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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))
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 !
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
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
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
Type Ia supernovae
43
white dwarf accreted matter and grows beyond the Chandrasekhar limit
star explodes – no remnant
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
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:
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
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