescope at Lick Observa-
tory, as well as with unfiltered KAIT images, beginning
within 3 days of discovery and interrupted only by nights
with poor observing conditions. We followed the SN un-
til it was no longer accessible from either KAIT or the
Nickel telescope. We show a KAIT image of the field in
Figure
1
, with the SN indicated. From our KAIT data
we measured an absolute position for the SN of 11h 52m
53.26s +44
◦
07
′
26
.
′′
2 (J2000). All images were reduced
using a custom pipeline (
Ganeshalingam et al. 2010
).
Point-spread function (PSF) photometry was then ob-
tained using DAOPHOT (
Stetson 1987
) from the IDL
Astronomy Users Library
2
. We also tried to perform the
photometry after using the image template-subtraction
method, in order to estimate the host contribution to
the light. We find that the differences between using
subtraction and not using subtraction are very small
(
<
0
.
1 mag) for all epochs, implying that the host contri-
bution is minimal. We thus adopted the results without
template subtraction.
Six stars in the KAIT and Nickel fields were cho-
sen as local calibrators; they are indicated in Figure
1
.
This field was observed by Pan-STARRS and included in
the release of the Mean Object Catalog
3
. We obtained
from that release the mean PSF photometry for the six
stars in the PS1 bands and used the transformations to
BV R
C
I
C
provided by
Tonry et al.
(
2012
). We list the
resulting magnitudes for the stars in Table
1
and use
these for calibration of the SN photometry.
2
http://idlastro.gsfc.nasa.gov/.
3
Available at https://archive.stsci.edu/panstarrs/.
4
Table 2.
KAIT and Nickel Photometry of SN 2017ein
a
MJD
B
V
R
Unfiltered
I
Source
(mag)
(mag)
(mag)
(mag)
(mag)
57902.27 17.18(16) 16.49(19) 16.19(20)
· · ·
15.92(22) KAIT
57903.23 16.81(18) 16.23(17) 15.95(20) 15.89(17) 15.68(1
9) KAIT
57903.23 16.72(01) 16.23(06) 15.89(07)
· · ·
15.63(07) Nickel
57905.24 16.42(17) 15.81(13) 15.44(12) 15.45(12) 15.16(1
4) KAIT
57906.22 16.18(16) 15.63(13) 15.33(14)
· · ·
15.01(15) KAIT
57907.26 16.03(19) 15.56(20) 15.29(20) 15.18(17) 14.91(2
2) KAIT
57908.25 16.04(18) 15.39(17) 15.21(20) 15.11(20) 14.77(2
0) KAIT
57909.22 15.95(17) 15.36(17) 15.05(18)
· · ·
14.73(18) KAIT
57910.20 15.84(08) 15.27(07) 14.96(12)
· · ·
14.63(10) Nickel
57910.23 16.01(20) 15.26(17) 14.98(18) 14.93(21) 14.59(1
9) KAIT
57917.21 16.51(27) 15.32(21) 14.87(21)
· · ·
14.52(22) KAIT
57918.21 16.60(18) 15.36(19) 14.89(20) 15.02(23) 14.50(2
1) KAIT
57919.21 16.62(18) 15.41(15) 14.92(16) 15.01(17) 14.50(1
8) KAIT
57920.21 16.82(16) 15.49(16) 14.96(17) 15.06(18) 14.50(2
0) KAIT
57922.21 16.98(14) 15.64(14) 15.08(16) 15.17(11) 14.57(1
7) KAIT
57923.21 17.17(19) 15.68(18) 15.15(20) 15.26(18) 14.62(2
2) KAIT
57924.21 17.20(16) 15.79(14) 15.21(16) 15.31(22) 14.68(1
6) KAIT
57924.23 17.10(08) 15.79(12) 15.21(11)
· · ·
14.70(14) Nickel
57925.21 17.24(14) 15.90(13) 15.27(15) 15.34(11) 14.73(1
5) KAIT
57925.21 17.20(09) 15.90(11) 15.28(11)
· · ·
14.76(12) Nickel
57926.23 17.40(17) 15.96(17) 15.37(18)
· · ·
14.80(19) KAIT
57927.22 17.49(17) 16.08(17) 15.46(20) 15.50(16) 14.84(1
9) KAIT
57928.21 17.47(08) 16.13(09) 15.48(10)
· · ·
14.90(11) Nickel
57928.22 17.56(17) 16.16(18) 15.50(18) 15.55(18) 14.88(2
0) KAIT
57929.19 17.62(24) 16.13(18) 15.56(17) 15.58(17) 14.96(1
9) KAIT
57930.21 17.76(17) 16.26(15) 15.65(14) 15.67(15) 14.97(1
4) KAIT
57931.23 17.78(20) 16.34(17) 15.61(19) 15.68(16) 14.99(1
9) KAIT
57932.21 17.82(25) 16.41(21) 15.74(22) 15.87(31) 15.12(2
2) KAIT
57933.21 17.83(16) 16.46(13) 15.81(12) 15.82(15) 15.13(1
3) KAIT
57933.21 17.87(04) 16.51(02) 15.77(05)
· · ·
15.12(06) Nickel
57934.19
· · ·
· · ·
· · ·
15.84(17)
· · ·
KAIT
57938.19 18.13(27) 16.76(16) 16.08(15) 16.09(15) 15.32(1
5) KAIT
57939.20 18.13(19) 16.79(14) 16.15(15) 16.13(16) 15.39(1
5) KAIT
57940.19 18.32(42) 16.78(19) 16.15(19) 16.24(21) 15.42(2
0) KAIT
57941.20 18.13(19) 16.91(14) 16.23(15) 16.14(16) 15.47(1
5) KAIT
57942.19 18.10(35) 16.86(16) 16.21(13) 16.20(14) 15.49(1
6) KAIT
57943.19 18.26(31) 16.87(17) 16.31(16) 16.26(17) 15.46(2
1) KAIT
57944.19 18.43(38) 16.92(18) 16.30(19) 16.24(16) 15.58(2
1) KAIT
57945.19 17.98(33) 16.83(19) 16.26(14) 16.26(18) 15.47(1
8) KAIT
57947.20 18.22(23) 16.98(14) 16.30(12) 16.29(12) 15.51(1
5) KAIT
57947.20 18.15(08) 17.07(09) 16.33(09)
· · ·
15.55(09) Nickel
a
Uncertainties are provided in parentheses in hundredths of
a magnitude.
5
Apparent magnitudes of the SN and the six calibra-
tors were all measured in the KAIT4/Nickel2 natural
system. The final results were transformed to the stan-
dard system, using the local calibrators and color terms
for KAIT4 (
Ganeshalingam et al. 2010
; their Table 4)
and updated Nickel color terms (
Shivvers et al. 2017
).
We present the early-time KAIT and Nickel photome-
try of SN 2017ein in Table
2
.
2.2.
Spectroscopy
Over a two-month period beginning on 2017 June 2,
eight optical spectra of SN 2017ein were obtained with
the Kast Spectrograph mounted on the 3-m Shane tele-
scope (
Miller & Stone 1993
) at Lick Observatory. They
were taken at or near the parallactic angle (
Filippenko
1982
) to minimize slit losses caused by atmospheric dis-
persion. Data were reduced following standard tech-
niques for CCD processing and spectrum extraction
(
Silverman et al. 2012
) utilizing IRAF
4
routines and
custom Python and IDL codes
5
. Low-order polyno-
mial fits to comparison-lamp spectra were used to cal-
ibrate the wavelength scale, and small adjustments de-
rived from night-sky lines in the target frames were ap-
plied. Observations of appropriate spectrophotometric
standard stars were used to flux calibrate the spectra.
We obtained 3
×
1200 s exposures with the Blue Chan-
nel spectrograph on the MMT on 2017 June 24 (JD
2,457,928.69). The data were taken with the 1200 lines
mm
−
1
grating with a central wavelength of 6360
̊
A and a
1
.
′′
0 slit width. The seeing was 1
.
′′
2. Standard reductions
were carried out using IRAF, and wavelength solutions
were determined using internal He-Ne-Ar lamps. Flux
calibration was achieved using spectrophotometric stan-
dards at a similar airmass taken throughout the night.
A log of all spectroscopic observations is provided in Ta-
ble
3
.
We note that we cross-checked the calibrations of our
photometry and our spectra at four nearly contempora-
neous epochs (MJD 57906.76, 57925.77, 57931.76, and
57935.74) by comparing observed colors with colors syn-
thetically generated from the spectra using pysynphot
6
,
and found that the two datasets differed by at most
∼
0
.
1 mag, within the uncertainties of the photometry;
thus, we are confident that our calibrations are sufficient.
2.3.
HST
Imaging
4
IRAF is distributed by the National Optical Astronomy Ob-
servatory, which is operated by AURA, Inc., under a cooperat
ive
agreement with the NSF.
5
https://github.com/ishivvers/TheKastShiv.
6
https://github.com/spacetelescope/pysynphot.
The SN 2017ein site is in publicly available archival
images obtained with the
HST
Wide Field Planetary
Camera 2 (WFPC2) on 2006 October 20 (by program
GO-10877, PI A. Filippenko, originally to observe SN
2005ay) in bands F555W (total exposure time 460 s)
and F814W (700 s). The image mosaics were obtained
from the Hubble Legacy Archive, whereas the individual
WFPC2 frames were obtained from the
HST
Archive at
MAST.
The SN site is just off the edge of the field of view
(NIC3) of archival Near-Infrared Camera and Multi-
Object Spectrograph (NICMOS) images obtained on
2003 April 3 in bands F160W, F187N, and F190N, and
thus we do not consider these.
We also observed the transient itself with
HST
on 2017
June 12 using the Wide-Field Camera 3 (WFC3) UVIS
channel in subarray mode, as part of our Target of Op-
portunity (ToO) program (GO-14645, PI S. Van Dyk),
in F438W, with 10 s individual frame times and total
exposure time 270 s. The F438W band was chosen in
this case to attempt to avoid saturation by the SN, at
which we were successful. All of these data had been ini-
tially processed via the default pipelines at STScI and
were obtained from the
HST
Archive. The individual
flc
frames, which had been corrected by the pipeline for
charge-transfer efficiency losses, were individually com-
bined into an image mosaic using AstroDrizzle within
PyRAF.
3.
THE SUPERNOVA: EARLY RESULTS
3.1.
Photometry
We show the combined KAIT and Nickel light curves
in Figure
2
. Here we have also included the photom-
etry at
R
from
Im et al.
(
2017
). We compared the
light curves with those of a number of SNe Ic and
found that the best match was with the light curves
from
Hunter et al.
(
2009
) for the normal SN Ic 2007gr
in NGC 1058. The light curves for SNe 2017ein and
2007gr show quite similar behavior. We also compared
the SN 2017ein light curves with those of SN 2004aw
(
Taubenberger et al. 2006
), which may agree at early
times but diverge post-peak, with SN 2004aw being gen-
erally more luminous at all bands. (We have made the
comparison with SN 2004aw here, since, as we show in
Section
3.5
, SNe 2017ein and 2004aw could be similar
bolometrically.) One can see that we missed observa-
tions of maximum light at
V
[
V
(max)] for SN 2017ein,
owing to poor observing conditions. Based on the com-
parison with SN 2007gr, we estimate, however, that
V
(max) occurred on JD
≈
2,457,913.1 (i.e., June 8.6).
That the SN was not detected to
R >
18
.
4 mag in
images obtained by
Im et al.
(
2017
) on May 24.63 (JD
6
Table 3.
Log of Spectroscopic Observations of SN 2017ein
UT date
MJD Age
a
Instrument Wavelength Resolution
Range (
̊
A)
(
̊
A)
2017 Jun 02.26 57906.76
−
6.3
Kast
3640–10,630
2.0
2017 Jun 21.27 57925.77 12.7
Kast
3626–10,710
2.0
2017 Jun 24.19 57928.69 15.6
MMT
5711–7022
0.5
2017 Jun 27.26 57931.76 18.7
Kast
3630–10,712
2.0
2017 Jul 01.24 57935.74 22.6
Kast
3636–10,710
2.0
2017 Jul 17.21 57951.71 38.6
Kast
3612–10,700
2.0
2017 Jul 26.20 57960.70 47.6
Kast
3622–10,670
2.0
2017 Jul 30.20 57964.70 51.6
Kast
3620–10,704
2.0
2017 Aug 01.19 57966.69 53.6
Kast
3620–10,680
2.0
a
Day since estimated time of
V
maximum, approximately 2,457,913.1 (June
8.6).
2,457,898.13), but then was seen at
R
= 17
.
7 mag by
these investigators on May 25.77 (JD 2,457,899.27),
and later discovered by Arbour on May 25.99, would
imply that the explosion date was likely sometime
between JD 2,457,898 and 2,457,899. We analyzed
the
R
-band light curve, including the
Im et al.
(
2017
)
points, by fitting a simple analytic model for H-free
SNe (
Vacca & Leibundgut 1997
; see Figure
2
), and
found that the explosion could have been as early as JD
2,457,898.1, or May 24.6, which we adopt. The analytic
model further implies that
R
-band maximum (again,
missed by our observations) occurred on approximately
JD 2,457,916, or
∼
3 days after
V
-band maximum.
In Figure
3
we show the early-time
B
−
V
,
V
−
R
,
and
V
−
I
color curves. We have initially corrected the
observed color curves for Galactic foreground reddening
(
Schlafly & Finkbeiner 2011
; via the NASA/IPAC Ex-
tragalactic Database, NED). Again, comparing with SN
2007gr (
Hunter et al. 2009
), we can see that the color
evolution of SN 2017ein follows much the same behavior,
to within the uncertainties. We also found that
V
−
R
of SN 2017ein evolved at early times in a fashion similar
to that of SN 2004fe (not shown;
Drout et al. 2011
).
3.2.
Spectroscopy
We show the early-time spectra of SN 2017ein in Fig-
ure
4
. The first spectrum was obtained on June 2, about
one week after discovery (and explosion). From the
light-curve comparison with SN 2007gr, we estimate the
ages of the spectra relative to
V
(max). In the figure
we also show a comparison with spectra of SN 2007gr
from
Valenti et al.
(
2008
) and
Chen et al.
(
2014
), and
of SN 2004aw from
Taubenberger et al.
(
2006
), at ap-
proximately the same ages. The spectra of SN 2004aw
and the
Valenti et al.
spectra of SN 2007gr were ob-
tained from WISeREP
7
(
Yaron & Gal-Yam 2012
), and
the
Chen et al.
SN 2007gr spectra were obtained from
the UC Berkeley Supernova Database
8
One can see that
the spectra of SN 2017ein resemble quite closely those
of SN 2007gr, as was also true for the light curves. SN
2017ein also somewhat resembles SN 2004aw, spectro-
scopically, although the latter may have had somewhat
broader lines. Many of the spectral features, which we
have indicated in the figure, are the same between SN
2017ein and SN 2007gr, and the line widths are compar-
atively narrow. SN 2017ein, like SN 2007gr, is clearly
not a broad-lined SN Ic.
In the pre-maximum spectrum, the features at 6350
̊
A
and at
∼
7000
̊
A, attributed to C
ii
λ
6580 and
λ
7235
in SN 2007gr (
Valenti et al. 2008
), are evident in SN
2017ein, and the Ca
ii
λλ
8542, 8662, 8498 near-infrared
(NIR) feature and, interestingly, the
∼
10
,
400
̊
A feature,
attributed primarily to C
i
, both make an early appear-
ance. In the later spectra, Ca
ii
, Ti
ii
, Mg
ii
λ
4481,
the well-resolved Fe
ii
λλ
4924, 5018, 5169, Sc
ii
, Na
i
,
O
i
λ
7774, and C
i
are also seen, as in SN 2007gr. For
both SN 2017ein and SN 2007gr at later times, weak
interstellar H
α
emission appears.
We measured (see Table
4
) the photospheric expansion
velocity of SN 2017ein from the Fe
ii
λλ
4924, 5018, 5169
and O
i
λ
7774 absorption lines, by fitting a Gaussian to
each of the lines with the routine
splot
in PyRAF. The
measurements from the Fe
ii
lines were averaged. We
7
https://wiserep.weizmann.ac.il/.
8
http://heracles.astro.berkeley.edu/sndb/.
(SNDB;
Silverman et al. 2012
).
7
Figure 2.
Early-time optical KAIT (solid squares) and
Nickel (open squares)
BV R
C
I
C
and KAIT unfiltered (open
circles) light curves of SN 2017ein. Also shown are
R
ob-
servations by
Im et al.
(
2017
; open diamonds). Addition-
ally, for comparison we show the light curves for the SNe
Ic 2007gr (open triangles;
Hunter et al. 2009
) and 2004aw
(crosses;
Taubenberger et al. 2006
). The curves for both SNe
were shifted in time to match approximately the
V
maximum
for SN 2017ein. After correction for reddening appropriate
for these two SNe, their light curves were then reddened by
the Galactic foreground for SN 2017ein and an additional
host contribution of
A
V
= 1
.
05 mag with
R
V
= 3
.
1. The
SN 2007gr light curves were further adjusted by a difference
in distance modulus of 1.32 mag, which is consistent with
the difference between the distances to the SN 2007gr and
SN 2017ein hosts, to within the uncertainties; the SN 2004aw
curves were adjusted by a distance modulus difference of 2.02
mag, which is a smaller difference than implied by the dis-
tance given in
Taubenberger et al.
The dotted lines indicate
the epochs of SN 2017ein spectroscopy.
show these velocities for both sets of lines in Figure
5
,
with a comparison to SN 2007gr (
Chen et al. 2014
), SN
2004aw (
Taubenberger et al. 2006
), and several SNe Ic
in the sample from
Liu et al.
(
2016
). The expansion
velocities of SN 2017ein are not initially (soon after ex-
plosion) as high at those of SN 2007gr and SN 2004aw,
and overall appear to be intermediate between these two
other SNe.
Unlike SN 2007gr, though, as the Ca
ii
NIR triplet in
SN 2017ein appeared to narrow in width with advancing
age, a notch of additional absorption appeared in the
red wing of the triplet feature. It is unclear to what to
attribute this feature; it could be the Ca
ii
λ
8664 line
becoming more apparent as the triplet feature narrowed,
or it could be the C
i
λ
8727 line, which is normally
Figure 3.
Early-time
B
−
V
,
V
−
R
, and
V
−
I
color curves
of SN 2017ein (solid squares), after initial correction for
red-
dening attributed to the Galactic foreground contribution
(
Schlafly & Finkbeiner 2011
). For comparison we show the
color curves of SN Ic 2007gr (open triangles;
Hunter et al.
2009
), corrected for Galactic foreground reddening and then
reddened and shifted in time to match the curves of SN
2017ein. Additionally, we show the best fit of the SN Ic
color template at
B
−
V
from
Stritzinger et al.
(
2018
; solid
curve). See discussion in Section
3.3
.
blended with the Ca
ii
NIR triplet, as the other C
i
absorption features also become stronger.
Also notably exceptional for SN 2017ein is the pres-
ence of significant Na
i
D absorption in the June 2 spec-
trum (day
−
6
.
3), implying more interstellar extinction
to SN 2017ein than to SN 2007gr. This is particularly
seen in Figure
6
, based on the MMT spectrum, in which
the well-resolved Na
i
doublet is strong (at 5904.9 and
5910.6
̊
A, respectively), attributable internally to the
host galaxy. The doublet is far weaker in the Galactic
foreground, consistent with the assumed low reddening
from this contribution. We have also indicated in Fig-
ure
6
the presumed location from the host galaxy of
the diffuse interstellar band (DIB) feature at
λ
5780, the
strength of which can be used to infer visual extinction
A
V
(
Phillips et al. 2013
) and which can change over time
in a subset of SNe Ib/c (
Milisavljevic et al. 2014
).
3.3.
Extinction to SN 2017ein
The extinction to SN 2017ein, as it turns out, is not
well determined from our data. We first attempted
to estimate the visual extinction from the SN spec-
tra.
Stritzinger et al.
(
2018
) presented for SESNe a
correlation between the equivalent width (EW) of the
Na
i
D feature and host-galaxy
A
V
. From the blended
8
Figure 4.
Early-time optical spectra obtained with the Kast
spectrograph on the 3-m Shane telescope at Lick Observa-
tory (plus one narrow-range MMT spectrum on day 15.6).
The ages are relative to
V
maximum. For comparison we
show spectra of SN 2007gr (
Valenti et al. 2008
;
Chen et al.
2014
) and SN 2004aw (
Taubenberger et al. 2006
) at roughly
similar ages. Various absorption lines and features seen in
the spectra are indicated. The SN 2017ein spectra shown in
this figure are available in the electronic journal and are al
so
posted on WISEReP, https://wiserep.weizmann.ac.il/.
Figure 5.
Expansion velocity evolution of SN 2017ein,
measured from the Fe
ii
λλ
4924, 5018, 5169 lines (filled
squares) and O
i
λ
7774 (filled circles). For comparison
we show the velocities of SN 2007gr from the Fe
ii
lines
(
Chen et al. 2014
; open squares), SN 2004aw from the O
i
λ
7774 line (
Taubenberger et al. 2006
; open triangles), and
several SNe Ic in the sample from
Liu et al.
(
2016
; crosses)
from Fe
ii
lines.
Table 4.
Photospheric Expansion Velocities for SN
2017ein
UT date
MJD
v
exp
(Fe
ii
)
v
exp
(O
i
)
(km s
−
1
(km s
−
1
)
2017 Jun 02.26 57906.76 9822
±
28 10362
±
49
2017 Jun 21.27 57925.77 8383
±
37 8973
±
122
2017 Jun 27.26 57931.76 8866
±
42 8427
±
98
2017 Jul 01.24 57935.74 8494
±
58 8306
±
105
2017 Jul 17.21 57951.71 8243
±
30 8097
±
230
2017 Jul 26.20 57960.70 8042
±
37 8391
±
137
2017 Jul 30.20 57964.70 7905
±
101 8405
±
89
2017 Aug 01.19 57966.69 8027
±
61 8368
±
97
Figure 6.
A moderate-resolution spectrum of SN 2017ein
obtained at the MMT on 2017 June 24. The interstellar Na
i
D doublet at the expected wavelengths for both the Milky
Way (“Gal.”) and NGC 3938 (“Host”) are indicated. Also
indicated is the expected location of the diffuse interstel-
lar band (“DIB”) near rest wavelength 5780
̊
A for the host
galaxy.
Na
i
D feature in the June 2 (day
−
6
.
3) spectrum, aris-
ing entirely from the host galaxy, we measured EW=
1
.
61
±
0
.
06
̊
A. From the MMT spectrum, we measured
from the resolved Na
i
D1 and D2 lines an EW of
0
.
67
±
0
.
04
̊
A and 0
.
72
±
0
.
01
̊
A, respectively. The
total EW of the feature from this spectrum is then
1
.
39
±
0
.
05
̊
A. Referring to
Stritzinger et al.
(
2018
; their
Figure 17), we then can see that
A
V
is somewhere in
the range of 0 to
∼
2 mag. Their best fit to the rela-
tion,
A
V
= 0
.
78(
±
0
.
15)
×
EW(Na
i
D), results in host
A
V
= 1
.
26
±
0
.
29 and
A
V
= 1
.
08
±
0
.
25 mag from the
9
Table 5.
Summary of Host Extinction Estimates for SN
2017ein
Method
A
V
R
V
(mag)
EW(Na
i
D)
a
∼
0–2
· · ·
1
.
26
±
0
.
29
· · ·
1
.
08
±
0
.
25
· · ·
EW(DIB
λ
5780)
b
.
1
.
06
· · ·
E
(
V
−
R
)
c
1
.
07
±
1
.
35 3.1
E
(
B
−
V
),
E
(
V
−
R
),
E
(
V
−
I
)
a
1.3–1.9 3.3–6.8
E
(
B
−
V
) and
E
(
V
−
R
) only
a
1.3–1.7
4.3
1.0–1.2
3.1
a
Following
Stritzinger et al.
(
2018
).
b
Following
Phillips et al.
(
2013
).
c
Following
Drout et al.
(
2011
).
blended feature in the Lick spectrum and the sum of the
resolved lines in the MMT spectrum, respectively.
Unfortunately, the DIB
λ
5780 feature is not par-
ticularly distinct in the MMT spectrum. Following
Phillips et al.
(
2013
), we fit a Gaussian of 2.1
̊
A full
width at half-maximum intensity (FWHM) to the spec-
trum and placed a 3
σ
upper limit of 203 m
̊
A to its equiv-
alent width, which, from their Equation 6 (which has a
50% systematic uncertainty), corresponds to
A
V
.
1
.
06
mag.
We can also estimate the amount of extinction to SN
2017ein through a color-curve comparison.
Drout et al.
(
2011
), from their systematic study of SN Ib/c light
curves, found that the intrinsic color, (
V
−
R
)
0
=
0
.
26
±
0
.
06 mag on
V
(10 d), 10 days past
V
(max) (see
also
Bianco et al. 2014
and
Dessart et al. 2016
). For
SN 2017ein, we estimate that
V
(10 d) was about JD
2,457,923.1, and on that day,
V
−
R
= 0
.
49
±
0
.
27 mag.
The Galactic foreground contribution to
E
(
V
−
R
) is
quite small, 0.012 mag (
Schlafly & Finkbeiner 2011
).
This would indicate that the host contribution to
E
(
V
−
R
) is 0
.
22
±
0
.
28 mag for SN 2017ein, where the uncer-
tainty includes both the measurement uncertainty in the
SN 2017ein
V
−
R
color from KAIT and the uncertainty
in the
Drout et al.
(
2011
) intrinsic color. Assuming a
Fitzpatrick
(
1999
) reddening law and
R
V
= 3
.
1, this
corresponds to
A
V
= 1
.
07
±
1
.
35 mag.
We compared the SN 2017ein
B
−
V
curve, corrected
for the low Galactic foreground reddening (0.019 mag),
with that of SN 2007gr, for which the Galactic red-
dening is
E
(
B
−
V
) = 0
.
055 mag (
Chen et al. 2014
;
via
Schlafly & Finkbeiner 2011
) and the host redden-
ing is assumed to be quite low (
Hunter et al. 2009
;
Drout et al. 2011
;
Chen et al. 2014
); see Figure
3
. Ad-
ditionally,
Stritzinger et al.
(
2018
), based on their sam-
ple from the Carnegie SN Project (CSP), provided color
templates, after correction for Galactic foreground red-
dening, for SESNe, specifically for SNe Ic, from 0 d to
+20 d relative to
V
(max). The relevant template to ap-
ply here is for
B
−
V
, which we show compared to the
color curve for SN 2017ein in Figure
3
. (Two other tem-
plates from
Stritzinger et al. 2018
,
V
−
r
, and
V
−
i
,
might have been applicable here; however, it is not read-
ily evident how to transform accurately for SNe Ic be-
tween the SDSS
r
and
i
bands used by the CSP and
the Johnson-Cousins
R
C
and
I
C
bands which we used
for this study.) Both the template and the SN 2007gr
B
−
V
color curve imply that
E
(
B
−
V
) = 0
.
34
±
0
.
07 mag
for SN 2017ein. From a minimum
χ
2
fitting of the SN
2017ein color curves to the corresponding SN 2007gr
ones, we estimate that
E
(
V
−
R
) = 0
.
24
±
0
.
06 and
E
(
V
−
I
) = 0
.
57
±
0
.
06 mag for SN 2017ein (the Galac-
tic foreground contribution to
E
(
V
−
I
) is 0.026 mag).
The value of
E
(
V
−
R
) for the entire color curve is con-
sistent with that inferred from the
Drout et al.
(
2011
)
fiducial value at
V
(10 d).
We performed an analysis of the three color excesses
in a manner similar to that of
Stritzinger et al.
(
2018
),
assuming a
Fitzpatrick
(
1999
) reddening law. We fit the
values of host
A
V
and
R
V
that were best constrained by
all three excesses. We found a rather large range in both
parameters,
A
V
= 1
.
3–1.9mag for
R
V
= 3
.
3–6.8. Im-
mediately, these imply that the visual extinction to SN
2017ein is appreciable, consistent with the large range
in
A
V
inferred from the Na
i
D feature strength, and
that the dust is dissimilar from the diffuse Galactic com-
ponent (for which
R
V
= 3
.
1).
Stritzinger et al.
(
2018
)
have indeed argued that
R
V
= 4
.
3 is typical for SNe Ic,
consistent with their general location in dusty environ-
ments presumably comprised of larger dust grains (we
note, however, that this inference is based on only one
of the events in their sample, which had more extreme
reddening). Values of
R
V
larger than 3.1 are typical for
massive star-forming regions; for example,
Smith
(
2002
)
found that the local
R
V
= 4
.
8 for clouds in the Carina
Nebula. This effect likely arises from the strong ultra-
violet radiation from young O-type stars which destroys
some of the smallest grains, leading to a flatter redden-
ing law.
The high values of
A
V
and
R
V
are most strongly
driven in our analysis by the value of
E
(
V
−
I
). If
we relax our analysis and consider only
E
(
B
−
V
) and
E
(
V
−
R
), and if we further assume that
R
V
= 4
.
3 (see
10
above), then
A
V
would be in the range of 1.3–1.7 mag,
consistent with the range in
A
V
when including all three
color excesses. If we, however, assume that
R
V
= 3
.
1,
similar to Galactic diffuse interstellar dust, then
A
V
would be distinctly lower, 1.0–1.2mag. For the sake of
discussion below, specifically when we analyze the na-
ture of the SN progenitor in Section
4.3
, we will consider
all three ranges in extinction and their implications.
We summarize all of the host-galaxy extinction esti-
mates for SN 2017ein in Table
5
. If a value for
R
V
is
either estimated or assumed, it is indicated in the table
as well.
3.4.
Distance to SN 2017ein
Another quantity for SN 2017ein that is not well
known is its distance,
D
. Several estimates exist for the
distance to the host galaxy, NGC 3938.
Poznanski et al.
(
2009
), based on assuming that SNe II-P are standard-
izable candles, estimated that the distance modulus to
SN 2005ay is
μ
= 31
.
27
±
0
.
13 mag (
D
= 17
.
9 Mpc).
Rodr ́ıguez et al.
(
2014
), invoking a similar, color-based
standardization of the absolute brightness of SNe II-P,
found that
μ
= 31
.
75
±
0
.
24 mag (22.4 Mpc) for their
method in the
V
band and 31
.
70
±
0
.
23mag (21.9 Mpc) in
I
. Several early Tully-Fisher estimates (
Bottinelli et al.
1984
;
1986
) resulted in far shorter distances, with
μ
≈
28
.
7
±
0
.
7 mag (
D
≈
5
.
7 Mpc), although
Tully
(
1988
)
lists
μ
= 31
.
15
±
0
.
40 mag (17.0 Mpc)
9
.
We therefore consider hereafter a range in possible
distance moduli to SN 2017ein of 31.15 to 31.75 mag
(
D
= 17
.
0 to 22.4 Mpc). Given this distance range
and the inferred SN extinction, above, for SN 2017ein,
M
V
(max)
≈ −
16
.
9 to
−
17
.
5 mag, which is consistent
with
−
17
.
2 mag for SN 2007gr (
Hunter et al. 2009
),
but less luminous than
−
18
.
0 mag for SN 2004aw
(
Taubenberger et al. 2006
).
3.5.
Quasi-Bolometric Light Curve
We constructed an early-time bolometric light curve
of SN 2017ein from our observed
BV R
C
I
C
light curves.
Given the relative sparseness in the photometric cover-
age and the large ranges in both distance and reddening,
high precision was not a particular concern. For that
reason, rather than performing our own detailed black-
body fitting of each set of photometric points, we uti-
lized the relations for SESNe from
Lyman et al.
(
2014
)
for estimating bolometric corrections to the absolute
B
and
V
light curves based on the reddening-corrected col-
ors, in order to generate the SN 2017ein bolometric light
9
See also the Extragalactic Distance Database,
http://edd.ifa.hawaii.edu/.
Figure 7.
Quasi-bolometric light curve of SN 2017ein (open
squares), assuming the bolometric corrections for SESNe
from
Lyman et al.
(
2014
). The uncertainties shown with
each data point arise primarily from the photometric mea-
surements and from the uncertainties in the bolometric cor-
rections. An error bar is also given, representing the addi-
tional uncertainty in both the reddening and the distance.
For comparison we show the bolometric light curves for SN
2007gr (
Hunter et al. 2009
;
Chen et al. 2014
; open triangles)
and SN 2004aw (
Taubenberger et al. 2006
; crosses). Addi-
tionally, we display the mean best-fit
Arnett
(
1982
) semi-
analytical model (solid curve), powered by radioactive dec
ay
of
56
Ni and
56
Co.
curve. We show this curve in Figure
7
. Although there
are uncertainties in both the photometry and the bolo-
metric corrections, these are dwarfed by the overall un-
certainties in both the distance and the extinction to the
SN.
For comparison we also show in Figure
7
the bolo-
metric light curves of SN 2007gr (
Hunter et al. 2009
;
Chen et al. 2014
) and SN 2004aw (
Taubenberger et al.
2006
). The bolometric luminosity of SN 2017ein is con-
sistent at peak, to within the large uncertainties, with
that of SN 2007gr; however, overall, the former is gener-
ally more luminous than the latter. SN 2004aw, which
was more luminous than SN 2007gr, is in agreement with
SN 2017ein, to within the uncertainties.
We have modeled the bolometric luminosity via a
semi-analytical light curve, following the method of
Arnett
(
1982
; see also
Valenti et al. 2008
;
Cano 2013
;
Taddia et al. 2015
;
Lyman et al. 2016
;
Prentice et al.
2016
;
Arnett et al. 2017
). The best-fit model to the
bolometric luminosity implies that the maximum bolo-
metric luminosity,
L
bol
(max), is (1.7–3.4)
×
10
42
erg s
−
1
,
and the maximum occurred around day 14 since explo-
11
Table 6.
Comparison of Light-Curve-Derived Parameters for SN 2017e
in
SN
M
(
56
Ni)
M
ej
E
K
Source
(
M
⊙
)
(
M
⊙
)
(10
51
erg)
2017ein
0.10–0.17
0.96–1.76 0.54–0.99
This work
a
· · ·
0.05–0.2
1
.
2
+0
.
3
−
0
.
2
0
.
7
+0
.
2
−
0
.
1
This work
b
2007gr
0
.
061
±
0
.
014
· · ·
· · ·
Chen et al.
(
2014
)
a
· · ·
0
.
076
±
0
.
020 2.0–3.5
1–4
Valenti et al.
(
2008
);
Hunter et al.
(
2009
)
a
· · ·
· · ·
∼
1
.
0
· · ·
Mazzali et al.
(
2010
)
c
· · ·
0
.
07
+0
.
01
−
0
.
01
1
.
2
+0
.
6
−
0
.
4
0
.
8
+0
.
3
−
0
.
3
Drout et al.
(
2011
)
b
2004aw
0.25–0.35
3.5–8.0
3.5–9.0
Taubenberger et al.
(
2006
)
a
;
Mazzali et al.
(
2017
)
c
· · ·
0
.
27
+0
.
05
−
0
.
05
4
.
5
+2
.
1
−
1
.
3
2
.
8
+1
.
3
−
0
.
8
Drout et al.
(
2011
)
b
SN Ic sample 0
.
33
±
0
.
07 5
.
75
±
2
.
09 1
.
75
±
0
.
24
Taddia et al.
(
2015
)
a
· · ·
0
.
24
±
0
.
15
1
.
7
+1
.
4
−
0
.
9
1
.
0
+0
.
9
−
0
.
5
Drout et al.
(
2011
)
b
a
From the quasi-bolometric light curve.
b
From multiband light curves.
c
From analysis of nebular spectra.
sion. We show the mean of the best-fit model in Fig-
ure
7
.
The
Arnett
(
1982
) model (see, e.g., Equations A1
and A2 of
Valenti et al. 2008
; Equations 1 and 2
of
Cano 2013
) fits two parameters, the nickel mass
M
(
56
Ni) and the diffusion timescale
τ
m
, where
τ
2
m
=
(
Cκ
opt
M
ej
)
/
(
βcv
sc
). In this expression,
M
ej
is the SN
ejecta mass,
c
is the speed of light,
β
is a constant of
integration
∼
13
.
8 (
Arnett 1982
),
v
sc
is the scale ve-
locity,
κ
opt
is the optical opacity (often assumed to be
0.07 cm
2
g
−
1
), and
C
is a constant of proportionality.
Note that the
M
ej
estimate is just for the mass involved
in the diffusion of SN light; if nonionized ejecta mass also
exists, the value of
M
ej
could be larger (
Wheeler et al.
2015
). The quantity
v
sc
is generally assumed to be the
photospheric velocity (
v
ph
) at maximum light. Note
that, for example,
Lyman et al.
(
2016
) assumed that
C
= 2, whereas
Chatzopoulos et al.
(
2012
) assumed
C
= 10
/
3 (since, generally, the mean photospheric ve-
locity is adopted, rather than the peak velocity, at
optical depth 1/3). The kinetic energy of the ejecta
is
E
K
= (3
/
10)
M
ej
v
2
sc
. From our fitting, the value of
τ
m
= 11
.
6 d. Based on our measurement of velocities
from the Fe
ii
lines (assumed to be nearly photospheric)
in the observed spectra, we have set
v
ph
= 9700km s
−
1
.
The results of the modeling are that
M
(
56
Ni) is
in the approximate range 0.09–0.18
M
⊙
. If
κ
opt
=
0
.
07 cm
2
g
−
1
, then for SN 2017ein,
M
ej
≈
1
.
45
M
⊙
and, thus,
E
K
≈
8
.
1
×
10
50
erg. For comparison with
SN 2007gr,
M
(
56
Ni)
≈
0
.
06
M
⊙
(
Chen et al. 2014
) to
∼
0
.
08
M
⊙
(
Valenti et al. 2008
;
Hunter et al. 2009
),
and
Hunter et al.
(
2009
) estimated that
M
ej
≈
2
.
0–
3.5
M
⊙
and
E
K
≈
(1–4)
×
10
51
erg. From nebular
spectra of SN 2007gr,
Mazzali et al.
(
2010
) estimated
a lower
M
ej
≈
1
M
⊙
. In contrast,
Taubenberger et al.
(
2006
) found for SN 2004aw significantly higher values
of
M
(
56
Ni) = 0
.
25–0.35
M
⊙
,
M
ej
= 3
.
5–8.0
M
⊙
, and
E
K
= (3.5–9.0)
×
10
51
erg (see also
Mazzali et al. 2017
).
We note that
Taddia et al.
(
2015
), from their sample of
SNe Ic, estimated mean values of
M
(
56
Ni) = 0
.
33 M
⊙
,
E
K
= 1
.
75
×
10
51
erg, and large
M
ej
= 5
.
75
M
⊙
. We
summarize this comparison in Table
6
.
We can also analyze these parameters,
M
(
56
Ni),
M
ej
, and
E
K
, from the observed light curve, follow-
ing
Drout et al.
(
2011
). From the simple analytic
model we constructed for the
R
-band light curve (see
Section
3.1
), we found that
M
R
(max)
≈ −
17
.
1 to
−
18
.
6 mag for SN 2017ein, given the distance and ex-
tinction range. From this model, the quantity ∆
m
15
,
the decrease in brightness between maximum and 15
days post-maximum, is 0.87 mag. This ∆
m
15
value
is consistent with that of SN 2007gr, to within the
uncertainties (
Drout et al. 2011
). From
M
R
(max)
and ∆
m
15
, we found that
M
(
56
Ni)
≈
0
.
05–0.2
M
⊙
,
which agrees well with that found from the bolomet-
ric curve. From the simple model, the characteris-
tic width of the light curve is
τ
c
= 10
±
1 d (again,
consistent with SN 2007gr;
Drout et al. 2011
). For
v
ph
= 9700 km s
−
1
, we can then infer that
M
ej
≈
1
.
2
M
⊙
and
E
K
≈
7
×
10
50
erg for SN 2017ein. Again, these
12
are completely consistent both with what we found
(above) for the bolometric light-curve analysis and with
SN 2007gr. We note that
Drout et al.
(
2011
) found
much higher values for these various parameters for SN
2004aw:
τ
c
≈
19,
M
(
56
Ni)
≈
0
.
27
M
⊙
,
M
ej
≈
4
.
5
M
⊙
,
and
E
K
≈
2
.
8
×
10
51
erg.
4.
THE SUPERNOVA PROGENITOR
4.1.
Progenitor Identification
To possibly identify a progenitor candidate in the
2006
HST
images, we astrometrically registered the 2006
WFPC2 F555W image mosaic to the 2017 WFC3 image
mosaic. Using 27 fiducial stars, we were able to regis-
ter the images to 0.26 WFPC2/WF pixel [1
σ
root-mean
square (rms); 26 milliarcsec]. As a result, SN 2017ein
would be at (798.55, 1766.58) in the WFPC2 image mo-
saic. An object can be seen at (798.33, 1766.54). This
is a difference of 0.22 pixel, which is within the rms
uncertainty in the astrometry. We therefore consider
this object to be a candidate for the progenitor of SN
2017ein. Note that we first preliminarily identified this
object in
Van Dyk et al.
(
2017
). We show the WFC3
and WFPC2 images to the same scale and orientation
in Figure
8
.
We have analyzed this object further by performing
PSF-fitting photometry of the
HST
WFPC2 images us-
ing Dolphot (
Dolphin 2000
;
2016
). Prior to this step,
we processed the individual WFPC2 frames with As-
troDrizzle, to attempt to flag cosmic-ray hits. We set
the Dolphot parameters FitSky=3 and RAper=8, us-
ing the TinyTim PSFs provided with Dolphot and set-
ting the parameters InterpPSFlib and WFPC2useCTE
to “true.” The SN site is contained in the WFPC2 chip
2. We find 24
.
56
±
0
.
11 and 24
.
58
±
0
.
17 VEGAMAG in
F555W and F814W, respectively. Dolphot also outputs
“object type”=“1” for the source, which indicates that
the routine considers it to be a “good star.” Other indi-
cators also point to a star-like character: the sharpness
and crowding parameters in each band are 0.002 and
−
0
.
007, and 0.030 and 0.076, respectively. The signal-
to-noise ratio (SNR) in each band for the star is appre-
ciable,
∼
10 and
∼
6 in F555W and F814W, respectively.
The output photometric quality flag from Dolphot is
also “0” for both bands, meaning (from the Dolphot
documentation) that the star was recovered very well
in the image data. The observed color of the source is
F555W
−
F814W=
−
0
.
02
±
0
.
20 mag.
4.2.
Metallicity of the Supernova Site
We can estimate the metallicity at the SN 2017ein site
from the oxygen abundance gradient for NGC 3938 pre-
sented by
Pilyugin et al.
(
2014
). Oxygen abundance is
1"
E
N
(a) WFC3 F438W 2017 June
SN 2017ein
(b) WFPC2 F555W 2006 Oct
SN 2017ein site
1"
(c) 2006 Oct zoom-in
SN 2017ein site
Figure 8.
(a) A portion of the WFC3/UVIS image of SN
2017ein obtained in F438W on 2017 June 12. The position
of the SN is indicated with tickmarks. (b) A portion of the
WFPC2 image of the host galaxy, NGC 3938, in F555W
on 2006 October 20, at the same orientation and scale as
panel (a). (c) A zoom-in of panel (b), with the SN position
indicated by a white circle with radius equal to the 0.26 pixe
l
rms astrometric uncertainty. The location of the SN is also
indicated in all three panels with tickmarks. We consider
the identified object to be the candidate progenitor of SN
2017ein. North is up, east is to the left in the figure.
13
often used as a proxy for metallicity. In that study the
central oxygen abundance is 12+log(O
/
H) = 8
.
79
±
0
.
05
and the gradient is
−
0
.
0413
±
0
.
0066 dex kpc
−
1
. We de-
projected an image of the host galaxy from the Digitized
Sky Survey, assuming the position angle (29
◦
) and in-
clination (18
◦
) from
Jarrett et al.
(
2003
), and calculate
that the SN is 44
′′
from the host nucleus. Given the
range in host-galaxy distance (above), this corresponds
to a
∼
3
.
7–4.8 kpc nuclear offset, and with this range we
estimate, from the published abundance gradient, that
12 + log(O
/
H)
≈
8
.
59–8.64.
Nebular H
α
and [N
ii
]
λ
6583 emission are also seen
in the MMT spectrum from June 24. Although not an
ideal strong-line indicator, the ratio of the line fluxes,
[N
ii
]
λ
6583/H
α
, known as the
N
2
index (
Pettini & Pagel
2004
), provides an estimate of metallicity. Dereddening
the spectrum assuming both
A
V
= 1 mag (with
R
V
=
3
.
1), we obtain
N
2
=
−
0
.
30. If the extinction in the
SN environment were as high as
A
V
= 2 mag, the index
varies slightly, to
−
0
.
31. Following the calibration of
this index by
Marino et al.
(
2013
), this would imply that
12 + log(O
/
H)
≈
8
.
60.
If we adopt the oxygen abundance for the Sun as
12+log(O
/
H)
≈
8
.
69 (
Asplund et al. 2009
), we can infer
that the metallicity at the SN site is somewhat subsolar
(i.e., [Fe/H]
≈ −
0
.
10).
4.3.
The Nature of the Progenitor Candidate
Assuming that the extinction to the SN we have es-
timated applies to the progenitor candidate as well,
and given the assumed range in distance, we find that
the object is consistent with being quite luminous and
blue. (Note that we are only considering here inter-
stellar extinction, which provides a lower limit on the
total extinction in the presence of any circumstellar ex-
tinction destroyed by the SN breakout; although not
typically expected for hot blue stars, some binary WC
stars are surrounded by dust formed by colliding winds;
e.g.,
Williams 2008
.) We show in Figure
9
three pos-
sible ranges of loci for the object in a color-magnitude
diagram (CMD), based on our assumptions about the
amount of reddening, as discussed in Section
3.3
.
We compared these loci in Figure
9
with single-
star evolutionary tracks from the Modules for Experi-
ments in Stellar Astrophysics (MESA) Isochrones and
Stellar Tracks (MIST;
Paxton et al. 2011
;
2013
;
2015
;
Choi et al. 2016
) v1.1 with rotation at
v
rot
/v
critical
=
0
.
4, solar-scaled abundances (assuming
Z
= 0
.
0142;
Asplund et al. 2009
), and subsolar metallicity [Fe/H]=
−
0
.
10. The tracks have been interpolated to the
WFPC2 VEGAMAG system by the MIST online in-
terface
10
. Among the tracks, we find that the candi-
date’s locus, assuming the lowest reddening, is roughly
consistent with, although somewhat blueward of, the
terminus of a star with initial mass 48–49
M
⊙
. The net
effect of rotation is to lead to a bigger He core mass,
and thus more luminous progenitor, for the same ini-
tial mass. If the star was not a rapid rotator, then the
implied initial mass (if single) would be even higher.
We note that some evolutionary models for single stars
with these initial masses at solar metallicity experience
“failed” explosions, with the cores collapsing directly to
black holes (e.g.,
Sukhbold et al. 2016
).
Additionally, we compared in Figure
9
with the
endpoints of binary-star models from BPASS v2.1
(
Eldridge et al. 2017
) at two different metallicities, solar
and somewhat subsolar (note that for BPASS,
Z
= 0
.
020
is considered solar metallicity). What is shown in the
figure is the combined light of both the primary and the
secondary in each of the binary models. (Here we have
assumed that the F555W and F814W bandpasses used
for BPASS are generic enough to apply for WFPC2.)
We have considered only systems for which the pri-
mary’s surface H mass fraction = 0, surface He mass
fraction
≤
0
.
27, and the ratio of the mass of remain-
ing He to the total ejecta mass
≤
0
.
1 at the model
endpoint. These criteria should be sufficient to approx-
imate a He-stripped SN Ic progenitor (J. J. Eldridge,
private communication). The systems allowed by the
color and luminosity range for the progenitor, again
assuming the lowest reddening, are to the lower-right
corner of this color-luminosity range. At
Z
= 0
.
020 the
lone system has a 60
M
⊙
primary, a mass ratio of the
two components
q
= 0
.
9 (i.e., the system consists of
two stars of nearly the same mass), and initial orbital
period log
P
(day) = 1
.
0. The primary of this system
ejects
M
ej
≈
10
M
⊙
and leaves behind a remnant with
∼
6
M
⊙
. At
Z
= 0
.
014 there is one system with these
same parameters, as well as systems with a 80
M
⊙
pri-
mary,
q
= 0
.
7, and a range of log
P
= 0
.
8 to 2.0. These
primaries eject
M
ej
≈
6
M
⊙
and have remnants with
∼
7–15
M
⊙
. At
Z
= 0
.
010 the systems have primaries
which, similarly, range from 60 to 80
M
⊙
, with
q
= 0
.
7–
0.8, and log
P
= 1–2.8. Ejecta masses are in the range
5–8
M
⊙
, and remnant masses are
∼
8–12
M
⊙
. The
remnants for all of these systems are most likely to be
black holes.
It may seem counterintuitive that the model binary
systems should possess a primary that is more massive
than the single-star models, with both binary and single-
10
http://waps.cfa.harvard.edu/MIST/.