Evolution of a star

1

gas cloud - free fall - ionization

protostar - hydrostatic

equilibrium - convective

White dwarf H burning

radiative

Main sequence - pp chain

End of MS - H burning shell

Main sequence - CNO cycle

radiative

convective

convective

End of MS - H burning shell

H. Karttunen et al., Fundamental

astronomy

Evolution of a star

2

Red giant phase - Expansion - 3α - He flash

He burning expansion

planetary nebula

white dwarf

He burning

He burning

He burning (shell)

He burning (shell)

C flash

supernova

C,Ne,O,Si burning

supernova

Remnant: neutron star or

black hole

Possibly total

destruction of the star

H. Karttunen et al., Fundamental astronomy

Endpoints of stellar evolution

3

The end of stellar evolution: an inert core of spent fuel that cannot maintain gas pressure to balance gravity

Chandrasekhar Mass:

In more massive cores: • electrons become relativistic • gravitational collapse

Θ≈ MYM eCh285.5

For N=Z : MCh=1.46 M0

M

If M < MCh: core can be balanced against gravitational collapse by electron degeneracy pressure

Mass and composition of the core

4

Depend on the ZAMS mass and the previous burning stages:

0.3- 8 M0 He burning C,O

8-12 M0 C burning O,Ne,Mg

> 8-12 M0 Si burning Fe

MZAMS Last stage Core

M<MCh core survives

M>MCh collapse

Mass Result

< 0.3 M0 H burning He

How can 8-12M0 mass star get below Chandrasekhar limit ?

Death of a low mass star: a “Planetary Nebula”

5

image: HST Little Ghost Nebula distance 2-5 kLy blue: OIII green: HII red: NII

And here’s the core ! a “white dwarf”

Envelope of star blown into space

Why “white dwarf” ?

6

• core shrinks until degeneracy pressure sets in and halts collapse

star is HOT (gravitational energy !)

star is small

WD M-R relation Hamada-Salpeter Ap.J. 134 (1961) 683

3/1~ −MR

White dwarfs: magnitudes

7

Perryman et al. A&A 304 (1995) 69 HIPPARCOS distance measurements

nearby stars:

Where are the white dwarfs ?

there (small but hot white (B~V))

White dwarfs in the L vs T plot

8

The Stellar Onion

9 http://cococubed.asu.edu/pix_pages/87a_art.shtml

Inner structure of a presupernova star

10

Rolfs&Rodney Fig. 8.10

Supernovae

11

MCORE > MCh : collapse supernova explosion at least the outer part of a star is blown off into space

But why would a collapsing core explode ?

a) CO or ONeMg cores that accrete matter from a companion star can get beyond the Chandrasekhar limit:

Further collapse heats star and CO or ONeMg burning ignites explosively

Whole star explodes – no remnant

b) collapsing Fe core in massive star (but not too massive) neutron star

Fe cannot ignite, but collapse halted once densities of ~2x nuclear density are reached (repulsive nuclear force)

Core collapse supernova mechanism

12

Fe core

inner core

pre SN star 1.

infalling outer core

outgoing shock from rebounce

proto neutron star 2.

infalling outer core proto neutron star

stalled shock

3.

revived shock

proto neutron star

matter flow gets reversed - explosion

4.

neutrinos

neutrino heated layer

Supernovae

13

1. Luminosity: might be the brightest objects in the universe can outshine a whole galaxy (for a few weeks) Energy of the visible explosion: ~1051 ergs (= 1 foe = 1 Bethe) Luminosity : ~109-10 L0 2. Frequency: ~ 1-10 per century and galaxy 3. Observational classes (types): Type I: no hydrogen lines

-depending on other spectral features there are sub types Ia, Ib, Ic, ... Type II: hydrogen lines

Why are there different types ? Answer: progenitor stars are different Type I: two possibilities: Ia: white dwarf accreted matter from companion Ib,c: collapse of Fe core in star that blew its H (or He) envelope into space prior to the explosion Type II: collapse of Fe core in a normal massive star (H envelope)

Mass of a star effects the abundances

14

Rolfs&Rodney Fig. 8.11

Lower mass star thinner layer of

heavier elements around the Fe/Ni core

different abundances

Yield function: average over

• stars of different masses • number of stars of a mass M • the amount of ejected matter

contribution to Solar System abundances

Supernova light curves

Plateau !

15

Origin of plateau:

H-envelope outer part: transparent (H) inner part: opaque (H+)

photosphere

earlier: later: As star expands, photosphere moves inward along the T=5000K contour (H-recombination) T, R stay therefore roughly fixed = Luminosity constant (as long as photosphere wanders through H-envelope)

Supernova SN1987A

16

- in the Large Magellanic Cloud, a nearby dwarf galaxy

- February 1987

- Type II

supernova

- progenitor star: Sanduleak -69° 202 (a blue supergiant with M~20MSun)

- no neutron star remnant has been found

Evolution of the star that became Supernova 1987a

17 http://cococubed.asu.edu/pix_pages/87a_art.shtml

Supernova light curve

18

56Ni: T½~6 days 56Co: T½~ 77 days

Note! About two to three hours before the visible light from SN 1987A reached the Earth, a burst of neutrinos was observed at three separate neutrino observatories SN1987A: the first time neutrinos emitted from a supernova had been observed directly

Supernova 1987A seen by Chandra X-ray observatory, 2000

19 Shock wave hits inner ring of material and creates intense X-ray radiation

Crab Nebula in Taurus

20

• a supernova remnant • corresponds to the

bright SN 1054 supernova

The nitrogen in our DNA, the calcium in our teeth, the iron in our blood, the carbon in our apple pies were made in the interiors of collapsing stars. We are made of starstuff. -Carl Sagan

Cas A supernova remnant

21

… seen over 17 years

youngest supernova in our galaxy – possible explosion 1680 (new star found in Flamsteeds catalogue)

Light curves and radioactivity

22

(Frank Timmes)

Radioactive isotopes are produced during the explosion there is explosive nucleosynthesis !

44Ti

23

59.2+-0.6 yr

3.93 h

1157 γ-ray

Galactic longitude and latitude

24

Galactic Longitude l

Galactic Latitude b

Galactic coordinates

25

North pole: latitude b = +90 degrees

Atlas Image [or Atlas Image mosaic] courtesy of 2MASS/UMass/IPAC-Caltech/NASA/NSF. South pole:

latitude b = -90 degrees

Galactic center: l=0 degrees

l=+180 degrees l=-180 degrees

44Ti gamma rays

26

Galactic longitude (deg)

Gal

actic

latit

ude(

deg)

Log likelihood ratio

Cross = location of Cassiopeia A

1157 keV gamma ray from 44Ti

27 Iyudin et al. 1997

Measure the half-life of 44Ti

28

It’s not so easy: Status as of 1997:

Method 1:

29

Prepare sample of 44Ti and measure activity as a function of time

teNtN λ−= 0)(number of sample nuclei N:

activity = decays per second:

teNtNtA λλλ −== 0)()(

Measure A with γ-ray detector as a function of time A(t) to determine N0 and λ

2/1

2lnT

=λ

Half-life from Berkeley:

30

Norman et al. PRC57 (1998) 2010

T1/2=59.2 yr

Half-life from ANL

31 Ahmad et al. PRL 80 (1998) 2550

Method 2: measure A and N0 at the same time

32

teNtNtA λλλ −== 0)()(Measure Α AND N0 at a one time

44Ti

Cyclotron Pulse

Time of flight

Use this setup from time to time:

44Ti

Standard Setup:

energy loss dE

Si detector Plastic det.

NSCL, MSU

33

Cyclotron 1 Cyclotron 2

Ion Source

Fragment Separator

Make 44Ti by fragmentation of 46Ti beam

46Ti/s 106/s 44Ti 1010

Selectivity

34

TOF Stop (fast scintillator)

TOF Start (fast scintillator)

Energy loss dE (Si-PIN diode or ionization chamber)

Bρ selection by geometry/slits and fields

Bρ = mv/q (relativistic Bρ=γmv/q !) m/q = Bρ/v

dE ~ Z2

v=d/TOF

measure m/q: Measure Z:

New half-life value

35

determine number of implanted 44Ti

60.3 +- 1.3 years

Goerres et al. Phys. Rev. Lett. 80 (1998) 2554

Explosive Nucleosynthesis

36

Explosive Si burning: Deepest layer: full NSE 28Si 56Ni Further out: α-rich freezeout • density low, time short 3α cannot keep up and α drop out of NSE (but a lot are made from 2p+2n !) • result: after freezeout lots of α ! • fuse slower – once one 12C is made quickly captures more result: lots of α-nuclei (44Ti !!!)

Explosive C-Si burning • similar final products • BUT weak interactions unimportant for >= Si burning (but key in core !!!) • BUT somewhat higher temperatures • BUT Ne, C incomplete (lots of unburned material)

composition before and after core coll. supernova:

mass cut somewhere here not ejected ejected

Shock wave rips through star and compresses and heats all mass regions

The “mass zones” in “reality”

37 1170s after explosion, 2.2*106 km width, after Kifonidis et al. Ap.J.Lett. 531 (2000) 123L

Contribution of Massive Stars to Galactic Nucleosynthesis

38 calculation with grid of massive stars 11-40M0 (from Woosley et al. Rev. Mod. Phys. 74 (2002)1015)

Overproduction factor X/Xsolar: the fraction of matter in the Galaxy that had to be processed through the scenario (massive stars here) to account for todays observed solar abundances. To explain the origin of the elements one needs to have

• constant overproduction (then the pattern is solar) • sufficiently high overproduction to explain total amount of elements observed today

“Problem” zone: these nuclei are not produced in sufficient quantities

Type Ia supernovae

Novae

low mass stars

Supernova remnants – neutron stars

39 SN remnant Puppis A (Rosat)

Neutron star kicked out with ~600 mi/s

An isolated neutron star seen with HST (HST=Hubble Space Telescope)

40

Its estimated that there are ~100’s of millions of neutron stars in our Galaxy

Neutron star properties

41

Mass:

Radius: about 10 km!

Compare to the Sun’s radius: 696 342 km!

Neutron star: composition

42

Type Ia supernovae

43

white dwarf accreted matter and grows beyond the Chandrasekhar limit

star explodes – no remnant

Discovery rate of type Ia supernovae

44

SN cosmology “super-nova”

D. Kasen, presentation

Nucleosynthesis contribution from type Ia supernovae

45

(Pagel 5.27)

Iron/Nickel Group

CO or ONeMg core ignites and burns to a large extent into NSE

Has to be consistent with solar abundances Nucleosynthesis is a prime constraint for models

Nucleosynthesis beyond iron (A > 60)

46

Rolfs&Rodney Fig. 9.1

Abundance pattern: • Fusion reactions no longer possible (B/A curve: max around Fe) • Charged particle captures: Coulomb barrier inhibits abundances should be very low… • Most probable explanation: neutron captures!!

Where to get neutrons? • Helium-burning or in red giants :

13C(α,n)16O, 22Ne(α,n)25Mg

• High neutron densities (probably) in supernova explosions

Supporting facts

• Only 3 % of the iron-peak elements are needed to synthesize all heavy elements (enough seed nuclei)

• Neutron capture cross sections of heavy elements are very large (high level density)

• Observed Tc lines from S-type stars (red giants) ongoing nucleosynthesis

47

Neutron captures

(Z,A) + n → (Z,A+1) + one or more γ−rays (Z,A+1) → (Z+1,A+1) + e- + νe (β- decay)

48

(n,γ)

β-

• Neutron capture can be followed by β- decay • Depends on the capture reaction rate λn and the beta-decay rate λβ

• Neutron captures occur until λn < λβ

• Two different neutron capture processes: rapid (r) and slow (s) processes

A-1 A A+1

(n,γ)

A-1

A A+1

Z+1

Z Z

β-

(n,γ)

A-1

A A+1

Z+1

Z

Branching!

Neutron capture cross sections

• Neutrons quickly thermalized through eleastic scattering

Maxwell-Boltzmann distribution most probable energy near E=kT

most probable thermal velocity: 𝑣𝑣𝑇𝑇 = 2𝑘𝑘𝑇𝑇𝑚𝑚

where m = reduced mass • (α,n) reactions in red giants: T=0.1-0.6 GK E~30 keV

• 𝜎𝜎𝑛𝑛,𝛾𝛾 ∝1𝑣𝑣∝ 1

𝐸𝐸 σv = constant = σTvT

• Reaction rate/particle pair: <σv>=constant=<σ>vT

49

M-B distribution

50

Rolfs & Rodney Fig. 9.2

Neutron capture rate λ

51

𝜆𝜆 𝑛𝑛, 𝛾𝛾 = 𝑁𝑁𝑛𝑛 𝜎𝜎𝑣𝑣 = 𝑁𝑁𝑛𝑛 𝜎𝜎 𝑣𝑣𝑇𝑇

where Nn = neutron density (1/cm3) <σv> = average reaction rate per particle pair (cm3/s)

𝜏𝜏 𝑛𝑛, 𝛾𝛾 = 1/𝜆𝜆𝑛𝑛 = lifetime agianst neutron capture

• slow (s) neutron capture: τ(n,γ) >> τβ λ(n,γ) << λβ capture rates much slower than the beta decay • rapid (r) neutron capture: τ(n,γ) << τβ λ(n,γ) >> λβ capture rates much faster than the beta decay

Neutron capture cross sections

52

Rolfs&Rodney Fig. 9.6

• Dips at magic numbers • Increases for heavier nuclei

Level densities

s process

λn << λβ: • slow neutron capture process • “low”neutron density nn ~1014 1/m3

• Proceeds close to stable nuclei up to 209Bi • A>209: no stable or metastable nuclei

available for the captures (alpha decay etc)

53

209Bi + n → 210Bi + γ 210Bi → 210Po + e- + νe t½ = 5 d 210Po → 206Pb + α t½= 139 d

Equilibrium in the s process

54

dn(A)/dt ∝ n(A-1)σ(A-1)-n(A)σ(A) = 0

In equilibrium

• analogously n(A)σ(A) = n(A+1)σ(A+1) → n(A)σ(A) ~ constant • observed: ~ constant when A > 100

• In equilibrium, a nucleus (with mass number A) is produced and destroyed at the same rate:

s process: (n,γ) cross sections

55

σANA=σA-1NA-1= constant is OK!

Krane Lilley

s process path

56

Lilley

Branching at 176Lu

57

Branching ratio R = λβ/λ(n,γ)

T1/2 (176Lu) = 18.5exp(14.7/T8) hours for T8>1

s-process ”thermometer” kT ~18-42 keV

Rolfs&Rodney Fig 9.11

s process path

58

Rolfs&Rodney Fig. 9.7

Different processes and abundances

59 Rolfs&Rodney Fig. 9.8

r process

λn >> λβ: • rapid neutron capture process • Requires a high (temporal) neutron flux (~1032 m-2s-1) → supernovae, neutron star

mergers? • Several neutron captures before β- decay • along the neutron-rich nuclei • Continues until fission becomes more likely • Explains how nuclei with A>209 have beeen

formed (e.g. 232Th, 235U, 238U)

60

61

Krane: Fig. 19.14: • s process close to stable nuclei • r-process runs towards neutron-rich nuclei

T½(β-)=0.1 s

T½(β-)=0.2 s

62

Krane:Fig.19.16: • 120Sn: both s- and r-process produce • 122Sn and 124Sn: r process only • 122,123,124Te: s process only

63

Krane:Fig. 19.17: • at closed neutron shells (N=50,82,126): t½ (β-) is short • r-process ends when fission becomes more probable • possibility to produce superheavy nuclei

64

Krane: Fig. 19.18: • peaks corresponding to closed neutron shells (Fig.19.17) → β- decay • peaks at A=80,130 ja 195: r-process path at N=50,82,126 • peaks at A= 90, 138 ja 208: s-process stable nucleu for which N=50,82,126

r process abundances

65

Nr ~ Nsolar – Ns pronounced maxima A=80 A=130 A=195

Rolfs&Rodney Fig 9.12

s- and r-process paths

66

Rolfs&Rodney Fig 9.13

Waiting-point approximation

67

Iliadis Fig. 5.69

Waiting point nucleus

Sn (Qn,γ) values and waiting points

68

Iliadis Fig. 5.70 T=1.25 GK, Nn=1022 cm-3