Modelling SN Type II

From Woosley et al. (2002)

Woosley Lecture 8

Iben (1985; Ql. J. RAS 26, 1)

5 M¯ evolution

Semiconvection

Semiconvection is the term applied to the slow mixing that goeson in a region that is stable by the strict Ledoux criterion but unstable by the Schwarzschild criterion.

Generally it is thought that this process does not contribute appreciably to energy transport (which is by radiation diffusion in semiconvectivezones), but it does slowly mix the composition. Its efficiency can be measured by a semiconvective diffusion coefficient that determineshow rapidly this mixing occurs.

Many papers have been written both regarding the effects of semiconvectionon stellar evolution and the estimation of this diffusion coefficient.

There are three places it is known to have potentially large effects:

• Following hydrogen burning just outside the helium core• During helium burning to determine the size of the C-O core• During silicon burning

One of the major effects of semiconvectionis to adjust the H/He abundance profilejust outside the H-depleted core (the helium core)

H-convective core

Langer, El Eid, and Fricke, A&A, 145, 179, (1985) (see also Grossman and Taam, MNRAS, 283, 1165, (1996))

30 M¯

Woosley & Weaver (1988; Phys. Rep. 163, 79)

No overshoot, semiconvection With overshoot, semiconvection

20 M¯

Semiconvection

No semiconvection

5000 yr between x

Langer & Maeder (1995; A&A 295, 685)

Woosley et al. (2002; RMP 74, 1015)

Mass loss – general features

See Chiosi & Maeder, ARAA, 24, 329 (1986) for a review

For how mass loss rates are measured see Dupree, ARAA, 24, 377(1986) – high resolution spectroscopy in IR, optical and uv; also radio measurements

For a review of the physics of mass loss see Castor in Physical Processes in Red Giants, ed. Iben and Renzini, Dordrecht: Reidel. See also Castor, Abott, & Klein, ApJ, 195, 157 (1975)

In massive stars, mass loss is chiefly a consequence of radiation pressure on grains and atoms. In quite massive stars, shocks and turbulence may be very important.

Humphreys – Davidson limit

Humphreys & Davidson (1979; ApJ 232, 409)

HD limit

HD limit

Humphreys (1984; IAU Symp 105, p. 279)

DEvolution with

mass loss

Maeder & Meynet (1988; A&AS 76, 411)

HD line

Mass Loss – Implications in Massive Stars

1) May reveal interior abundances as surface is peeled off ofthe star. E.g., CN processing, s-process, He, etc.

2) Structurally, the helium and heavy element core – onceits mass has been determined is insensitive to the presence of the envelope. If the entire envelope is lost however,the star enters a phase of rapid Wolf-Rayet mass loss that does greatly affect everything – the explosion, light curve,nucleosynthesis and remnant properties. A massive hydrogen envelope may also make the star more difficult to explode.

3) Mass loss sets an upper bound to the luminosity of redsupergiants. This limit is metallicity dependent.For solar metallicity, the maximum mass star that

dies with a hydrogen envelope attached is about 35 solar masses.

4) Mass loss – either in a binary or a strong wind – may be necessary to understand the relatively small mass of Type Ib supernova progenitors. In any case it is necessary to removethe envelope and make them Type I.

5) The nucleosynthesis ejected in the winds of starscan be important – especially WR-star winds.

6) In order to make gamma-ray bursts in the collapsarmodel for gamma-ray bursts, the final mass of the helium core must be large. Also the mass loss rateinferred from the optical afterglows of GRBs implya relatively low mass loss rate.

7) The winds of presupernova stars influence the radio luminosity of the supernova

8) Mass loss can influence whether the presupernova staris a red or blue supergiant.

9) The calculation of mass loss rates from theory is an important laboratory test ground for radiation hydrodynamics.

The Wolf-Rayet star WR224is found in the nebula M1-67which has a diameter of about 1000 AU

The wind is clearly veryclump and filamentary.

Nieuwenhuijzen and de Jager, A&A, 231, 134, (1990)

across the entire HR-diagram. This is multiplied by a factor toaccount for the metallicity-dependence of mass loss.

Studies by of O and B stars including B-supergiants, by Vink et al, A&A, 369, 574, (2001), indicate a metallicity sensitivity with scaling approximately as Z0.65.

Kudritzski, ApJ, 577, 389 (2002) in a theoretical treatmentof stellar winds (non-LTE, 2 million lines). Mass loss rate approximately proportional to ~Z1/2 down to Z = 0.0001times solar.

Wolf-Rayet stars – Langer, A&A, 220, 135, (1989)

More recently this has been divided by 2 - 3 to account foroverestimates made when clumping was ignored. Hamann andKoesterke, A&A, 335, 1003, Wellstein & Langer, A&A, 350, 148, (1998)

Models for optically thick radiation winds – Nugis and Lamers,A&A, 389, 162 (2002).

Parameterized results – Nugis and Lamers, A&A, 360, 227, (2000)

Y here is heliummass fraction at the surface. Z is metallicity atat the surface.

Wellstein and Langer (1998) corrected for Z-dependence and divided by 3 to correct for clumping is what we currently use.

Here Xs is the surface hydrogen mass fraction (WN stars)and the result should be multiplied by 1/3 (Z/Z¯)1/2..

Evolution with mass loss

Maeder & Meynet (1987; A&A 182, 243)

Wolf-Rayet stars

Maeder & Meynet (1987; A&A 182, 243)

Evolutionary sequences with mass loss –Chiosi and Maeder (1986; ARAA 24, 329)

time !

Chiosi and Maeder (1986 ; ARAA 24, 329)

Woosley et al. (2002; RMP 74, 1015)

Woosley et al. (2002; RMP 74, 1015)

Quirrenbach (2007; Science 317, 325)

Effects of rotation

Effects of rotation

Teff4 / F / geff

Observed gravity darkening

Domiciano de Souza et al. (2005; A&A 442, 567)

Altair ( Aquilae)

veq ' 230 km/sec

Teff4 / geff

*KH

Effects of rotation

See Kippenhahn & Weigert (1990; Sect. 42)

Meridional circulation

20 M¯

Solar composition

Meynet & Maeder (2002; A&A 390, 561)

Other instabilities that lead to mixing and the transport of angular momentum:

Eddington-Sweet and shear dominate.

energy available from shear adequate to (dynamically) overturn a layer. Must do work against gravity and any compositional barrier.

0 for stabili ytj

r

See Heger et al, ApJ, 528, 368 (2000)

1

12

.

12

1

)(

m

eff

G

gAM

iso mass loss

Maeder (1999; A&A 347, 185)

STELLAR WINDS & ROTATION

64.0

10

00030

2

06

LL

K

Enables a massive starto lose lots of mass andlittle angular momentum GRBs

André Maeder

= L /(4 c G M) = grad/g

Teff =25000 K

1

1

1,2

2

11

.

.

94

1

1

)0(

)(

critvvM

M

LARGEENHANCEMENTS !

André Maeder

Eta Carina

STRUCTURE• Oblateness (interior, surface)• New structure equations • Shellular rotation

MASS LOSS• Stellar winds • Anisotropic losses of mass and angular momentum

MIXING• Meridional circulation• Shear instabilities + diffusion• Horizontal turbulence• Advection + diffusion of angular momentum• Transport + diffusion of elements

André Maeder

Effects on evolution

Results:

• Fragile elements like Li, Be, B destroyed to a greater extent when rotational mixing is included. More rotation, more destruction.

• Higher mass loss

• Initially luminosities are lower (because g is lower) in rotating models. Later luminosity is higher because He-core is larger

• Broadening of the main sequence; longer main sequence lifetime

• More evidence of CN processing in rotating models. He, 13C, 14N, 17O, 23Na, and 26Al are enhanced in rapidly rotating stars while 12C, 15N, 16,18O, and 19F are depleted.

• Decrease in minimum mass for WR star formation.

These predictions are in good accord with what is observed.

Heger, Langer, and Woosley (2000), ApJ, 528, 368

Evolution Including Rotation

N

C

O

N`

20 M¯ with and without rotation

Heger, Langer, and Woosley (2000), ApJ, 528, 368

Without barrier With barrier

Without rotation

With rotation

Final angular momentum distribution is important to:

• Determine the physics of core collapse and explosion

• Determine the rotation rate and magnetic field strength of pulsars

• Determine the viability of the collapsar model for gamma-ray bursts.

Binary evolutionEquipotentials

Separate evolution

Binary evolutionEquipotentials

Mass transfer

Binary evolutionEquipotentials

Common envelope

Cases of mass

transfer

Paczynski (1971; ARAA 9, 183)

Binary evolutionAssume: 50% of all massive stars in binaries having P < 100 yr

Case A: During H core burning

Case B: After H core burning before He ignition

Case C: After He ignition

Common envelope: both stars fill their Roche envelope, either by birth or evolution

Podsiadlowski et al. (1992; ApJ 391, 246)

Binary evolution to Type Ia SN

Iben & Tutukov (1984; ApJS 54, 335

Triple-star evolution

Iben & Tutukov (1999; ApJ 511, 324)

Triple-star evolution

Iben & Tutukov (1999; ApJ 511, 324)