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Home , Superluminous supernova

what powers the brightest supernovae? time-domain astronomy a data driven revolution Palomar-48 inch optical ASASSN-15lh superluminous PTF-13ajg supernovae scp06f6 2005ap ptf09cnd 2008es 2006gy

2007bi type Ia ordinary core collapse supernovae ultra-long duration gamma-ray bursts levan et al 2014 Two ways to blow up a massive thermonuclear powered by core collapse to powered by runaway nuclear burning or no compact object formed

hongfeng Yu Argonne NL F. Ropke MPA Type I - no hydrogen spectral classification: Type II - hydrogen evolution of a shock core collapse breakout

H envelope

SN shock He core

C/O core

? shock revival core collapse shock stall Fe core

fallback neutrino neutron star spindown? GW emission cooling neutron star bounce

pre-bounce Core collapse supernova simulation 2D neutrino powered explosion Austin Harris (LBNL) with ORNL Chimera code core collapse supernova energetics ordinary case

gravitational energy released in neutron 2 GM 53 Eg GM 2 10 ergs ⇡ Rns ⇡ 53 Eg 2 10 ergs energy of supernova⇡ GMR explosionns ⇡ (kinetic53 and thermal energy) Eg Mni✏ni 10 ergs R td/tni Lni⇡ 1 ns 2 ⇡ e51 Eke ⇡Mvtni 10 ergs ⇡ 21 2⇡ 51 Eke Mv 10 ergs total energy radiated⇡ 2 in ordinary⇡ supernova light curve M M10 ✏ 15 M ni ni td/tni L ⇡ e49 Elcni Lt 10 ergs ⇡⇡ tni⇡ 42 43 1 L 10 10 ergs s ⇡ M M10ni✏ni15 Mtd/tni Lni ⇡ e ⇡ tni td 50 150 days 42 43 1 L 10⇡ 10 ergs s ⇡M 10 15 M ⇡ Esn 1 10 B td 4250⇡ 15043 days 1 L 10⇡ 10 ergs s ⇡ 13 R? 10 cm Esn ⇡1 10 B td 50⇡ 150 days ⇡ 1 R(t) Eth(t)=E013 R? 10 Rcm Esn⇡ 1 100 B ⇡  Esn Rsh 1 Rsh 45R(t) 1 Lsn 10 13 ergs s 4 ⇡ td ERthsn(tR)=?⇠E100 cm 10 R  ⇡ R   0 4 15 E RshR 10 R 10 cm1 R sn sh⇡ 45⇡R(t) 1 sh Lsn Eth(t)=10E0 ergs s 4 ⇡ td Rsn ⇠ R0 10 R   1  100 km s tsh = Rsh/vsh =4 2 years 15 EsnRshRsh10 R 45 10 cmv1sh Rsh Lsn ⇡ 10 ⇡ergs s 4 ⇡ td Rsn ⇠ 10 R   Esn R? 45 1 R1? 100 km s Lsn 4 10 ergs15 s 4 tsh⇡=tRdRshsh/vRsnsh 10=⇠ 2R years 10 cm 10 R  v  ⇡ ⇡ sh R R M 1 Etdsn= ⌧R? = ⇢45R 1001 km sR? Lsntsh = Rsh/vshc = 210 yearsergsc s⇠ Rc 4 ⇡ td Rsn ⇠  v 10 R  sh M Esn RR? td R M R? L td = ⌧ =⇠⇢10(45Rvt)ergsc s 1 sn c c ⇠ Rc 4 ⇡ td R sn ⇠  10 R 1/2  1/2  1/2 1/2 M M  v R M R M td 29 dayst 9 ⇠ vc td⇡= ⌧ d = M⇢R 0.1 10  c ✓⇠ (vt)◆c c ✓⇠ Rc◆ ✓ ◆   1/2 1 2 511/2 1/2 1/2 M Esn Mv M10M ergs  1 B v t ⇡292 daystd ⇡ ⌘ d ⇠ vc ⇡ ⇠M(vt)c 0.1 109  ✓ ◆ ✓ ◆ ✓ ◆ 2 4 1/2 L =4⇡R SB1/T2 1/2 1/2 M 1 2 M 51  v t Esn 29Mv days 10 ergs 1 B d ⇠ vc ⇡⇡ 2 ⇡M ⌘0.1 109  ✓ ◆ ✓ ◆ ✓ ◆ 1 1 1 credit: ASASSN Team 1051-1052 radiated energy!

super-luminous supernova

“ordinary” supernova superluminous supernova spectra Halpha Type II quimbySmith+ 2006et al. 2010

Type I Quimby+ 2007 Type I superluminous spectra SCP06f6 C/O model

FeII

CII/MgII OII/CII CII

stripped envelope progenitor CIII/CII Howell, kasen, et al., 2013 supernova light curve basics debris expands at v ~ 0.03c, cools by pdV work at t ~ weeks-months r ~ 1015 cm ~ 100 AU ρ ~ 10-13 g cm-3 translucent reheated to engine T ~ 5000-20000 K ? Z Z-1 9 e+ L >~ 10 Lsun

Υ radioactive decay νν 56Ni -> 56Co -> 56Fe supernova light curve basics light curve duration set by diffusion time the diffusion time of photons through optically thick remnant

but since the remnant is expanding, R = vt

solving for time (td ~ telapsed)

e.g., arnett (1979) supernova light curve basics of the light curve energy loses for adiabatically expanding radiation (pdV work)

simple estimate of emergent luminosity

assuming diffusion time td ~ 106 s How to power a super-luminous supernova light curve dump in energy after the ejecta has expanded (at t ~ tdiff) so radiation can escape immediately

• radioactivity: decay of freshly synthesized isotopes: e.g., 56Ni • shocks: interaction of the supernova ejecta with a dense surrounding medium • engines: later time energy injection from a central source (neutron star or black hole) Milisecond “Collapsar” Pulsational Pair instability Birth star: ~30-70 radioactivity ~1 MeV per 56Ni ASASSN-15lh need Mni >> Msun scp06f6 2005ap ptf09cnd 2008es 2006gy

2007bi ej = M MNi

type Ia

ej ordinary = 0.1 M core collapse MNi supernovae pair instability supernovae Rakavy, Shaviv, and Zinamon (1967), Bakrat, Rakavy, and Sack (1967) Bond, Arnett, and Carr (1984), Umeda and Nomoto (2001) Heger and Woosley (2002), Scannapeico et al 2005, Woosley (2007)

progenitor masses M ~ 150-260 Msun H H He He Si/Mg

C/O Si/O56 pairs trigger Ni e+/e- collapse and runaway thermonuclear burning

total exposion energy: 1051- 1053 ergs radioactive 56Ni produced: 0-50 Msun pair instability light curve models

M = 130 helium star

M = 250 M = 250 blue supergiant red supergiant

kasen, woosley, & heger (2011) type Ia type II pan, kasen, & Loeb (2012) ASASSN-15lh

scp06f6 2004ap2005ap ptf09cnd 2008es He 2006gy BSG RSG

2007bi type Ia ordinary core collapse supernovae pair instability supernovae SN2007bi as a pair instability SN? Gal Yam et al., Nature (2009)

helium bolometric New early time observations show rise too fast Nicholl et al 2013 shock powered light curves from interaction with circumstellar material

interacting“tamped” supernovae supernova models

supernova ejecta

slow moving debris at ~100 AU ejection ~2 years prior Mass loss from late stage nuclear burning? oxygen burning lasts ~1 year releases ~1052 ergs!

Tap that energy somehow: convectively driven waves, burning instabilitiies, pair instability

Quataert & Shiode (2012) Quataert, Fernandez, Kasen, et al (2016) Smith & Arnett (2014) Arnett & Meakin (2011) Woosley et al (2007) density

colliding shell velocity toy model colliding shell supernovae ~30% efficiency of conversion of kinetic energy to light

shell Esn = 1052 ergs Rsh = 1015 cm colliding shell model

pair instability (100 Msun He star) 4 Smith et al.

482 SMITH ET AL. Vol. 686

signatures of interaction Fig. 3.— Lick Observatory spectra of SN 2006gy at two different epochs,correctedforarangeofassumedhost-galaxyreddening corresponding to thenarrow values of A Rlinelisted emission at right (Cardelli et al. 1989). This is in addition to Galactic extinction of AR =0.43 mag. These are compared to the day 32 spectrum of the Type IIn SN2006tf(black)fromourdatabase,whichisaSNwithaspectrum similar to that of SNas 2006gy, in Type but seems II to SLSNe show little reddening. We adopt AR =1.25smith± 0.25 et mag al., for 2006, SN 2006gy; 2008 see text.

Fig. 16.—Cartoon illustration of the components of SN 2006tf at about 60 days after discovery, during the decline from the main light-curve peak. The primary feature is the massive postshock shell of gas, composed of the swept-up opaque pre-SN envelope around the star ejected in the decade before core collapse. Most of the mass is in the cold dense shell (CDS), bounded by the forward shock ( FS) and the reverse shock ( RS). Diffusion of radiation from this shocked shell produces the main continuum (1) and the intermediate-width component of H . This shell expands at constant speed into the preshock CSM (dense wind of the progenitor). The interior of the shell is filled by freely expanding SN ejecta, the outermost parts of which are ionized by radiation (wavy lines) propagating inward from the reverse shock, exciting the broad He i and O i features seen in the spectrum. There is also a second photosphere (2) in the SN ejecta, which is fainter than the main photosphere and can only be seen if the main shell thins or develops clumps as time proceeds. Right: More detailed depiction of the postshock gas, including the clumpy structure that forms due to instabilities in the cold dense shell layer. The dashed line here represents the photosphere at some arbitrary early time, working its way from left to right through the clumpy CDS as the SN expands. When it reaches a dense clump, the recombination photosphere will proceed through that clump, but for the regions between clumps it will eventually break through, allowing an observer to see the underlying SN ejecta.

needed to power the late-time luminosity (see previous point), shock by this time after explosion (Fig. 15). The broad features Fig. 4.— Dereddened visual-wavelength spectra ofand SN fully 2006gy consistent at t within= 36 the d anduncertainty 96 d after of the explosion, late-time lu- obtainedare also at seen Lick in Observatory P Cygni absorption and in He i k5876 and O i k7774. with the Keck II telescope, respectively. Several narrowminosity ab estimate.sorption This lines is also in aour factor high-resolution of 10 lower than Keck the spectrumThe absorption have be requiresen marked, some additional but background continuum  there are some remaining unidentified lines. Also plottednecessary is mass-lossaspectrumoftheTypeIaSN1991Tat rate in the decade just before core collapse,t = 35 dlight (Filippenko source, which et is al. likely 1992) to be for the diffusion of radiation from comparison with our day 36 spectrum of SN 2006gy;signifying there is a esssharpentially boost in noM˙ immediately similarity between before the star’s the two death. spectra.the inner SN ejecta deposited by shock energy or radioactive 6. The intermediate-width component of the H line arises decay. The luminosity required for the absorption strength im- imum light from a portion of the same Keckmostly spectrum in a swept-up, in dense,out postshock to ±6,000 cooling km shell s− expanding1 may be causedplies that the either underlying by electron SN was overluminous as well, indepen- 1 Fig. 4, with the flux normalized to the underlyingat a constant contin- speed of 2000scattering km sÀ (Fig. or 16). by This the is fastest the dom- SNdent ejecta. of CSM-interaction. inant speed of the forward shock plowing into the CSM. This 9. A possible explanation for why the broad features are seen uum level, and the velocity scale chosen withspeed the does narrow not change perceptiblyThe blue from edge day 32 of onward. the broad, Since blueshiftedonly from day 64 Hα throughabsorption 95 is that before that time, the shocked −1 −1 Hα emission feature at v =0kms . The Htheα shellprofile doesin not deceleratein Figure even though 5 indicates it is emitting an almost outflowshell speed was highly of 4,000 opaque km (the broad s , features reside interior to the Figure 5 reveals several different characteristic1051 velocitiergs, the shelles must alreadywhere be the very massive emission by day jumps 32, con- backreverse up just shock; to Fig. the 16). level Long that after that time (by day 445), the relevant to interpretations of SN 2006gy. First,sistent the with very our estimateswould above. be expected for a symmetricSN ejecta profile. luminosity This has jump probably is dropped far below that of 7. The−1 nature of the Balmer emission changes with time. At the ongoing CSM-interaction region. narrow emission component (FWHM ≈ 100early km stimes,) the has H /H fluxreadily ratio is apparent consistent with when recombina- we take the10. redshifted The luminosity side of the of intermediate-width the component of an associated P Cygni absorption feature thattion, indicates whereas at late times,broad the H emission/H ratio rises profile to more and than reflectH is it not to correlated the blue with side, the continuum to luminosity of the SN outflow speeds of 130 km s−1 (the trough) to10, 260 suggesting km s− that1 it becomessimulate dominated what by a direct symmetric collisional profile(Fig.12).Itrisesasthecontinuumluminosityfades.Compared would look like (Fig. excitation. to other SNe IIn, the H equivalent width is lower, but rises to (the blue edge) in the unshocked circumstellar8. Broad gas. wings In of H 5).may Because be due in part this to electron absorption scatter- tracessimilar the values speed at late of times the more dom- than 1 yr after explosion (Fig. 13). addition to Hα, several lines identified in Figsing, but 4 there and also 5 appearsinant to be an absorbing underlying broad material emission alongThis the is another line of clue sight that SN at 2006tf this has some additional source of − also have narrow absorption features. component, seen almostepoch, exclusively we at blueshifted take this speeds speed up to of 4,000continuum km luminosity s 1 to represent at early times, which is likely to be the about 7500 km s 1 (Fig. 8). This broad component appears slow diffusion of radiation from the massive swept-up opaque A broad Hα emission component has an apparentÀ À dense material swept up by the SN blast wave in the cir- −1 sometime after day 41, is seen on days 64 through 95, and dis- shell that mimics a normal H-recombination SN atmosphere, FWHM ≈ 2400 km s that is similar toappears Hβ at again early at very latecumstellar times. We propose material that this feature (CSM) cor- interactionbut at constant hypothesis, velocity. which times (Harutyunyan et al. 2006). The trueresponds unabsorbed to the outermostshould parts of the closely SN ejecta trace that have the almost speed of11. the The blast intermediate-width wave itself. postshock H emission has pro- FWHM of this broad Hα component is largerreached because the reverse of shock (seeThe Fig. broad-line 16). Material traveling profile at this differsnounced from asymmetrythe smooth at late broad times, showing an asymmetric and speed would, in fact, just about reach the radius of the reverse blueshifted profile at velocities within roughly 1000 km s 1. the broad blueshifted absorption. Extended faint wings parts of Hα profiles normally seen in SNe IIn (e.g., Æ À crab wind nebula centralfrom gaenslar engine and slane (2006) power from neutron star spindown

pulsar B ~ 5x1012 g from gaenslar and slane (2006) P ~ 20 ms neutron star spindown magnetized neutron stars release their rotational energy by magnetic dipole emission see Ostriker & Gunn (1974), Bodenheimer & Ostriker (1974), Gaffet (1977) rotational energy

spindown timescale ultra-luminous supernovae from to dump in energy at the right time (≈ months) requires the right magnetic field B and period P (e.g., the Crab nebula) 12 B ~ 5 x 10 g; P ~ 20 ms 49 Em ~ 5 x 10 ergs; tm ~ 2,000 years too slow.... magnetar model of gamma-ray bursts e.g., Thompson et al., (2004) 15 Bucciantini et al., (2007, 2008) B ~ 3 x10 g; P ~ 1 ms Uzdensky and MacFadyen (2007) Em ~ 2 x 1052 ergs; tm ~ 1 minute too fast...

magnetar powered super-luminous supernovae 14 B ~ 1x10 g; P ~ 4 ms 51 Em ~ 10 ergs; tm ~ months just right... dynamics of magnetar energy injection

chen et al 2016 magnetar energy deposition magnetic dipole spindown in vacuum:

mechanism and thermalization efficiency

magnetized pulsar wind particle acceleration hard photon production (inverse compton/synchrotron) x-ray absorption thermal optical radiation Metzger et al 2015, Kasen, Metzger, & Bildsten (2016) geometry spherical or bipolar? 1D dynamics of magnetar powered supernova Kasen & Bildsten (2010) theoretical maximum P = 1 ms 2-3 x 1045 ergs/s

P = 2 ms

P = 5 ms

longer spindown time magnetar theoretical maximum ~ 2x1045 ergs/s ASASSN-15lh

scp06f6 2005ap ptf09cnd 2008es P =1 ms 2006gy

2007bi ej = M Ni M P =5 ms

ej = 0.1 M MNi bolometric magnetar models kasen and bildsten (2010) an early signature of the engine?

nicholl+ (2015)

c.f. leloudas+ (2012) (2012)

> bright supernova + magnetar? > CSM + CSM (moriya+ 2012) > CSM interaction + magnetar? (piro 2015) two types of magnetar heating shell structure, t = 10 days

density

temperature

velocity double peaked light curves from magnetar driven shock breakout kasen, metzger, bildsten 2016 double peaked light curves from magnetar driven shock breakout kasen, metzger, bildsten 2016

with inefficient pulsar nebula thermalization at early times kasen, metzger, bildsten 2016 Early peaks are common in SLSNe

smith+ 2016 black hole central engines

an inefficiently cooled disk blows energetic winds

MacFadyen and Woosley (1999) rotating core collapse and disk formation rodrigo fernandez (UCB) C/O core He core H envelope 1D core collapse explosion model dexter and kasen (2013)

lower explosion energy and/or strong reverse shocks give continuous fallback and black hole feeding at later times

escape speed fallback accretion rate from low energy explosions of massive stars

dexter and kasen (2013) red supergiantquataert and kasen (2014) (Type II)

blue supergiant (Type II) M dot ~ t He star -5/3 (Type Ib)

˙ ˙ 2 44 1 ✏ M LBH = ✏Mc 10 ergs s 3 7 1 ⇡ 10 10 M s ✓ ◆✓ ◆ ˙ 2 1 ˙ 2 45 1 M vw LBH ⇠Mvw 10 ergs s 7 1 ⇡ 2 ⇡ 10 M s 0.1c ✓ ◆✓ ◆ 2 GM 53 Eg 10 ergs ⇡ Rns ⇡ 1 E Mv2 1051 ergs ke ⇡ 2 ⇡ E Lt 1049 ergs lc ⇡ ⇡

Mni✏ni td/tni Lni e ⇡ tni

M 10 15 M ⇡

42 43 1 L 10 10 ergs s ⇡

t 50 150 days d ⇡

E 1 10 B sn ⇡

R 1013 cm ? ⇡

1 R(t) E (t)=E th 0 R  0

Esn Rsh 45 1 Rsh Lsn 10 ergs s 4 ⇡ td Rsn ⇠ 10 R   4 15 Rsh 10 R 10 cm ⇡ ⇡

100 km s 1 t = R /v = 2 years sh sh sh v  sh

Esn R? 45 1 R? Lsn 10 ergs s 4 ⇡ td Rsn ⇠ 10 R   R R M t = ⌧ = ⇢R d c c ⇠ Rc  

1 Accretion Powered Supernova Light Curves 5

1/2 where t0 (2Gρ0)− (cf. Eq. 2 of QK12). For α < 0, the enclosed≡ mass is roughly constant, and the accretion rate is: sn2008es simple toy light sn1998bw (2α 3)/3 sn2008d 3 − curve models 8π ρ0r0 t 44 sn2010x M˙ = , (7) 10 3 t0 t0 comparison ! " ) 3/2 -1 to observed where now t0 πr0 / 2GM(r0). In this way, the freefall accretion≡ rate provides information about the supernova density profile of the progenitor# star. 1043 For bound material with vesc v, the maximum ra- light curves 2 2 1 ≃ dius, r r (1 v /v )− , becomes much larger than 1 ≡ 0 − esc the initial one, r0. Then the asymptotic fallback rate, ˙ 5/3 M t− , applies (Michel 1988; Chevalier 1989). This 42 asymptotic∝ scaling applies at the latest times in all three Luminosity (ergs s 10 curves in Figure 2. Using the ballistics solution from Chevalier (1989), we can bridge these two asymptotic limits to analytically estimate the fallback accretion rate at all times for comparison with our numerical calcula- 1041 tions. For each mass shell, the downstream shock velocity is 0 20 40 60 80 100 taken from the analytic formulae in Matzner & McKee Days since explosion (1999), which are typically an excellent approximation to the numerical calculations. Then the total fallback time Figure 5. Comparison of fallback powered light curves (solid for each mass element can be calculated from Eq. 3.7 of lines) with some observed supernovae. The parameters for these Chevalier (1989), and its time derivative is an approxi- events are given in Table 1. The orange dashed curve assumes toff = 7 days. Data points are taken from Gezari et al. (2009, mate accretion rate. This assumes that pressure effects SN 2008es), Mazzali et al. (2008, SN 1998bw and SN 2008D), and are negligible, which is incorrect. However, the true ac- Kasliwal et al. (2010, SN 2010X). celeration measured from the numerical calculations de- scribed below turns out to usually be roughly half of the We explore possible outcomes for supernova light gravitational acceleration. curves from the injection of accretion energy. A wide This ballistic estimate reproduces the fallback accre- range of explosion energies is used for each progenitor to tion rate at all times in many progenitors. However, explore the full range of possible outcomes. Only explo- in some cases (particularly blue supergiants such as sions with positive total energy at late times are consid- SN1987A, Chevalier 1989) the reverse shock formed at ered. Approximate one zone light curves are calculated the hydrogen-helium interface is strong enough to decel- using the methods described in Appendix A. For these erate portions of the ejecta below the escape speed. This calculations, we need the diffusion time: enhances the accretion rate at late times, and can sig- nificantly add to the remnant mass (Zhang et al. 2008). 3 Mκ 3 (Mej + Mfb)κ The reverse shock formation and evolution is analagous td = = , (8) 4π vc %4π vf c to that formed when the forward shock breaks out of $ the star and into the interstellar medium (e.g., McKee where Mfb = ξ M˙ fbdt is the total outflow mass, 1974; Chevalier 1982). As the simplest possible reverse E = ϵM˙ c2 is the injected accretion energy, and shock prescription, we solve the strong shock jump con- fb fb & ditions for the reverse shock velocity and the downstream vf = (Esn + Efb)/(Mej + Mfb) is the final ejecta ve- locity. We assume an opacity κ =0.2 g cm 1. Note velocity at the boundary of 100% helium and hydrogen # − layers: vRS v0 0.6v. For simplicity, we take vRS that there is an ambiguity in determining Mfb, depend- to be constant≤ at its≃ initial value. Then the location of ing on the interpretation of the fudge factor ξ. If ξ in- intersection between ejecta and the reverse shock can be dicates the fraction of outflow mass that interacts with found, as well as the resulting ballistic t(M) for mate- the supernova ejecta, then the above expression for Mfb rial that is recaptured after passing through the reverse applies. If on the other hand, the mass transfer to the shock. The reverse shock prescription is important for ejecta is more efficient while the specific energy of the the blue Z29 curve in Figure 2. outflow is lower, Mfb could be significantly larger. The ballistic approximation does a reasonable job re- From the light curves, we measure the time to peak, producing the numerical calculations in all cases. The tp, and the peak luminosity, Lp. The results are shown 3 largest disagreement is in the reverse shock cases, where in Figure 3 for ϵ = 10− . Each point represents a the semi-analytic accretion rate overestimates (underes- single explosion energy and progenitor model, color- timates) the numerical results at early (late) times. The coded by the radius of the pre-supernova star (red for resulting remnant mass vs. initial mass distribution from R>1013cm (RSGs), purple for 1012cm < R < 1013cm, these explosions is in excellent agreement with Zhang blue for 1011cm < R < 1012cm (BSGs), and green for et al. (2008). R<1011cm (He or C/O stars). This radius also corre- sponds to the zero age : solar 4. POSSIBLE OUTCOMES 4 for RSGs, zero for BSGs, and 10− Z for stars in be- ⊙ diversity of outcomes

dexter and kasen (2013) ultra-long duration gamma-ray bursts levan et al 2014 an ultra-long (104 s) GRB with a luminous supernova Greiner et al 2015 diversity of transients from magnetars metzger, margalit, kasen & quataert (2015)