Supernovae & Neutrinos - Technische Universität...

Supernovae & Neutrinos

Hans Peter Loens

Seminar: Relativistic Heavy-Ion-Physics

February the 2nd 2006

What is a Supernova?Some Preparations

How A Supernova Explodes (Type-II)What Was Missing? What Needs To Be Improved?

Table of Contents

1 What is a Supernova?The PhenomenonSN1987AClassification Of Supernovae

2 Some PreparationsThe Chandrasekhar-MassStellar Burning Stages

3 How A Supernova Explodes (Type-II)Single Steps Of A Type-II SupernovaThe Role Of The NeutrinosWhat We see...What Is Left Behind?

4 What Was Missing? What Needs To Be Improved?

Hans Peter Loens Seminar: Relativistic Heavy-Ion-Physics Supernovae & Neutrinos



The PhenomenonSN1987AClassification Of Supernovae

What Is A Supernova? - Phenomenological Approach

Figure: SN1994D in NGC 4526

Some Facts

huge stellar explosion!

extreme light-radiation (seenext slide!)

luminosity comparable to awhole galaxy

the fate for all stars withmass

M > 8 ·M�

M� = 1.98 · 1033g





Carrying Away That Much Energy

Extreme Conditions & Characteristics

high temperatures: > 1 MeV ≈ 1.16 · 1010K

high densities: 107 · gcm3 < ρ < 1015 · g

cm3

high luminosity: up to 1010 · L�huge size: radii ∼ 1013m ≈ 1.44 · 104 · R�

Relative Energy Distribution

total energy: ∼ 1053erg ≡ Egrav

(visible) radiation: set as 1 (∼ 1049erg)

kin. energy of the matter: 100 (∼ 1051erg =: 1Bethe)

neutrino energy: 10000 ≡ 99% of total energy ≡ 1053erg=⇒ this is why we concentrate on the neutrinos!





SN1987A

Figure: SN1987A

SN1987A-Facts

on 24.02.1987

on 23.02.1987⇒ neutrino det.

@GreatMagellanicCloud

distance 50kPc

M ≈ 20 ·M�





Supernova Classification




The Chandrasekhar-MassStellar Burning Stages

The Chandrasekhar-Mass

White Dwarf: fate of a star with M < 8 ·M�

degenerate electron gas!

the degenerate electron gas creates a pressure that ’presses’against the gravitational force=⇒ star only collapses to a minimum size

Definition

The Chandrasekhar-Mass is the largest mass of a star that canwithstand gravitational collapse by degenerate electron pressure:

MCh ≈ 1.457 · (2Ye)2 ·M�

with Ye = # of electrons# of nucleons as the ’electron abundance’


Figure: evolution of a star which will end in a Supernova



Single Steps Of A Type-II SupernovaThe Role Of The NeutrinosWhat We see...What Is Left Behind?

A Schematic Overview





Pre-Supernova Situation

above table: evolution of a 15 ·M� star

especially note the time scale, temperature, density &neutrino loss

the fusion stops at iron (= Fe has the highest binding energyper nucleon =⇒ no fusion beyond iron!)

’sudden’ stop of fusion in the Fe zone (= no moretemperature output) has dramatic consequences!





The Onion-Structure

Collapse

But if MCore > MCh ≈ 1.4 ·M�=⇒ Fe-core becomes unstable=⇒ core-collapse starts!





Why The Electrons Cannot Withstand

The Processes

There are basically two processes which lead to the core-collapse:

1 Si-burning creates Fe, which is added to the core ⇒ mass ofcore increases

2 electron capture on nuclei:

e− + ANXZ −→ A

N+1XZ−1 + νe

thus Ye changes ⇒ MCh ≈ 1.457 · (2Ye)2 ·M� decreases

EC On Nuclei More Important!

Although the cross-section for EC on nuclei is smaller than on freeprotons, the EC on nuclei is more important because theabundance of free protons is low!





Figure: EC on nuclei dominates relative to EC on protons





What Happens Next?

the core collapses with ≈ 1/4 · cthe core collapses to a hot, dense, neutron-rich sphere=⇒ Proto-Neutron-Star (PNS)

when reaching nuclear-matter densities, the repulsive nuclearforce comes into play

this happens at densities ∼ 1014g/cm3





Collapse-Phase - Properties

Properties/Characteristics

high temperature: reactions nominated by1 strong interaction2 electromagnetic interaction

are in equilibrium ⇒ matter composition in NSE

reactions nominated by weak interaction are NOT inequilibrium

the neutrino interaction with matter must be considered indetail (Boltzmann transport!)





Collapse Phase - Homologous Core





Collapse Phase - Shockwave Formation

@nuclear-matter densities: sudden strong resistance againstgravitation

⇒ sound waves propagate outwards (analogon: vibrations of ahammer hitting an anvil)

these waves stop @sonic-point, while more material keeps onfalling on the core

sound waves collect @sonic-point ⇒ pressure rises there!

=⇒ discontinuous change in velocity=⇒ a shock wave forms which travels outside!

Important

This shock wave is the main ’force’ for the whole explosion. If theshock wave has enough energy to reach the outer burning shells ofthe star, the star will explode!





Core-Collapse Again

Figure: schematic figure of the first stage of the core collapse





The Neutrino-Trapping

With increasing density another important effect comes into play:

Neutrino-Trapping

ρ ∼ 1012g/cm3 the neutrinos are trapped due to elastic scattering!

mean free path: λ = 1ρσ = 1

ρ12·(

108cm([N2/6A]Xnuclei+Xfree neutrons)·E2

ν

)with N = #neutrons

nucleus , A = #nucleonsnucleus , Xi = mass fraction

from this follows: λ ≈ 0.4km - compared to a diffusion lengthD ≈ 8km=⇒ neutrinos cannot leave core in the relevant time-scale

Important

The neutrinos form a Fermi gas!

This effect is very important! Due to the trapping the backward EC comesinto play and thus the EC-rate becomes equilibrated with its backwardreaction!=⇒ Ye can only fall to a minimum value =⇒ thus a significant amount ofelectrons & protons remains!





Neutrino Reactions

elastic scattering:ν + A � ν + A

absorption:νe + (N,Z ) �e− + (N − 1,Z + 1)

ν-e-scattering:ν + e− � ν + e−

inelastic scattering:ν + A � ν ′ + A∗

cross-sections: ∼ E 2ν





Core Bounce

The Bounce

ρ ' (4− 8) · 1014g/cm3

’Point of maximumscrunch’

pressure by the nucleons

the sound-waves traveloutwards=⇒ shock wave

inside the shock: Tincreases





Neutrino Burst

What Does The Shock Wave Do?

The energy of the shock wave is about 1051erg

while travelling outwards, the shock dissociates the Fe-nuclei

The Burst

huge amount of free protons (after shock passage)

huge EC rates (EC on Fe-bound protons much more unlikelythan on free proton - large threshold!)

Neutrinosphere = min. radius where neutrinos are not trapped

sudden neutrino burst from regions beyond theNeutrinosphere!





Neutrino Burst - Graphics





Shock Scenarios

Definition

Depending on the energy of the shock wave, there are differentoutcomes for the shock:

1 prompt shock scenario: the shock has enough energy to reachthe Si-burning layers, there he gains energy (due to theburning) and disrupts the star

2 delayed shock scenario: the shock does not have enoughenergy ⇒ the shock stallssome mechanism(s) transfer(s) energy to the shock for theSupernova to occur!

Unlikely Prompt Shock

Yet, the ’Prompt Shock Scenario’ is unlikely to happen - might onlyhappen in less massive stars - but why?





Shock Stagnation

Due To The Dissociation...

=⇒ shock wave candissociate ≈ 0.6 ·M�core envelope:Menv = Mcore −Mhom

newest estimation:Menv ≈ 1 ·M�

=⇒ shock wave cannot reachthe Si-shell

Accretion Shock

The shock wave stops andmatter from outer regionsfalls through it towards thecore!





Other Neutrino-Flavors

ν-Reactions Behind The Shcok

pair-annihilation: e− + e+ ↔ ν + ν̄neutrino-bremstrahlung: N + N → N + N + ν + ν̄EC/PC: νe + n ↔ p + e− and ν̄e + p ↔ n + e+

µ’s & τ ’s ...

the temperature in the core behind the shock is very high(≈ 1011K )many e−e+-pairs exist (because of the deleptonization theelectron-degeneracy decreases)

rate of energy transfer: r ≈ 1025 · T 9 · (1 + 0.19n) erg/cm3

stwo other flavours besides eletron-flavour: n=2(cannot be created by charged current interaction)Note: Rν different for the neutrino types!f.ex., the matter is neutron-rich: Rν̄e < Rνe





Reviving The Shock

Crucial Reactions

The most important reactions just behind the shock are

νe + n ↔ p + e−

ν̄e + p ↔ n + e+

=⇒ competition: emisson (cooling) VS. absorption (heating)

Definition

due to the different ρ & T profiles, there is a gain-radius Rg :R < Rg : emission dominates!R > Rg : absorption dominates!=⇒ the matter directly behind the stalled shock wave gets heatedby neutrinos!





KABOOOOOM!

The Shock Wave Is Revived!

Within ≈ 0.4s the Neutrinoshave deposited enough energyfor the shock wave to make its

way through the rest of thecore!

Which Model?

Nowadays, it is believed, thatthe delayed shock model isthe ’main’ model forsupernovae.





Supernova Light

the light comes from radioactive decays of nuclei created inthe SN

by scattering with free electrons, these photons arethermalized and radiated away in IR-, UV- & visible-light

main part by 56Ni (0.15 - 0.2 solar masses)

56Ni −→ 56Co in 6 days

56Co −→ 56Fe in 77 days


Figure: Massive Stars Remnants; from Heger et al., Astrophy.J. 591



What I left out & what are future tasks?

convection: in the newer models, the role of convection is verycrucial: some think that convection has an essential rolebringing hot material to the stalled shock! (needed:huge entropy gradient)

symmetry: usually supernovae do not have a spherical symmetry -this is quite sure, because of the observation that theremnant neutron star gets a kick to a high velocity(102 − 103 km

s ) by the explosion=⇒ the symmetry is now being accepted as veryimportant concerning the simulations!

magnetic fields: might support supernovanuclear physics: much needs to be done, especially the neutrino

energy transportneutrino oscillations: their role is still widely discussedgravitational waves: from supernovae are also being studiedγ-ray bursts: supernovae are thought to be their sources




S. Woosley & H.-Th. Janka, The Physics of Core-Collapse Supernovae, Nature Physics 1(2005) 147http://arxiv.org/abs/astro-ph/0601261

Hans A. Bethe, Supernovae, Physics Today Sept. 1990http://www.physicstoday.org/vol-43/iss-9/vol43no9p24_27.pdf

Hans A. Bethe & G. Brown, How a Supernova Explodes, Scientific American 252, 1985

H.-Th. Janka, Astroteilchenphysik in Deutschland - Workshop 2003http://www-ik.fzk.de/workshop/atw/Session7/Janka/janka.pdf

K. Kotake, K. Sato & K. Takahashi, Explosion Mechanism, Neutrino Burst andGravitational Wave in Core-Collapse Supernovaehttp://www.heap.phys.waseda.ac.jp/kkotake/IOP_paper.html

A. Burrows, S. Reddy & T.A. Thompson, Neutrino Opacities in Nuclear Matterhttp://arxiv.org/abs/astro-ph/0404432

K. Langanke, lecture: Nuclear Astrophysics at TU Darmstadt 05/06http://theory.gsi.de/~langanke/


http://arxiv.org/abs/astro-ph/0601261

http://www.physicstoday.org/vol-43/iss-9/vol43no9p24_27.pdf

http://www-ik.fzk.de/workshop/atw/Session7/Janka/janka.pdf

http://www.heap.phys.waseda.ac.jp/kkotake/IOP_paper.html

http://arxiv.org/abs/astro-ph/0404432

http://theory.gsi.de/~langanke/

Supernovae & Neutrinos - Technische Universität...

Documents

Transcript of Supernovae & Neutrinos - Technische Universität...