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PRESUPERNOVA EVOLUTION AND EXPLOSION OF MASSIVE STARS: CHALLENGES OF THE NEXT CENTURY. Marco Limongi INAF – Osservatorio Astronomico di Roma, ITALY email: marco@oa-roma.inaf.it. Alessandro Chieffi INAF – Istituto di Astrofisica Spaziale e Fisica Cosmica, ITALY

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slide1

PRESUPERNOVA EVOLUTION AND EXPLOSION OF MASSIVE STARS: CHALLENGES OF THE NEXT CENTURY

Marco Limongi

INAF – Osservatorio Astronomico di Roma, ITALY

email: marco@oa-roma.inaf.it

Alessandro Chieffi

INAF – Istituto di Astrofisica Spaziale e Fisica Cosmica, ITALY

email: alessandro.chieffi@iasf-roma.inaf.it

slide2

WHY ARE MASSIVE STARS IMPORTANT IN THE GLOBAL EVOLUTION OF OUR UNIVERSE?

Chemical Evolution of Galaxies

Light up regions of stellar birth  induce star formation

Production of most of the elements (those necessary to life)

Mixing (winds and radiation) of the ISM

Production of neutron stars and black holes

Cosmology (PopIII):

Reionization of the Universe at z>5

Massive Remnants (Black Holes)  AGN progenitors

Pregalactic Chemical Enrichment

High Energy Astrophysics:

Production of long-lived radioactive isotopes:

(26Al, 56Co, 57Co, 44Ti, 60Fe)

GRB progenitors

The understanding of these stars, is crucial for the interpretation of many astrophysical events

slide3

OVERVIEW OF MASSIVE STARS EVOLUTION

Grid of 15 stellar models: 11, 12, 13, 14, 15, 16, 17, 20, 25, 30, 35, 40, 60, 80 and 120 M

Initial Solar Composition (A&G89)

All models computed with the FRANEC (Frascati RAphson Newton Evolutionary Code) release 5.050419

(Limongi & Chieffi 2006, ApJ, 647, 483)

  • Evolution followed from the Pre Main Sequence up to the beginning of the core collapse
  • 4 physical + N chemical equations (mixing+nuclear burning) fully coupled and solved simultaneously (Henyey)
  • Nuclear network very extended

282 nuclear species (H to Mo) and ~ 3000 processes (Fully Automated)

(NO Quasi (QSE) or Full Nuclear Statistical Equilibrium (NSE) approximation)

  • Convective Core Overshooting: d=0.2 Hp
  • Mass Loss: Vink et al. (2000,2001) (Teff>12000 K),De Jager (1988) (Teff<12000 K), Nugis & Lamers (2000) (Wolf-Rayet)/Langer 1898 (WNE/WCO)
slide4

CORE H BURNING

Convective Core

g

g

g

CNO Cycle

g

g

Core H Burning Models

g

g

g

Mmin(O) = 14 M

t(O)/t(H burning): 0.15 (14 M ) – 0.79 (120 M)

MASS LOSS

slide5

CORE He BURNING

3a+ 12C(a,g)16O

He Convective Core

g

g

g

Core He Burning Models

g

g

M ≤ 30 M RSG

g

g

M ≥ 35 M BSG

g

H burning shell

H exhausted core (He Core)

slide6

H

M ≥ 30 M

M ≥ 40 M

M ≥ 35 M

Log (Teff)>4.0

H/He

He/CO

H

WNL

H<0.4

WNE

(N)

He

WCO

He C,O

massive stars mass loss during h he burning
MASSIVE STARS: MASS LOSS DURING H-He BURNING

WR : Log10(Teff) > 4.0

  • WNL: 10-5< Hsup <0.4 (H burning, CNO, products)
  • WNE: Hsup<10-5 (No H)
  • WN/WC: 0.1 < X(C)/X(N) < 10 (both H and He burning products, N and C)
  • WC: X(C)/X(N) > 10 (He burning products)

O-Type: 60000 > T(K) > 33000

  • M < 30 Mסּexplode as Red SuperGiant (RSG)
  • M ≥ 30 Mסּ explode as Blue SuperGiant (BSG)
slide8

ADVANCED BURNING STAGES

Neutrino losses play a dominant role in the evolution of a massive star beyond core He burning (T>109 K)

g

n

n

H-burn. shell

g

g

CO core

n

n

Core burning

He core

g

g

n

g

n

He-burn. shell

g

g

n

n

g

The Nuclear Luminosity (Lnuc) closely follows the energy losses

Evolutionary times of the advanced burning stages reduce dramatically

slide10

ADVANCED BURNING STAGES

Four major burnings, i.e., carbon, neon, oxygen and silicon.

Central burning  formation of a convective core

Central exhaustion  shell burning  convective shell

Local exhaustion  shell burning shifts outward in mass

 convective shell

slide11

ADVANCED BURNING STAGES

The details of this behavior (number, timing, overlap of convective shells) is mainly driven by the CO core mass and by its chemical composition (12C, 16O)

CO core mass

Thermodynamic history

Basic fuel for all the nuclear burning stages after core He burning

12C, 16O

At core He exhaustion both the mass and the composition of the CO core scale with the initial mass

slide12

ADVANCED BURNING STAGES

...hence, the evolutionary behavior scales as well

In general, one to four carbon convective shells and one to three convective shell episodes for each of the neon, oxygen and silicon burnings occur.

The number of C convective shells decreases as the mass of the CO core increases (not the total mass!).

slide13

PRESUPERNOVA STAR

The complex interplay among the shell nuclear burnings and the timing of the convective zones determines in a direct way the final distribution of the chemical composition...

16O,24Mg, 28Si,29S, 30Si

14N, 13C, 17O

14N, 13C, 17O

12C, 16O

28Si,32S, 36Ar,40Ca, 34S, 38Ar

12C, 16O s-proc

20Ne,23Na,

24Mg,25Mg,

27Al,

s-proc

56,57,58Fe, 52,53,54Cr, 55Mn, 59Co, 62Ni

NSE

slide14

PRESUPERNOVA STAR

...and the density structure of the star at the presupernova stage

MFe=1.20-1.45 M for M ≤ 40 M

MFe=1.45-1.80 M for M > 40 M

The final Fe core Masses range between:

In general the higher is the mass of the CO core, the more

compact is the structure at the presupernova stage

slide15

THE SUPERNOVA PROBLEM

The most recent and detailed simulations of core collapse SN explosions show that:

  • the shock still stalls  No explosion is obtained
  • the energy of the explosion is a factor of 3 to 10 lower than usually observed

Work is underway by all the theoretical groups to better understand the problem and we may expect progresses in the next future

The simulation of the explosion of the envelope is mandatory to have information on:

  • the chemical yields (propagation of the shock wave  compression and heating explosive nucleosynthesis)
  • the initial mass-remnant mass relation
slide16

Shock Wave

Compression and Heating

Matter Falling Back

Matter Ejected into the ISM

Ekin1051 erg

Induced Expansion and Explosion

Mass Cut

Initial Remnant

Final Remnant

Initial Remnant

Fe core

EXPLOSION AND FALLBACK

  • Piston (Woosley & Weaver)
  • Thermal Bomb (Nomoto & Umeda)
  • Kinetic Bomb (Chieffi & Limongi)

Different ways of inducing the explosion

FB depends on the binding energy: the higher is the initial mass the higher is the binding energy

slide17

Z=Z

E=1051 erg

NL00

No Mass Loss

WIND

SNII

SNIb/c

RSG

WNL

Final Mass

WC/WO

WNE

He-Core Mass

Fallback

Remnant Mass

He-CC Mass

CO-Core Mass

Black Hole

Fe-Core Mass

Neutron Star

THE FINAL FATE OF A MASSIVE STAR

slide19

CONCLUSIONS. I

  • Stars with M<30 M explode as RSG Stars with M≥30 M explode as BSG

WNL: 25-30 M

  • The minimum masses for the formation of the various kind of Wolf-Rayet stars are:

WNE: 30-35 M

WNC: 35-40 M

MFe=1.20-1.45 M for M ≤ 40 M

  • The final Fe core Masses range between:

MFe=1.45-1.80 M for M > 40 M

30-35 M

  • The limiting mass between SNII and SNIb/c is:

SNII

SNIb/c

Salpeter IMF

25-30 M

  • The limiting mass between NS and BH formation is:

NS

BH

  • M>35 M (SNIb/c) do not contribute to the intermediate and heavy elements (large fallback)
slide20

MAIN UNCERTAINTIES

PRESUPERNOVA EVOLUTION:

  • Mass Loss during Blue and Red supergiant phases, and Wolf-Rayet stages
  • Treatment of Convection: extension of the convective zones (overshooting, semiconvection), interaction mixing-nuclear burning
  • 12C(a,g)16O cross section
  • Rotation

EXPLOSION:

  • Lack of an autoconsistent hydrodynamical model (neutrino transport)
  • Induced explosion [Explosion energy (where and how), time delay, fallback and mass cut (boundary conditions), mixing (inner and outer borders), extra-fallback, Ye variation, aspherical explosions]
slide21

THE ROLE OF THE MASS LOSS FOR WNE/WCO IN THE ADVANCED BURNING PHASES

Nugis & Lamers (2000) (NL00)

Langer (1989) (LA89)

Strong reduction of the He core during early core He burning

slide23

CONSEQUENCES ON THE FALLBACK

Final kinetic energy = 1 foe (1051 erg)

slide24

Z=Z

E=1051 erg

LA89

No Mass Loss

SNII

SNIb/c

RSG

WIND

WNL

WNE

Final Mass

WC/WO

He-Core Mass

He-CC Mass

CO-Core Mass

BH

Remnant Mass

Fe-Core Mass

NS

NS

THE FINAL FATE OF “LA89” MASSIVE STARS

slide26

TREATMENT OF CONVECTION

Convection is, in general, a hydrodynamical multi-D phenomenon  its inclusion in a hydrostatic 1-D stellar evolution code consititutes a great source of uncertainty

Mixing-Length theory:

  • Extension of the convective zones (stability criterion, overshooting, semiconvection)?
  • Temperature Gradient?
  • Interaction between nuclear burning and convective mixing?

What about Mixing-Length theory for advanced burning stages of massive stars?

Does it make sense?

slide27

He

X

T>4 108

produced

60Fe

22Ne, a

M

X

M > 35 MO

M

TREATMENT OF CONVECTION

PRODUCTION OF 60Fe IN MASSIVE STARS:

60Fe synthesize within the He convective shell

Convection

  • Preserves 60Fe from destruction
  • Brings new fuel (a, 22Ne)

He convective shell forms in a zone with variable composition

slide28

Si conv. shell

28Si

Si exhausted Core

Mass Fraction

Final Fe Core Mass

M

Si conv. shell

28Si

Si exhausted Core

Mass Fraction

Final Fe Core Mass

M

TREATMENT OF CONVECTION

THE MASS OF THE Fe CORE:

Core Collapse and Bounce

Fe core

Shock Wave

n

n

n

n

Energy Losses

1 x 1051 erg/0.1M

slide29

UNCERTAINTY ABOUT 12C(a,g)16O

C0.2 X(12C)=0.2

C0.4 X(12C)=0.4

C0.4

C0.2

(Imbriani et al. ApJ 2001)

slide30

THE ROLE OF ROTATION

Cells of Meridional Circulation

Increasing rotation

Von Zeipel Theorem

OBLATENESS

GRATTON-

ÖPIK CELL

Advection of Angular Momentum

Shear Instabilities:

- Mixing of chemical species

- Transport (diffusion) of angular momentum

slide31

THE ROLE OF ROTATION

How include this multi-D phenomenon in a 1-D code?

Cylidrical Symmetry:

Average values over characteristic surfaces

1D problem

Isobars/Equipotentials

slide32

MAJOR UNCERTAINTIES IN THE SIMULATION OF THE EXPLOSION (REMNANT MASS – NUCLEOSYNYHESIS):

  • Prompt vs Delayed Explosion (this may alter both the M-R relation and Ye of the presupernova model)
  • How to kick the blast wave:
  • Thermal Bomb – Kinetic Bomb – Piston
  • How much energy to inject and where:
  • Thermal Bomb (Internal Energy)
  • Kinetic Bomb (Initial Velocity)
  • Piston (Initial velocity and trajectory)
  • How much kinetic energy at infinity
  • [SN(~1051 erg)/HN(~1052 erg)]
  • Extension and timing (before/after fallback) of mixing
  • Efficiency of n-process and changing of Ye
slide33

INDUCED EXPLOSION

Hypernova model without mixing-fallback

(25M, E51=10)

Normal SN model (25M, E51=1)

Hypernova model

(25M, E51=10, mixing-fallback, Ye)

Hypernova model with mixing-fallback

(25M, E51=10)

slide34

STRATEGIES FOR IMPROVEMENTS

  • Convection: hydrodynamical simulations in 3D  derive simple prescriptions to be used in 1D hydrostatic models (Arnett)
  • Rotation: implementation of 3D stellar models
  • 12C(a,g)16O: ask to nuclear physicists
  • Explosive Nucleosynthesis and Stellar Remnants: solve the “Supernova Problem”  improve the treatment of neutrino transport