A hot and fast ultra-stripped supernova that likely
formed a compact neutron star binary
K. De,
1
∗
M. M. Kasliwal,
1
E. O. Ofek,
2
T. J. Moriya,
3
J. Burke,
4
,
5
Y. Cao,
6
S. B. Cenko,
7
,
8
G. B. Doran,
9
G. E. Duggan,
1
R. P. Fender,
10
C. Fransson,
11
A. Gal-Yam,
2
A. Horesh,
12
S. R. Kulkarni,
1
R. R. Laher,
13
R. Lunnan,
11
I. Manulis,
2
F. Masci,
13
P. A. Mazzali,
14
,
15
P. E. Nugent,
16
,
17
D. A. Perley,
14
T. Petrushevska,
18
,
19
A. L. Piro,
20
C. Rumsey,
21
J. Sollerman,
11
M. Sullivan,
22
and F. Taddia
11
∗
To whom correspondence should be addressed; E-mail: kde@astro.caltech.edu
1
Cahill Centre for Astrophysics, California Institute of Technology, 1200 East California Boule-
vard, Pasadena, CA 91125, USA.
2
Department of Particle Physics and Astrophysics, Faculty of Physics, The Weizmann Institute
of Science, Rehovot 76100, Israel.
3
Division of Theoretical Astronomy, National Astronomical Observatory of Japan, National
Institutes of Natural Sciences, 2-21-1 Osawa, Mitaka, Tokyo 181-8588, Japan.
4
Las Cumbres Observatory, 6740 Cortona Drive, Suite 102, Goleta, CA 93117-5575, USA.
5
Department of Physics, University of California, Santa Barbara, CA 93106-9530, USA.
6
Department of Astronomy, University of Washington, Box 351580, Seattle, WA 98195-1580,
USA.
7
Astrophysics Science Division, NASA Goddard Space Flight Center, Mail Code 661, Green-
belt, MD 20771, USA.
8
Joint Space-Science Institute, University of Maryland, College Park, MD 20742, USA.
9
Jet Propulsion Laboratory, California Institute of Technology, Pasadena, CA 91109, USA.
10
Department of Physics, Astrophysics, University of Oxford, Denys Wilkinson Building, Ox-
ford OX1 3RH, UK.
11
Oskar Klein Centre, Department of Astronomy, Stockholm University, 106 91 Stockholm,
Sweden.
12
Racah Institute of Physics, The Hebrew University of Jerusalem, Jerusalem, 91904, Israel.
13
Infrared Processing and Analysis Center, California Institute of Technology, MS 100-22,
Pasadena, CA 91125, USA.
14
Astrophysics Research Institute, Liverpool John Moores University, Liverpool L3 5RF, UK.
1
arXiv:1810.05181v1 [astro-ph.HE] 11 Oct 2018
15
Max-Planck-Institut f
̈
ur Astrophysik, Karl-Schwarzschild-Str. 1, D-85748 Garching bei M
̈
unchen,
Germany.
16
Lawrence Berkeley National Laboratory, Berkeley, California 94720, USA.
17
Department of Astronomy, University of California, Berkeley, CA, 94720-3411, USA.
18
Oskar Klein Centre, Department of Physics, Stockholm University, 106 91 Stockholm, Swe-
den.
19
Centre for Astrophysics and Cosmology, University of Nova Gorica, Vipavska 11c, 5270 Aj-
dov
ˇ
s
̆
cina, Slovenia.
20
The Observatories of the Carnegie Institution for Science, 813 Santa Barbara Street, Pasadena,
CA 91101, USA.
21
Astrophysics Group, Cavendish Laboratory, 19 J J Thomson Avenue, Cambridge CB3 0HE,
UK.
22
Department of Physics and Astronomy, University of Southampton, Southampton, SO17 1BJ,
UK.
2
Compact neutron star binary systems are produced from binary massive stars
through stellar evolution involving up to two supernova explosions. The final
stages in the formation of these systems have not been directly observed. We
report the discovery of iPTF 14gqr (SN 2014ft), a Type Ic supernova with a
fast evolving light curve indicating an extremely low ejecta mass (
≈
0
.
2
so-
lar masses) and low kinetic energy (
≈
2
×
10
50
ergs). Early photometry and
spectroscopy reveal evidence of shock cooling of an extended He-rich envelope,
likely ejected in an intense pre-explosion mass loss episode of the progenitor.
Taken together, we interpret iPTF 14gqr as evidence for ultra-stripped super-
novae that form neutron stars in compact binary systems.
Core-collapse supernovae (SNe) are the violent deaths of massive stars when they run out of
nuclear fuel in their cores and collapse, forming a neutron star (NS) or black hole (BH) (
1
). For
massive stars that have lost some or all of their outer hydrogen (H) and helium (He) envelope,
the resulting collapse produces a stripped envelope SN (
2
). The amount of material stripped
from the star is a sensitive function of the initial mass of the star and its environment; if the
star was born in a binary system, it also depends on the orbital properties of the system and the
nature of the companion (
2, 3
).
Because most massive stars are born in close binary systems (
4
), stripping via binary in-
teractions likely plays a large role in producing the observed diversity of stripped envelope
SNe (
5, 6
). For the most compact companions in close orbits, the stripping of massive stars may
be large enough to completely remove its outer layers, leaving behind a naked metal core close
to the minimum mass required for the core to collapse (the Chandrasekhar mass). If massive
enough, the highly stripped core eventually collapses to produce a faint and fast evolving SN
explosion which ejects a small amount of material (
7, 8
). Although it has been difficult to se-
3
curely identify these explosions, such ‘ultra-stripped’ SNe have been suggested to lead to the
formation of a variety of compact NS binary systems (i.e. a NS in orbit around another NS,
white dwarf (WD) or BH) (
7, 9
).
Discovery and follow-up of iPTF 14gqr
iPTF 14gqr was discovered by the intermediate Palomar Transient Factory (iPTF; (
10, 11
)) on
2014 October 14.18 UT (Coordinated Universal Time) at a
g
-band optical magnitude of
≈
20
.
2
mag. The source was not detected in the previous observation on 2014 October 13.32 (0.86 days
before discovery), with a limiting magnitude of
g
≥
21.5 mag. The transient was found in the
outskirts (at a projected offset of
≈
29
kpc from the center) of a tidally interacting spiral galaxy
(IV Zw 155) at a redshift
z
= 0
.
063
and luminosity distance
D
= 284
.
5
Megaparsecs (Figure
1). We obtained rapid ultraviolet (UV), optical and near-infrared (NIR) follow-up observations
of the source, including a sequence of four spectra within 24 hours from the first detection (
12
).
We also obtained multi-epoch X-ray and radio observations and found that the source re-
mained undetected at these wavelengths (
12
). These upper limits rule out luminous non-thermal
emission, such as typically seen in relativistic and gamma-ray burst (GRB) associated SNe, but
are not stringent enough to constrain the environment of the progenitor (Figure S11, Figure
S12).
Our photometric follow-up indicated that the source rapidly faded within a day of detection,
followed by re-brightening to a second peak on a longer timescale (rising over
≈
7
days; (
12
))
(Figure 2). The early decline was detected in all optical and UV photometric bands, and char-
acterized by a blackbody spectrum which cooled rapidly from a temperature
T >
32000
K
4
near first detection to
T
∼
10000
K at one day after discovery (Figure 3; Figure 4). Our early
spectra also exhibit blackbody continua with temperatures consistent with those inferred from
the photometry, superimposed with intermediate width emission lines of He
II
, C
III
and C
IV
.
Such high ionization lines, which are typically associated with elevated pre-explosion mass loss
episodes in massive stars, have not been seen in early spectra of previously observed hydrogen-
poor SNe. Although similar features are present in the early spectra of some hydrogen-rich
core-collapse SNe (
13–15
) (Figure S6), the relatively large widths of the lines [Full Width at
Half Maximum (FWHM)
∼
2000 - 4000 km s
−
1
] as well as the rapid evolution of the 4686
̊
A
emission feature (Figure 3) are not.
Spectra obtained near the second peak are dominated by emission from the expanding pho-
tosphere and exhibit relatively blue continua, with broad absorption features reminiscent of
normal stripped envelope SNe of Type Ic, that do not exhibit absorption lines of H or He in the
spectra (
16
) (Figure S7). We find associated absorption velocities of
∼
10,000 km s
−
1
(
12
). The
photometric properties of the second peak are broadly consistent with a number of previously
observed fast Type Ic events (Figure S3, Figure S4), but the rapidly declining first peak and
the fast rise time to the second peak are unlike previously observed events. The source quickly
faded after the second peak, declining at a rate of 0.21 mag day
−
1
in the
g
band (
12
). Our final
spectrum taken at
≈
34
days after explosion shows that the source exhibited an early transition
to the nebular phase, and on a timescale faster than previously observed core-collapse SNe. The
nebular phase spectrum exhibits prominent [Ca
II
] emission similar to several other Type Ic SNe
(Figure S8).
Multi-color photometry at multiple epochs allow us to trace the evolution of the optical
/ UV Spectral Energy Distribution (SED), which we use to construct bolometric light curves
5
that contain flux integrated over all wavelengths (Figure 3; Figure 4; (
12
)). We fit the pseudo-
bolometric light curve of iPTF 14gqr with a simple Arnett model (
17
) to estimate the explosion
parameters. Allowing the explosion time to vary as a free parameter, we estimate an ejecta mass
M
ej
≈
0.15 – 0.30 solar masses (M
), an explosion kinetic energy
E
K
≈
(1.0 – 1.9)
×
10
50
ergs
and synthesized Ni mass
M
Ni
≈
0
.
05
M
(
12
) (Figure 4, Figure S9). The inferred ejecta mass
is lower than known core-collapse Type Ic SNe (
18–20
), which have ejecta masses in the higher
range of
∼
0
.
7
−
15
M
, and with a mean of
2
−
3
M
over a sample of
≈
20 SNe. However,
the parameters of iPTF 14gqr are similar to those inferred for the rapidly evolving Type I SNe
SN 2005ek (
21
) and 2010X (
22
), whose physical origins remain a matter of debate.
The rapid decline of the first peak observed in iPTF 14gqr is reminiscent of shock cool-
ing emission from the outer layers of a progenitor after the core-collapse SN shock breaks
out (
23, 24
) (Figure S5). We consider alternative explanations (
12
) and find them to be incon-
sistent with the data. In particular, the observed double-peaked light curve in the redder optical
bands requires the presence of an extended low mass envelope around the progenitor (
24, 25
).
To constrain the properties of such an envelope, we use models (
25
) to construct multi-color
light curves for a range of masses and radii of the envelope (
M
e
and
R
e
respectively). We find
a best-fitting model of
M
e
∼
8
×
10
−
3
M
and
R
e
∼
3
×
10
13
cm (
∼
450
solar radii (R
)) (
12
)
(Figure 2 & Figure S10). Even though the model considered here is simplified (e.g. it ignores
the density structure of the envelope), we expect the estimated parameters to be accurate within
an order of magnitude (
26
), leading us to conclude that the progenitor was surrounded by an
extended envelope with a mass of
∼
0
.
01
M
at a radius of
∼
500
R
.
We constrain the composition of the outer envelope using the early spectra. The emission
lines observed in the early spectra of iPTF 14gqr can be understood as arising from recombina-
6
tion in the outer regions of the extended circumstellar material (CSM), which was ionized by
the high energy radiation produced in the shock breakout (e.g. (
13, 15
); Figure S6). We estimate
the location and mass of the emitting He
II
from the luminosity of the early 4686
̊
A line, and
assuming a CSM density profile that varies with radius
r
as
∝
r
−
2
(
15
). We find the emitting
region to be located at
r
∼
6
×
10
14
τ
−
2
cm, and contain a helium mass
M
He
∼
0
.
01
τ
−
3
M
,
where
τ
is the optical depth of the region (
12
). The absence of prominent Lorentzian scatter-
ing profiles in the lines suggest that the optical depth is small and assuming
τ
≈
1
, we find
r
∼
6
×
10
14
cm (
8
.
5
×
10
3
R
) and
M
He
∼
0
.
01
M
. Because our calculations are based on
fitting a simple two-component Gaussian profile to the 4686
̊
A emission line (to estimate the
unknown contamination of C III at 4650
̊
A), these estimates are uncertain by a factor of a few.
Using the C IV 5801
̊
A lines and similar methods as above, we estimate a CSM carbon mass
of
∼
4
×
10
−
3
M
, while the hydrogen mass is constrained to be
<
10
−
3
M
. Additional
constraints based on light travel time arguments also suggest that the envelope was located at
r
≤
6
×
10
15
cm from the progenitor (
12
). The flash-ionized emission lines exhibit complex
asymmetric profiles (Figure 3) that we attribute to light travel time effects, given the large size
of the envelope and the high inferred wind velocities (
12, 27
).
An ultra-stripped progenitor
The low ejecta mass and explosion energy, as well as the presence of an extended He-rich en-
velope, indicate an unusual progenitor channel for iPTF 14gqr. The detection of the early shock
cooling emission indicates a core-collapse origin of the explosion, while the bright radioactivity
powered emission suggests that this explosion is associated with the class of iron core-collapse
explosions. The low ejecta mass together with the small remaining amount of He in the progen-
7
itor rule out models of single star evolution as well as a non-degenerate massive star companion
for the progenitor of iPTF 14gqr (
12
), leaving only the most compact companions (such as a
NS, WD or BH) as possible explanations of the highly stripped (or ‘ultra-stripped’) progenitor.
Ultra-stripped explosions have been modeled in the case of He star - NS binaries, where
stripping of the He star by a NS in a close orbit leads to the subsequent collapse of an ultra-
stripped He star (
7, 8, 28
). Hence, we compare theoretical bolometric light curves for ultra-
stripped explosions (
28
) to those of iPTF 14gqr in Figure 5, for a model with
M
ej
= 0
.
2
M
,
M
Ni
= 0
.
05
M
and
E
K
= 2
×
10
50
ergs. To account for the early declining emission, we
also add a component corresponding to shock cooling of an extended envelope, for
M
e
= 0
.
01
M
and
R
e
= 6
×
10
13
cm. The two component light curve matches the light curve data. We
also compare the spectroscopic properties of iPTF 14gqr to those of ultra-stripped SN models
in Figure 5. The models (
28
) assumed fully mixed ejecta that led to the production of strong
line blanketing features below
4000
̊
A , unlike this source. Thus, we re-calculated the models
for ejecta with no mixing (as with the light curve calculations), and were able to match to the
spectra of iPTF 14gqr near the second peak (Figure 5, Figure S13).
Our observations indicate the presence of an extended He-rich envelope around the progen-
itor at the time of collapse, thus providing insight into the terminal evolution of the progenitors
of ultra-stripped SNe, and more broadly, the lowest mass progenitors of core-collapse SNe. Us-
ing the line widths in our early spectra, we estimate that the emitting envelope was expanding
with a velocity of
∼
1000
−
2000
km s
−
1
at the time of collapse, consistent with the escape
velocity from a compact He star (
12
). When considered with the inferred size of the envelope
(at least
∼
500 R
), the velocities suggest that the envelope was ejected
∼
8
−
20
days prior to
the explosion.
8
The temporal coincidence of the ejection with the final SN suggests that the envelope was
likely associated with an intense pre-SN mass loss episode of the progenitor (
12
). Despite the
close stripping, ultra-stripped progenitors are expected to retain a small amount of He (
∼
0
.
01
M
) in their outer layers. The prominent He and C lines in the early spectra are consistent
with eruptive mass loss when considering the expected surface compositions of ultra-stripped
progenitors (
8
). The timescale of the ejection is similar to that expected for silicon flashes (
∼
2
weeks before explosion) in the terminal evolution of low mass metal cores (
29
), that have been
suggested to lead to elevated mass loss episodes prior to the explosion. Such mass loss episodes
are relevant to ultra-stripped progenitors as well (
28–30
).
iPTF 14gqr exhibits a projected offset of
∼
15
kpc from the nearest spiral arms of its star
forming host galaxy (
12
), which is puzzling when compared to the expected locations of ultra-
stripped SNe (
8
). While we do not find evidence of an underlying stellar association or of galaxy
emission features in late-time imaging and spectroscopy, the limits are not sensitive enough to
rule out the presence of a dwarf galaxy or a star forming H-II region (characterized by its H
α
emission) at or near the transient location (
12
). Nonetheless, the tidally interacting environment
of the host galaxy suggests that outlying star formation in collisional debris is likely in this
system (
12, 31
), which could harbor young stellar systems (with ages of
∼
5 - 100 Myrs) in the
faint tidal tails (Figure S14). Hence, the discovery of a core-collapse SN in these outskirts is
consistent with our interpretation.
While a number of previously observed fast Type Ic SNe (e.g. SN 2005ek (
21
) and SN 2010X
(
22
)) were suggested to be members of the ultra-stripped SN class, it has been difficult to con-
firm a core-collapse origin for these explosions because these events were discovered only near
9
maximum of the radioactively powered peak. Specifically, without early photometry and spec-
troscopy that can reveal the presence of a shock cooling component, these fast transients are
also consistent with variants of models involving thermonuclear detonations on white dwarfs
(e.g. (
32–34
)). The early discovery and prompt follow-up of iPTF 14gqr establish the presence
of a shock cooling emission component that requires an extended progenitor consistent with a
core-collapse explosion. In the probable scenario that iPTF 14gqr formed a NS in the explosion
(we find a BH remnant to be unlikely given the observed properties of the SN (
12
)), the low
ejecta mass in the system suggests that the SN results in the formation of a bound and compact
NS binary system (
12
).
Implications for formation of compact NS binaries
Our interpretation of iPTF 14gqr as an ultra-stripped SN has implications in the wider context
of stellar evolution. Compact NS binary systems evolve from binary massive stars that undergo
several phases of mass transfer over their lifetime (Figure 6). The initial phases of such evolu-
tion, in which two massive stars evolve into interacting binaries consisting of a compact object
in orbit around a massive star (X-ray binaries) have been observed in several systems in the
local Universe (
35, 36
). However, the subsequent phases that lead to the formation of compact
NS binary systems, have not been observed. This is due to the low occurrence rates of such
systems, the short lifetimes (
∼
10
6
years) of the final stages and observational selection effects
disfavoring their detection (
8, 37, 38
).
Binary evolution models suggest that the subsequent evolution proceeds via a common enve-
lope phase, during which the loss of angular momentum via dynamical friction leads to the for-
mation of a close He star - compact object binary (
9, 39, 40
). An additional phase of close grav-
10
itational stripping by the compact companion then leads to the formation of an ultra-stripped
SN progenitor (
9
), with properties which can be inferred from our observations of iPTF 14gqr.
The measured orbital properties of known double NS systems suggest that the second NSs were
created in weak and low ejecta mass explosions that impart a small natal kick to the newborn
NS (
41, 42
).
The presence of the extended He-rich envelope in iPTF 14gqr along with the lack of He in
the low mass of ejecta suggest that the progenitor was highly stripped by a compact companion,
such that only a thin He layer was retained on its surface. This He layer was then ejected in
an intense pre-SN mass loss episode, as shown by the high velocity of the envelope. Taken
together, these observations provide evidence of the terminal evolution of a post common enve-
lope He star - compact object binary leading to the formation of a compact NS binary system
(Figure 6).
While wide binaries containing a NS and another compact object may be formed in non-
interacting systems of binary massive stars, ultra-stripped SNe have been suggested to precede
the formation of almost all compact NS binary systems (
8
). Thus, these explosions likely repre-
sent the only channel to forming NS-NS and NS-BH systems that are compact enough to merge
within the age of the universe and produce observable merger signals for joint gravitational
wave (e.g. (
43
)) and electromagnetic (e.g. (
44–46
)) observations (
8, 47, 48
). Given that only a
fraction of the systems produced by these explosions will merge within that time, the rates of
ultra-stripped explosions must be higher than the rates of their mergers.
11
References and Notes
1. S. E. Woosley, A. Heger, T. A. Weaver,
Rev. Mod. Phys.
74
, 1015 (2002).
2. S. J. Smartt,
Annu. Rev. Astron. Astrophys.
47
, 63 (2009).
3. N. Langer,
Annu. Rev. Astron. Astrophys.
50
, 107 (2012).
4. H. Sana,
et al.
,
Science
337
, 444 (2012).
5. S.-C. Yoon, S. E. Woosley, N. Langer,
Astrophys. J.
725
, 940 (2010).
6. N. Smith, W. Li, A. V. Filippenko, R. Chornock,
Mon. Not. R. Astron. Soc.
412
, 1522
(2011).
7. T. M. Tauris,
et al.
,
Astrophys. J.
778
, L23 (2013).
8. T. M. Tauris, N. Langer, P. Podsiadlowski,
Mon. Not. R. Astron. Soc.
451
, 2123 (2015).
9. T. M. Tauris,
et al.
,
Astrophys. J.
846
, 170 (2017).
10. Y. Cao, P. E. Nugent, M. M. Kasliwal,
Publ. Astron. Soc. Pac.
128
, 114502 (2016).
11. F. J. Masci,
et al.
,
Publ. Astron. Soc. Pac.
129
, 014002 (2017).
12. Supplementary materials are available on Science online.
13. A. Gal-Yam,
et al.
,
Nature
509
, 471 (2014).
14. D. Khazov,
et al.
,
Astrophys. J.
818
, 3 (2016).
15. O. Yaron,
et al.
,
Nature Phys.
13
, 510 (2017).
16. A. Gal-Yam,
Observational and Physical Classification of Supernovae
(Springer Interna-
tional Publishing, Cham, 2017), pp. 1–43.
12
17. W. D. Arnett,
Astrophys. J.
253
, 785 (1982).
18. M. R. Drout,
et al.
,
Astrophys. J.
741
, 97 (2011).
19. J. D. Lyman,
et al.
,
Mon. Not. R. Astron. Soc.
457
, 328 (2016).
20. F. Taddia,
et al.
,
Astron. Astrophys.
609
, A136 (2018).
21. M. R. Drout,
et al.
,
Astrophys. J.
774
, 58 (2013).
22. M. M. Kasliwal,
et al.
,
Astrophys. J.
723
, L98 (2010).
23. E. Nakar, A. L. Piro,
Astrophys. J.
788
, 193 (2014).
24. N. Sapir, E. Waxman,
Astrophys. J.
838
, 130 (2017).
25. A. L. Piro,
Astrophys. J.
808
, L51 (2015).
26. A. L. Piro,
et al.
,
Astrophys. J.
846
, 94 (2017).
27. G. Gr
̈
afener, J. S. Vink,
Mon. Not. R. Astron. Soc.
455
, 112 (2016).
28. T. J. Moriya,
et al.
,
Mon. Not. R. Astron. Soc.
466
, 2085 (2017).
29. S. E. Woosley, A. Heger,
Astrophys. J.
810
, 34 (2015).
30. B. M
̈
uller, D. Gay, A. Heger, T. Tauris, S. A. Sim,
arXiv:1803.03388
(2018).
31. M. Boquien,
et al.
,
Astron. J.
137
, 4561 (2009).
32. K. J. Shen, D. Kasen, N. N. Weinberg, L. Bildsten, E. Scannapieco,
Astrophys. J.
715
, 767
(2010).
33. B. D. Metzger,
Mon. Not. R. Astron. Soc.
419
, 827 (2012).
13
34. S. Darbha,
et al.
,
Mon. Not. R. Astron. Soc.
409
, 846 (2010).
35. O. G. Benvenuto, M. C. Bersten,
Close Binary Stellar Evolution and Supernovae
(Springer
International Publishing, Cham, 2017), pp. 1–22.
36. R. Walter, A. A. Lutovinov, E. Bozzo, S. S. Tsygankov,
Astron. Astrophys. Reviews
23
, 2
(2015).
37. Y. G
̈
otberg, S. E. de Mink, J. H. Groh,
Astron. Astrophys.
608
, A11 (2017).
38. E. Zapartas,
et al.
,
Astron. Astrophys.
601
, A29 (2017).
39. D. Bhattacharya, E. P. J. van den Heuvel,
Phys. Rep.
203
, 1 (1991).
40. T. M. Tauris, E. P. J. van den Heuvel,
Formation and evolution of compact stellar X-ray
sources
(2006), pp. 623–665.
41. R. D. Ferdman,
et al.
,
Astrophys. J.
767
, 85 (2013).
42. P. Beniamini, T. Piran,
Mon. Not. R. Astron. Soc.
456
, 4089 (2016).
43. B. P. Abbott,
et al.
,
Astrophys. J.
848
, L13 (2017).
44. B. P. Abbott,
et al.
,
Astrophys. J.
848
, L12 (2017).
45. E. Pian,
et al.
,
Nature
551
, 67 (2017).
46. D. Kasen, B. Metzger, J. Barnes, E. Quataert, E. Ramirez-Ruiz,
Nature
551
, 80 (2017).
47. R. Voss, T. M. Tauris,
Mon. Not. R. Astron. Soc.
342
, 1169 (2003).
48. Compact systems may also form by dynamical capture in dense stellar environments (
49
).
49. J. Grindlay, S. Portegies Zwart, S. McMillan,
Nature Phys.
2
, 116 (2006).
14
50. P. A. Mazzali, L. B. Lucy,
Astron. Astrophys.
279
, 447 (1993).
51. B. Abolfathi,
et al.
,
Astrophys. J. Suppl. Ser.
235
, 42 (2018).
52. T. M. Tauris, T. Sennels,
Astron. Astrophys.
355
, 236 (2000).
53. N. M. Law,
et al.
,
Publ. Astron. Soc. Pac.
121
, 1395 (2009).
54. G. Rahmer,
et al.
,
Ground-based and Airborne Instrumentation for Astronomy II
(2008),
vol. 7014 of
Proc. SPIE
, p. 70144Y.
55. N. M. Law,
et al.
,
Ground-based and Airborne Instrumentation for Astronomy III
(2010),
vol. 7735 of
Proc. SPIE
, p. 77353M.
56. S. B. Cenko,
et al.
,
Publ. Astron. Soc. Pac.
118
, 1396 (2006).
57. C. Fremling,
et al.
,
Astron. Astrophys.
593
, A68 (2016).
58. E. F. Schlafly, D. P. Finkbeiner,
Astrophys. J.
737
, 103 (2011).
59. T. M. Brown,
et al.
,
Publ. Astron. Soc. Pac.
125
, 1031 (2013).
60. S. Valenti,
et al.
,
Mon. Not. R. Astron. Soc.
459
, 3939 (2016).
61. P. W. A. Roming,
et al.
,
Space Sci. Rev.
120
, 95 (2005).
62. D. N. Burrows,
et al.
,
Space Sci. Rev.
120
, 165 (2005).
63.
http://heasarc.nasa.gov/lheasoft/
.
64. J. A. Cardelli, G. C. Clayton, J. S. Mathis,
Astrophys. J.
345
, 245 (1989).
65. B. Winkel,
et al.
,
Astron. Astrophys.
585
, A41 (2016).
15
66. C. Benn, K. Dee, T. Ag
́
ocs,
Ground-based and Airborne Instrumentation for Astronomy
II
(2008), vol. 7014 of
Proc. SPIE
, p. 70146X.
67. I. M. Hook,
et al.
,
Publ. Astron. Soc. Pac.
116
, 425 (2004).
68. J. B. Oke,
et al.
,
Publ. Astron. Soc. Pac.
107
, 375 (1995).
69. S. M. Faber,
et al.
,
Instrument Design and Performance for Optical/Infrared Ground-
based Telescopes
, M. Iye, A. F. M. Moorwood, eds. (2003), vol. 4841 of
Proc. SPIE
, pp.
1657–1669.
70. O. Yaron, A. Gal-Yam,
Publ. Astron. Soc. Pac.
124
, 668 (2012).
71. J. T. L. Zwart,
et al.
,
Mon. Not. R. Astron. Soc.
391
, 1545 (2008).
72. T. D. Staley, G. E. Anderson, AMIsurvey: Calibration and imaging pipeline for radio data,
Astrophysics Source Code Library (2015).
73.
http://www.astro.caltech.edu/
̃
dperley/programs/lpipe.html
.
74. F. Taddia,
et al.
,
Astron. Astrophys.
574
, A60 (2015).
75. D. Poznanski,
et al.
,
Science
327
, 58 (2010).
76. M. W. Richmond,
et al.
,
Astron. J.
111
, 327 (1996).
77. M. M. Kasliwal,
et al.
,
Astrophys. J.
755
, 161 (2012).
78. F. Taddia,
et al.
,
Astron. Astrophys.
592
, A89 (2016).
79. P. J. Brown,
et al.
,
Astron. J.
137
, 4517 (2009).
80. F. B. Bianco,
et al.
,
Astrophys. J. Suppl. Ser.
213
, 19 (2014).
16
81. I. Arcavi,
et al.
,
Astrophys. J.
742
, L18 (2011).
82. P. A. Crowther,
Annu. Rev. Astron. Astrophys.
45
, 177 (2007).
83. A. Sander, W.-R. Hamann, H. Todt,
Astron. Astrophys.
540
, A144 (2012).
84. A. V. Filippenko,
et al.
,
Astrophys. J.
450
, L11 (1995).
85. S. Taubenberger,
et al.
,
Mon. Not. R. Astron. Soc.
371
, 1459 (2006).
86. M. Modjaz, Y. Q. Liu, F. B. Bianco, O. Graur,
Astrophys. J.
832
, 108 (2016).
87. R. Lunnan,
et al.
,
Astrophys. J.
836
, 60 (2017).
88. A. V. Filippenko, A. C. Porter, W. L. W. Sargent,
Astron. J.
100
, 1575 (1990).
89. S. Valenti,
et al.
,
Mon. Not. R. Astron. Soc.
437
, 1519 (2014).
90. S. Valenti,
et al.
,
Astrophys. J.
673
, L155 (2008).
91. D. Foreman-Mackey, D. W. Hogg, D. Lang, J. Goodman,
Publ. Astron. Soc. Pac.
125
, 306
(2013).
92. Z. Cano,
Mon. Not. R. Astron. Soc.
434
, 1098 (2013).
93. L. Dessart,
et al.
,
Mon. Not. R. Astron. Soc.
458
, 1618 (2016).
94. A. L. Piro, V. S. Morozova,
Astrophys. J.
792
, L11 (2014).
95. E. Waxman, B. Katz,
arXiv:1607.01293
(2016).
96. E. O. Ofek,
et al.
,
Astrophys. J.
768
, 47 (2013).
97. Y. Cao,
et al.
,
Nature
521
, 328 (2015).
17
98. G. H. Marion,
et al.
,
Astrophys. J.
820
, 92 (2016).
99. G. Hosseinzadeh,
et al.
,
Astrophys. J.
845
, L11 (2017).
100. D. Kasen,
Astrophys. J.
708
, 1025 (2010).
101. N. Smith,
arXiv:1612.02006
(2016).
102. P. J. Storey, D. G. Hummer,
Mon. Not. R. Astron. Soc.
272
, 41 (1995).
103. R. L. Kingsburgh, M. J. Barlow, P. J. Storey,
Astron. Astrophys.
295
, 75 (1995).
104. R. A. Chevalier,
Astrophys. J.
258
, 790 (1982).
105. R. A. Chevalier,
Astrophys. J.
499
, 810 (1998).
106. S. R. Kulkarni,
et al.
,
Nature
395
, 663 (1998).
107. A. M. Soderberg,
et al.
,
Nature
442
, 1014 (2006).
108. D. A. Perley,
et al.
,
Astrophys. J.
781
, 37 (2014).
109. A. Corsi,
et al.
,
Astrophys. J.
830
, 42 (2016).
110. Y. Suwa, T. Yoshida, M. Shibata, H. Umeda, K. Takahashi,
Mon. Not. R. Astron. Soc.
454
,
3073 (2015).
111. T. Yoshida, Y. Suwa, H. Umeda, M. Shibata, K. Takahashi,
Mon. Not. R. Astron. Soc.
471
,
4275 (2017).
112. S. Hachinger,
et al.
,
Mon. Not. R. Astron. Soc.
422
, 70 (2012).
113. K. Nomoto,
et al.
,
Nature
371
, 227 (1994).
114. L. Dessart,
et al.
,
Mon. Not. R. Astron. Soc.
414
, 2985 (2011).
18
115. L. Dessart, D. J. Hillier, C. Li, S. Woosley,
Mon. Not. R. Astron. Soc.
424
, 2139 (2012).
116. L. Dessart,
et al.
,
Mon. Not. R. Astron. Soc.
453
, 2189 (2015).
117. S. Valenti,
et al.
,
Mon. Not. R. Astron. Soc.
383
, 1485 (2008).
118. J. D. Lyman, A. J. Levan, R. P. Church, M. B. Davies, N. R. Tanvir,
Mon. Not. R. Astron.
Soc.
444
, 2157 (2014).
119. A. W. McConnachie,
Astron. J.
144
, 4 (2012).
120. W. E. Harris,
Astron. J.
112
, 1487 (1996).
121. M. F. Skrutskie,
et al.
,
Astron. J.
131
, 1163 (2006).
122. L. Bianchi, A. Conti, B. Shiao,
Adv. Space Res.
53
, 900 (2014).
123. F. B. Bianco,
et al.
,
Astronomy and Computing
16
, 54 (2016).
124. L. J. Kewley, M. A. Dopita,
Astrophys. J. Suppl. Ser.
142
, 35 (2002).
125. L. J. Kewley, S. L. Ellison,
Astrophys. J.
681
, 1183 (2008).
126. M. Pettini, B. E. J. Pagel,
Mon. Not. R. Astron. Soc.
348
, L59 (2004).
127. S. S. McGaugh,
Astrophys. J.
380
, 140 (1991).
128. N. E. Sanders, E. M. Levesque, A. M. Soderberg,
Astrophys. J.
775
, 125 (2013).
129. M. Asplund, N. Grevesse, A. J. Sauval, P. Scott,
Annu. Rev. Astron. Astrophys.
47
, 481
(2009).
130. M. Kriek,
et al.
,
Astrophys. J.
700
, 221 (2009).
131. C. Maraston,
Mon. Not. R. Astron. Soc.
362
, 799 (2005).
19
132. G. Leloudas,
et al.
,
Astron. Astrophys.
530
, A95 (2011).
133. N. E. Sanders,
et al.
,
Astrophys. J.
758
, 132 (2012).
134. I. Arcavi,
et al.
,
Astrophys. J.
721
, 777 (2010).
135. J. L. Prieto, K. Z. Stanek, J. F. Beacom,
Astrophys. J.
673
, 999 (2008).
136. N. Smith,
et al.
,
Mon. Not. R. Astron. Soc.
420
, 1135 (2012).
137. W. Fong, E. Berger,
Astrophys. J.
776
, 18 (2013).
138. R. C. Kennicutt, Jr.,
Annu. Rev. Astron. Astrophys.
36
, 189 (1998).
139. P. A. Crowther,
Mon. Not. R. Astron. Soc.
428
, 1927 (2013).
140. D. E. Osterbrock, G. J. Ferland,
Astrophysics of gaseous nebulae and active galactic nu-
clei
(2006).
141. J. P. Anderson, S. M. Habergham, P. A. James, M. Hamuy,
Mon. Not. R. Astron. Soc.
424
,
1372 (2012).
142. J. K. Werk,
et al.
,
Astron. J.
139
, 279 (2010).
143. B. Mullan,
et al.
,
Astrophys. J.
731
, 93 (2011).
144. K. A. Knierman,
et al.
,
Astron. J.
126
, 1227 (2003).
145. H. D. Tran,
et al.
,
Astrophys. J.
585
, 750 (2003).
146. F. Renaud, C. M. Boily, T. Naab, C. Theis,
Astrophys. J.
706
, 67 (2009).
147. G. Trancho, N. Bastian, F. Schweizer, B. W. Miller,
Astrophys. J.
658
, 993 (2007).
148. I. Saviane, J. E. Hibbard, R. M. Rich,
Astron. J.
127
, 660 (2004).
20
149. C. Frohmaier, M. Sullivan, P. E. Nugent, D. A. Goldstein, J. DeRose,
Astrophys. J. Suppl.
Ser.
230
, 4 (2017).
150. E. Bellm, S. Kulkarni,
Nature Astron.
1
, 0071 (2017).
151. J. L. Tonry,
et al.
,
Publ. Astron. Soc. Pac.
130
, 988 (2018).
152. G. Meynet, A. Maeder,
Astron. Astrophys.
429
, 581 (2005).
153. J. J. Eldridge, J. S. Vink,
Astron. Astrophys.
452
, 295 (2006).
154. C. Georgy, G. Meynet, R. Walder, D. Folini, A. Maeder,
Astron. Astrophys.
502
, 611
(2009).
155. E. Zapartas,
et al.
,
Astrophys. J.
842
, 125 (2017).
156. C. L. Fryer,
Astrophys. J.
522
, 413 (1999).
157. C. L. Fryer, V. Kalogera,
Astrophys. J.
554
, 548 (2001).
158. A. I. MacFadyen, S. E. Woosley, A. Heger,
Astrophys. J.
550
, 410 (2001).
159. J. D. M. Dewi, O. R. Pols,
Mon. Not. R. Astron. Soc.
344
, 629 (2003).
160. N. Ivanova,
et al.
,
Astron. Astrophys. Rev.
21
, 59 (2013).
161. P. Podsiadlowski,
Evolution of Binary and Multiple Star Systems
, P. Podsiadlowski,
S. Rappaport, A. R. King, F. D’Antona, L. Burderi, eds. (2001), vol. 229 of
Astronom-
ical Society of the Pacific Conference Series
, p. 239.
162. S. E. Woosley, N. Langer, T. A. Weaver,
Astrophys. J.
448
, 315 (1995).
163. N. Smith,
Annu. Rev. Astron. Astrophys.
52
, 487 (2014).
21
164. E. Quataert, J. Shiode,
Mon. Not. R. Astron. Soc.
423
, L92 (2012).
165. J. H. Shiode, E. Quataert,
Astrophys. J.
780
, 96 (2014).
166. D. Milisavljevic,
et al.
,
Astrophys. J.
846
, 50 (2017).
167. R. Margutti,
et al.
,
Astrophys. J.
778
, 18 (2013).
168. A. M. Soderberg,
et al.
,
Nature
463
, 513 (2010).
169. P. Salas, F. E. Bauer, C. Stockdale, J. L. Prieto,
Mon. Not. R. Astron. Soc.
428
, 1207
(2013).
170. E. Berger, S. R. Kulkarni, R. A. Chevalier,
Astrophys. J.
577
, L5 (2002).
171. A. Horesh,
et al.
,
Astrophys. J.
778
, 63 (2013).
172. M. R. Drout,
et al.
,
Astrophys. J.
821
, 57 (2016).
173. A. M. Soderberg, A. Brunthaler, E. Nakar, R. A. Chevalier, M. F. Bietenholz,
Astrophys.
J.
725
, 922 (2010).
174. A. M. Soderberg,
et al.
,
Astrophys. J.
621
, 908 (2005).
175. S. Wellons, A. M. Soderberg, R. A. Chevalier,
Astrophys. J.
752
, 17 (2012).
176. K. W. Weiler,
et al.
,
Astrophys. J.
740
, 79 (2011).
177. A. M. Soderberg,
et al.
,
Nature
453
, 469 (2008).
178. A. J. van der Horst,
et al.
,
Astrophys. J.
726
, 99 (2011).
179. Y. Cao,
et al.
,
Astrophys. J.
775
, L7 (2013).
180. M. Sullivan,
et al.
,
Astrophys. J.
732
, 118 (2011).
22
Acknowledgements
We thank the anonymous referees for a careful reading of the manuscript, that helped improve
the quality of the paper. We thank C. Steidel, N. Stone, D. Stern, P. Hopkins, S. de Mink,
Y. Suwa, A. Heger and T. M. Tauris for valuable discussions. MMK thanks J. Fuller, E. S.
Phinney, L. Bildsten and E. Quataert for stimulating discussions at the Skyhouse during a PTF-
TN meeting. We thank Tim Staley and Gemma Anderson for help with scheduling of the AMI
observations. Additional facility acknowledgments are given in the Supplementary Material.
Funding
The Intermediate Palomar Transient Factory project is a scientific collaboration among the Cal-
ifornia Institute of Technology, Los Alamos National Laboratory, the University of Wisconsin,
Milwaukee, the Oskar Klein Center, the Weizmann Institute of Science, the TANGO Program
of the University System of Taiwan, and the Kavli Institute for the Physics and Mathematics
of the Universe. This work was supported by the GROWTH (Global Relay of Observatories
Watching Transients Happen) project funded by the National Science Foundation under PIRE
Grant No 1545949. GROWTH is a collaborative project among California Institute of Technol-
ogy (USA), University of Maryland College Park (USA), University of Wisconsin Milwaukee
(USA), Texas Tech University (USA), San Diego State University (USA), Los Alamos Na-
tional Laboratory (USA), Tokyo Institute of Technology (Japan), National Central University
(Taiwan), Indian Institute of Astrophysics (India), Indian Institute of Technology Bombay (In-
dia), Weizmann Institute of Science (Israel), The Oskar Klein Centre at Stockholm University
(Sweden), Humboldt University (Germany), Liverpool John Moores University (UK).
A.H. acknowledges support by the I-Core Program of the Planning and Budgeting Com-
mittee and the Israel Science Foundation. A.G.-Y. is supported by the EU via ERC grant No.
24
725161, the Quantum Universe I-Core program, the ISF, the BSF Transformative program and
by a Kimmel award. E.O.O. is grateful for support by grants from the Willner Family Lead-
ership Institute Ilan Gluzman (Secaucus NJ), Israel Science Foundation, Minerva, BSF, BSF-
transformative, and the I-Core program by the Israeli Committee for Planning and Budgeting
and the Israel Science Foundation (ISF). F. T. and J. S. gratefully acknowledge the support
from the Knut and Alice Wallenberg Foundation. The Oskar Klein Centre is funded by the
Swedish Research Council. This research used resources of the National Energy Research Sci-
entific Computing Center, a DOE Office of Science User Facility supported by the Office of
Science of the U.S. Department of Energy under Contract No. DE-AC02-05CH11231. M. S.
acknowledges support from EU/FP7 ERC grant no. [615929]. P. E. N. acknowledges support
from the DOE through DE-FOA-0001088, Analytical Modeling for Extreme-Scale Computing
Environments. T. J. M. is supported by the Grants-in-Aid for Scientific Research of the Japan
Society for the Promotion of Science (16H07413, 17H02864). Numerical computations were
in part carried out on PC cluster at Center for Computational Astrophysics, National Astronom-
ical Observatory of Japan. Part of this research was carried out at the Jet Propulsion Laboratory,
California Institute of Technology, under a contract with the National Aeronautics and Space
Administration.
Author contributions
KD and MMK initiated the study, conducted analysis and wrote the manuscript. IM initiated the
follow-up of the young transient. DAP, GED and YC conducted Keck and Palomar observations
and contributed to data reduction and manuscript preparation. SBC conducted Keck and Swift
observations and contributed to data reduction and manuscript preparation. MS conducted the
WHT observations and data reduction. FT and JS conducted NOT observations, data analy-
sis and contributed to manuscript preparation. JB conducted the LCO observations and data
25
reduction. TP conducted Gemini observations and data analysis. CR and RPF conducted the
AMI observations and data reduction. AH conducted the VLA observations and data reduction.
SRK is iPTF PI and contributed to manuscript preparation. TJM and PAM prepared the ultra-
stripped SN models presented in the paper. EOO, CF, AGY, RL, PEN and ALP contributed to
manuscript preparation. GBD, RRL and FM contributed to the machine learning codes used to
search for young transients.
Competing interests
The authors declare no competing interests.
Data and materials availability
All photometric data used in this paper are provided in the supplementary material (Table
S1 and Table S2), while all spectra are available via the WISeREP repository at
https:
//wiserep.weizmann.ac.il/
. The codes used for the ultra-stripped SN modeling are
presented in (
50
), while the synthetic spectra presented in this paper are available at
https:
//goo.gl/9gkc9M
.
List of Supplementary materials
Materials and Methods
Supplementary text
Figures S1 - S15
Tables S1 - S7
References (51 - 181)
26
Figure 1:
Discovery field and host galaxy of iPTF 14gqr
. A. An optical image of the field
from the Sloan Digital Sky Survey (SDSS); r and g filter images have been used for red and
cyan colors respectively). B. Composite RGB image (r, g and B filter images have been used
for red, green and blue colors respectively) of the iPTF 14gqr field from images taken near
the second peak (19 October 2014) with the Palomar 60-inch telescope (P60), showing a blue
transient inside the white dashed circle at the discovery location. C. Late-time composite R+G
image (R and G filter images have been used for red and cyan colors respectively) of the host
galaxy taken with the Low Resolution Imaging Spectrograph on the Keck-I telescope.
27
A
B
Figure 2:
Multi-color photometric observations of iPTF 14gqr.
A. Multi-color light curves of
iPTF 14gqr from our photometric follow-up observations (magnitudes are corrected for galactic
extinction, and offset vertically as indicated in the legend). Inverted triangles denote 5
σ
upper
limits while other symbols denote detections. Hollow inverted triangles are upper limits from
P48/P60 imaging and the filled inverted triangles are upper limits from Swift observations (filled
green triangles are
V
band limits from Swift). Epochs when spectra were obtained are marked in
both panels by vertical black dashed lines. B. Zoom-in of the early evolution of the light curve.
The black solid line shows the assumed explosion epoch. The colored solid lines show the
best-fitting shock cooling model for extended progenitors (
25
). Only photometric data before
the cyan dot-dashed vertical line were used in the fitting (
12
).
28
A
B
C
Figure 3:
Spectroscopic evolution of iPTF 14gqr.
A. Observed spectra before (gray) and after
(black) binning. The epochs of the spectra along with the scaling and vertical shifts used are
indicated next to each spectrum. B. Zoom-in of the early spectra, indicated by the black dashed
box in (A), showing rapid evolution of the
λ
4686 feature within 24 hours of discovery. The
x-axis indicates the velocity shift from the He II
λ
4686 line. The orange and cyan lines mark
the locations of the
λ
4686 line and the C III
λ
4650 line respectively. For the +13.9 h and +25.2
h spectra, additional magenta lines show the profiles of the C IV
λ
5801 and the C III
λ
5696
features respectively, at the same epochs. C. Scaled optical / UV SEDs of the photometry
and spectra obtained within the first light curve peak (see Figure 2) in magenta, along with
photometry near the second peak in orange. The circles indicate observed photometric fluxes,
while the triangle is a 5
σ
upper limit. The dashed black lines indicate the best fitting blackbody
SEDs including all optical / UV data points for the first peak and including only the optical data
points for the second peak (
12
).
29