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How to Make a SNIa J. Craig Wheeler Department of Astronomy. University of Texas at Austin Supernova Cosmology and Looking to the Future Cook's Branch.

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Presentation on theme: "How to Make a SNIa J. Craig Wheeler Department of Astronomy. University of Texas at Austin Supernova Cosmology and Looking to the Future Cook's Branch."— Presentation transcript:

1 How to Make a SNIa J. Craig Wheeler Department of Astronomy. University of Texas at Austin Supernova Cosmology and Looking to the Future Cook's Branch Nature Conservancy April 13, 2012

2 It should explode Bolometric light curve Multi-color light curve Multi-wavelength spectral evolution, UV to NIR, pre-maximum to nebular Criteria for a decent model of SN Ia: Marion et al. (2009), NIR spectra Wheeler & Benetti (2000)

3 Progress on understanding deflagration/detonation transition; there may be no such thing as a “distributed” flame (Poludnenko & Oran 2011). Delayed detonation model (Khokhlov 1991) does a decent job of all of those. Höflich 1995, and following Blondin et al. (2012) => carbon/oxygen white dwarf is grown slowly to degenerate ignition of carbon, simmering/smoldering, deflagration, detonation. This does not demand, but is consistent with single-degenerate binary evolution. Also dominance of blue-shifted Na D absorption (Sternberg et al. 2011).

4 Double Degenerate models tend to be very dynamic, messy. Tend to form large envelope from smaller mass, disrupted white dwarf that corrupts both the light curves and the spectra (Fryer et al. 2010; Raskin et al. 2012; Bloom et al. 2012; Shen et al. 2012). Possible support for double degenerate in models of SN 2011fe (R Ö pke et al. 2012), but detonate disrupted dwarf matter at 2x10 6 g cm -3 versus > 10 7 g cm -3 for all other work in the literature. Subsequent model has very different internal composition structure, potential problems with polarization, nebular spectra. R Ö pke et al. (2012)

5 I have a hard time believing that double-degenerate models can do the spectral evolution overall as well as delayed-detonation models. This encourages me to continue to think about single-degenerate models. Issues: ★ Demographics, population synthesis constraints ★ Lack of supersoft source progenitors (but double-degenerate has similar issues; Di Stefano 2010) ★ Lack of observed surviving companions, especially new Schaefer & Pagotta (2012) limit, M V > 8.4 (Even tighter HST limits on Tycho, SN 1006 Schmidt, depend on companion ejection velocity)

6 Lack of supersoft sources is hardly an issue (JCW& Pooley 2012, in prep). Very small column depth, > cm -2, can absorb the soft X-rays. A wind has been suggested, but there are other possibilities: Ca NIR high-velocity feature Gunk in Recurrent Novae 0.1 keV black body absorbed by various column depths of solar abundance matter.

7 High-velocity Ca II, Si II in SN 2011ao (HET: Marion, Vinko, JCW) Dynamically, need shell of ~ 0.02 M  at less than cm (Gerardy et al. 2004), column depth σ ~ 1.4x10 30 cm -2 M 31 R Patat et al. (2012) - Recurrent Nova RS Oph, M  in CSM from wind, previous outburst, < 4x10 14 cm, σ ~ 3x10 24 cm -2 R High velocity Si II (must come from WD) Hi V Ca II, 22,000 km/s CSM??

8 Close binaries tidally locked to orbit will be rapidly rotating. Under some circumstances, the companion to a white dwarf can undergo nearly homogeneous evolution on the main sequence. => Enhanced helium abundance (Livio & Truran 1992). Might change systematics of thermonuclear burning on white dwarf. Chatzopoulos, Robinson & JCW (2012, submitted).

9 Where is the surviving companion? What kind of main sequence star can beat the Schaefer & Pagnotta, M V > 8.4 limit? An M dwarf 70% of the stars in the Galaxy are M dwarfs, => ~ M dwarfs (Bochanski et al. 2011) 1 M dwarf in 1000 has a white dwarf companion (Law et al. 2012), PTF M dwarf study  10 9 M dwarf/white dwarf pairs in the Galaxy Bochanski, Hawley & West (2011) ~ white dwarfs in the Galaxy, 17% thin disk, 34% thick disk, 49% halo (Napiwotzki 2009) => To zeroth order, 10% all white dwarfs have M dwarf companions

10 Hot DA field white dwarf mass distribution Most peak at M  ~ 19% have ~ 0.8 M  ~ 9% have ~1.1 M  These masses may be underestimated by ~ 10% (Falcon et al. 2010). Lower metallicity stars make more massive white dwarfs (Willson 2000).  ~10 8 white dwarfs with M dwarf companions have mass ~ 1.1 M  Only need to accrete ~ 0.3 M  to reach the Chandrasekhar limit. Are these all carbon/oxygen white dwarfs? Kepler et al. (2007) If reach Chandrasekhar limit in 10 Gyr, could have 0.01 SN Ia per year.

11 Mass of M dwarfs (Delfosse et al. 2000) M0V ~ 0.6 M  M4V ~ 0.2 M  Fully convective limit ~ 0.35 M  Mass distribution of the M dwarfs: Number density is essentially flat (Bochanski et al. 2010) M dwarfs flare. Can the flare rate account for the required mass transfer rate? Unfortunately, no, unless presence of white dwarf has big effect. flare ~ 5x M  yr -1 (independent of energy of flares; derived from Aarnio et al. 2012, Hilton, 2012)

12 M dwarfs are magnetic: Sudden transition in field strength at M4 ~ 0.2 M  (Stassun et al. 2010) M0 ~ 100 G M4 ~ 1000 G Do not know that this transition would occur during binary mass loss, but wouldn’t that be interesting?

13 Suppose the white dwarf has a modest magnetic field, ~ 10 5 to 10 6 G Version of an intermediate polar or polar, but companion is also magnetic. M star dipole field much stronger than white dwarf field at M star: => merged field structure. Stable orientation, aligned dipoles in orbital plane (King, Frank & Whitehurst 1990). M star synchronously locked by tidal torque. White dwarf synchronously locked by magnetic torque (time scale, 10s of years; Campbell 1983, … 2010) White dwarf may be slowly rotating! Period ~ hour.

14 Mass transferred from M star to white dwarf will be locked in a magnetic bottle. Presence of magnetic bottle may affect mass transfer, loss processes. Not standard Roche lobe overflow. This magnetic field configuration may suppress a Hachisu-like wind from the white dwarf: P Mstar ~ B 2 /8π ~ 4x10 2 B 2 2 (R Mstar /a) 6 dynes/cm 2 > P wind ~ ½ρv 2 ~ 2.5 M-dot -6 v 8 a dynes/cm 2 May disrupt any accretion disk: angular velocity in magnetic bottle is orbital, not Keplerian near white dwarf.

15 Angular momentum loss will be driven by gravitational waves and loss of wind, magnetic braking, from magnetic bottle. (dJ/dt) wind ~ (1 – f)η dM/dt trans a 2 Ω a = separation Ω = orbital and spin angular velocities (all locked) f = fraction of transferred mass lost to wind η = scale factor (dJ/dt) tot = (dJ/dt) GW + (dJ/dt) wind = (dJ/dt) orb + (dJ/dt) wd + (dJ/dt) M Assume fill Roche lobe a ~ 1.52x10 11 (M M /M  ) 2/3 (M M + M wd ) 1/3 cm., neglect (dJ/dt) wd, R wd /R  ~ (M wd /M  ) -1/3, R M /R  ~ M M /M , (a/a) = {2/3 + (1-f)/3 [M M /(M M + M wd )]} (M M /M M )

16 From (dJ/dt) GW + (dJ/dt) wind = (dJ/dt) orb + (dJ/dt) wd + (dJ/dt) M (dJ/dt) GW = Fcn(M M, M wd, f, η) (M M /M M ) J orb / (dJ/dt) GW ~ 0.2 Gyr Neglect (dJ/dt) GW, (M M /M M ) drops out, unconstrained, f = Fcn(M M, M wd, η) ??

17 Major open issue: Can a system like this provide the mass transfer rate to yield non- degenerate H, He shell burning, beat the “nova” limit, grow the white dwarf to central carbon ignition? The mass loss may be channeled by the magnetic flux connecting the two stars, landing on a concentrated polar area of the white dwarf, enhancing the effective local rate of accretion compared to spherical accretion (Livio, Shankar & Truran 1988; but lateral diffusion rapid?). X-ray illumination from pole cap of white dwarf pointed at the companion might drive self-sustained mass transfer: may not need to fill Roche lobe. This will not be standard Roche lobe overflow mass transfer.

18 Poster object: Recurrent nova/polar T Pyx (Schaefer, Pagnotta & Shara 2010; Schaefer et al. 2012) White dwarf 1.3 M  Companion 0.1 M  dM/dt ~ M  /yr onto magnetic pole cap No disk luminosity 50,000 K black body P/dP/dt = 313,000 years

19 Other issues: Mass trapped in magnetic bottle: M b /M  ~ B 3 2 R b,13 3 T 4 -1 Trapped mass might be related to high-velocity Ca II feature, covering most, but not all solid angles. Magnetic bottle may be optically thick: τ ~ 50 κ B 2 2 R b,11 T 4 -1 If radiating shell-burning luminosity, ~ erg s -1 (reduced by area of pole cap) T eff ~ 3x10 5 K L 37 1/4 R b,11 -1/2

20 Other issues: Could start more massive, brighter than Schaefer/Pagnotta limit, just need to be dimmer by time of explosion. Transition to larger field on M dwarf during binary mass loss? Tidal locking enhance M dwarf rotation, dynamo? Mass at transition to fully convective at 0.35 M , prediction of spherical, non-rotating, non-magnetic models: probably wrong, probably higher mass. Accretion from magnetic stream onto white dwarf magnetic pole cap, nature of dynamics, nuclear burning.

21 Other issues: Remove mass from M dwarf, luminosity lags by thermal timescale, brighter than instantaneous mass would predict. What would such an enshrouded system look like? R M /a ~ 0.46 [M M /(M M + M wd ] 1/3 ~ 0.2; R 2 /4  a 2 ~ Schmidt limits as function of luminosity, velocity, Tycho, M V > 13 for v ~ 100 km/s (typical velocity if fill Roche lobe ~ 400 km/s).

22 Conclusions: Is any of this definitive? No Is this worth thinking about some more? I think so.

23 Ideas: Brian Schmidt - what happens to the M dwarf? Might lose mass, maybe a lot, might cool quickly on the convective timescale, be even dimmer than before the explosion. Heating, stripping, ablation. Chris Stubbs - synchrotron radiation in flow along field lines? Tonry - toroidal currents in M star, rotating, what does it do, torques. Probably end up with poles not pointing at each other, but canted. That was King et al. oscillating solution. Tonry has large sample of PanStAARs DA, DB, and WD/Mstars. Selected by colors. Maybe 10% of sample is WD/M?? HETDEX look at DES SN?

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