Physics 120 Nuclear and Hadronic Physics Prof' C'E' HydeWright

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Physics 120Nuclear and Hadronic PhysicsProf.
C.E. Hyde-Wright

• Lecture 1.
• Nucleo-Synthesis
• Sources from http//www.nscl.msu.edu/schatz/phy98
  3.html

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Triumphs of 20th Century Nuclear Astrophysics

• Why the Earth is hot
• Why the sun shines
• Nucleosynthesis in stars
• Big-Bang Nucleosynthesis
• Neutrinos from the sun
• Solar neutrinos and the lack there-of.
• Supernovae
• Nucleosynthesis
• Neutrinos

Sun (41026 Watt) (5 109 yr)
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Galactic Radioactivity - detected by g-radiation
1 MeV-30 MeV g-Radiation in Galactic Survey
(26Al Half life 700,0000 years)
Nucleosynthesis happens!
44Ti in Supernova Cas-A Location
(Half life 60 years)
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Atomic Nuclei, 4He, 56Fe, 238U AZ

• Each atom of a given element has a nucleus of
  charge Ze.
• He, Z2 Fe, Z26 U, Z92
• Each element can have multiple isotopes with mass
  number A. Each isotope has Z protons and NA-Z
  neutrons.
• The strong force binds the protons neutrons
  together
• (and partially dissolves themsee lecture 3).
• E mc2 Each nucleon (n,p) has mass 938 MeV/c2.
  Typical binding energy in nuclei is 8 MeV per
  nucleon or 0.9 of total mass.
• Compare with binding energy of Hydrogen atom 13
  eV/ 938 MeV 10-8

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III. Nuclear Physics that determines the
properties of the Universe
Part I Nuclear Masses
1. Energy generation
nuclear fusion reaction A B C
if mAmB gt mC then energy Q(mAmB-mC)c2 is
generated by reaction
Q-value Q Energy generated (gt0) or consumed
(lt0) by reaction
2. Stability
with Qgt0 ( or mAgt mB mC)
if there is a reaction A B C
then decay of nucleus A is energetically possible.
nucleus A might then not exist (at least not for
a very long time) Decays AZ ? AZg or AZ ?
A-1Zn AZ ? A-2(Z-2) a or fission
3. Equilibria
for a nuclear reaction in equilibrium abundances
scale with e-Q
(Saha equation)
Masses become the dominant factor in determining
the outcome of nucleosynthesis
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2. Nucleons
size 1 fm
n ? p e- ne T1/2 10.4 minutes
3. Nuclei
nucleons attract each other via the strong force
( range 1 fm)
a bunch of nucleons bound together create a
potential for an additional
neutron
proton(or any other charged particle)
Density of all nuclei is ?constant 0.17/fm3
V
V
Coulomb Barrier Vc
R 1.3 x A1/3 fm
Potential
Potential
R
R
r
r
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4. Nuclear Masses and Binding Energy
Energy that is released when a nucleus is
assembled from neutrons and protons
mp proton mass, mn neutron mass, m(Z,N)
mass of nucleus with Z,N

• Bgt0
• With B the mass of the nucleus is determined.
• B is roughly A (8MeV/nucleon)

Masses are usually tabulated as atomic masses
m mnuc Z me Be
Nuclear Mass 1 GeV/A
Electron Binding Energy13.6 eV (H)to 116 keV
(K-shell U) / Z
Electron Mass 511 keV/Z
Most tables give atomic mass excess D in MeV
(mu defined as mass of C12/12, so for 12C D0)
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Nucleosynthesis in Supernovae explosions
Nucleosynthesis in final stages of a star before
Supernovae (or white dwarf)
Fe, Ni, Most stable nuclei
Unstable to a-decay, fission Heats the earth
Fusion of protons to 4He Drives the sun
Emc2
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b decay basically no barrier -gt if energetically
possible it usually happens
(except if another decay mode dominates)
therefore any nucleus with a given mass number A
will be converted into the most
stable proton/neutron combination with mass
number A by b decays
Even-even more stable than even-odd
(Bertulani Schechter)
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Typical part of the chart of nuclides
red proton excessundergo b decay
37 Cl Ray Davis Neutrino detection
blue neutron excessundergo b- decay
Z
N
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The pp-I chain
Step 1
The weak interaction regulates Nucleo synthesis
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To summarize the pp-I chain
On chart of nuclides
3He
4He
2
1H
2H
1
1
2
Or as a chain of reactions
bottle neck
(p,g)
(3He,2p)
(p,e)
1H
d
3He
4He
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3He has a much higher equilibrium abundance than
d
- therefore 3He3He possible
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Hydrogen burning with catalysts

• ppII chain
• ppIII chain
• CNO cycle

1. ppII and ppIII
once 4He has been produced it can serve as
catalyst of the ppII and ppIII chainsto
synthesize more 4He
out
in
(e-,n)
(p,4He)
(4He,g)
3He
7Be
7Li
4He
ppII (sun 14)
(b n)
(p,g)
decay
8B
8Be
24He
ppIII (sun 0.02)
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(Rolfs and Rodney)
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Summary pp-chains
ppI
ppI ppII ppIII
7Be
8Be
6Li
7Li
Why do additional pp chains matter ?
3He
4He
pp dominates timescale
2
1H
2H
1
1
2
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CNO cycle
Ne(10)
F(9)
O(8)
N(7)
C(6)
3
4
5
6
7
8
9
neutron number
All initial abundances within a cycle serve as
catalysts and accumulate at largest t
Extended cycles introduce outside material into
CN cycle (Oxygen, )
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Competition between the p-pchain and the CNO
Cycle
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Neutrino emission
ltEgt0.27 MeV
E0.39,0.86 MeV
ltEgt6.74 MeV
ppIII loss 28
ppII loss 4
ppI loss 2
note ltEgt/Q0.27/26.73 1
Total loss 2.3
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2 neutrino energies from 7Be electron capture ?
7Be e- ? 7Li ne
Mono-energetic neutrinos (initial electron has
zero energy)
En
En
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Continuous fluxes in /cm2/s/MeV Discrete fluxes
in /cm2/s
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Astronomy Picture of the Day June 5, 1998
Nucleosynthesis produces neutrinos. The center of
the sun has never been seen before!
Neutrino image of the sun by Super-Kamiokande
next step in neutrino astronomy
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Big Bang Nucleosynthesis

• http//teacher.nsrl.rochester.edu/Research/PHOBOS/
  Presentations/GeneralTalks/APSNYSS102000/sld021.ht
  m
• Equilibrium density of 1H, 2H, 3He, 4He, Li
  determines mean density of baryonic matter after
  big bang 0.02 of critical density necessary to
  close universe (re-collapse).
• How to measure elemental isotopic abundances?

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Solar Absorption Spectra
Celia Payne, Harvard Ph.D. 1925, The sun is
mostly Hydrogen (Published over the objection of
her advisor who was convincedwith The rest of
the astronomy communitythat the sun was made of
iron)
solar spectrum (Nigel Sharp, NOAO)
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Use carbonaceous chondrites (6 of falls)
Chondrites Have Chondrules - small 1mm size
spherical inclusions in matrix believed to
have formed very early in the pre-solar nebula
accreted together and remained largely
unchanged since thenCarbonaceous Chondrites
have lots of organic compounds that indicate
very little heating (some were never heated
above 50 degrees)
Chondrule
How find them ?
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a-nuclei12C,16O,20Ne,24Mg, . 40Ca
GapB,Be,Li
general trend less heavy elements
r-process peaks (nuclear shell closures)
s-process peaks (nuclear shell closures)
U,Th
Fe peak(width !)
Fe
Au
Pb
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Challenges for the 21st Century

• Solve the solar neutrino puzzle ?
• Neutrino masses,
• Neutrino spectroscopy of the universe
• Proton decay limits
• Dark matter searches (WIMPS)
• How do Supernovae explode
• Hydrodynamics
• Rapid neutron-capture (r-process) synthesis of
  the heavy elements.
• What are neutron stars?
• Neutron matter, quark matter
• The stability of stars, synthesis of heavy
  elements, existence of life depend upon fine
  tuning of the parameters of standard model

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Neutrino Astronomy
Photons emitted from sun are not the photons
created by nuclear reactions (heat is
transported by absorption and emission of photons
plus convection to the surface over
timescales of 10 Million years)
But neutrinos escape !
Every second, 109 solar neutrinos pass through
your thumbnail !
But hard to detect (they pass through 1033 g
solar material largely undisturbed !)
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First experimental detection of solar neutrinos
1964 John Bahcall and Ray Davis have the idea to
detect solar neutrinos using the reaction

• 1967 Homestake experiment starts taking data

• 100,000 Gallons of cleaning fluid in a tank 4850
  feet underground
• 37Ar extracted chemically every few months
  (single atoms !) and decay counted in counting
  station (35 days half-life)
• event rate 1 neutrino capture per day !

• 1968 First results only 34 of predicted
  neutrino flux !

solar neutrino problem is born - for next 20
years no other detector !
Neutrino production in solar core T25
nuclear energy source of sun directly and
unambiguously confirmed
solar models precise enough so that deficit
points to serious problem
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37Cl tank in Homestake mine (S.D.), 1967 2001
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Are the neutrinos really coming from the sun ?
Water Cerenkov detector
high energy (compared to rest mass) - produces
cerenkov radiation when traveling in water (can
get direction)
nx
nx
neutral current (NC)
Z
e-
e-
Super-KamiokandeDetector
ne
ne
chargedcurrent (CC)
W-
e-
e-
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many more experiments over the years with very
different energy thresholds
Ratio of experimental Neutrino yield divided by
prediction for each detector
All show deficit to standard solar model
Mostly sensitive only to electron neutrinos.
ne only
all flavors, but
nt,nm only 16 of ne cross section becauseno
CC, only NC
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The solution neutrino oscillations
Neutrinos can change flavor while travelling from
sun to earth An electron neutrino is really a
superposition of two neutrino types of different
mass Each component oscillates at frequency
m/h ? beats!
The arguments
1. SNO solar neutrino experiment
uses three reactions in heavy water
Charged Current (Cerenkov)
CC
EC
Electron Capture (Cerenkov)
(n-capture by 35Cl - g scatter - Cerenkov)
NC
key

• NC independent of flavor - should always equal
  solar model prediction if oscillations explain
  the solar neutrino problem
• Difference between CC and ES indicates
  additional flavors present

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Sudbury Neutrino Observatory
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With SNO results
Puzzle solved
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more arguments for neutrino oscillation solution
2. Indication for neutrino oscillations in two
other experiments

• 1998 Super Kamiokande reports evidence for nm
  --gt nt oscillations for neutrinos created
  by cosmic ray interaction with the atmosphere

• 2003 KamLAND reports evidence for disappearance
  of electron anti neutrinos from reactors

3.
There is a (single) solution for oscillation
parameters that is consistent with all solar
neutrino experiments and the new KamLAND results
KamLAND
Reactor prouduces from beta decay of
radioactive material in core
Detection in liquid scintillator tank in
Kamiokande mine 180 km away
check whether neutrinos disappear
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2003 Results
dashed Best fit LMA sin22Q0.833, Dm25.5e-5
eV2shaded 95 CL LMA from solar neutrino data
K. Eguchi, PRL 90 (2003) 021802
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Explosive Nucleosynthesis
Pb (82)
Sn (50)
Fe (26)
protons
H(1)
neutrons
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What happens at hydrogen exhaustion
(assume star had convective core)
1. Core contracts and heats
H shell burning
H,He mix
He rich corecontracts andgrows from H-burning
He rich core
? red giant
2. Core He burning sets in
He core burning
? lower mass stars become bluer low Z stars
jump to the horizontal branch
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Helium burning 1 the 3a process
First step
a a ? 8Be
unbound by 92 keV decays back to 2 a within
2.6E-16 s !
but small equilibrium abundance is established
Second step
8Be a ? 12C would create 12C at excitation
energy of 7.7 MeV
1954 Fred Hoyle (now Sir Fred Hoyle) realized
that the fact that there is carbon in
the universe requires a resonance in 12C at 7.7
MeV excitation energy
1957 Cook, Fowler, Lauritsen and Lauritsen at
Kellogg Radiation Laboratory at Caltech
discovered a state with the correct properties
(at 7.654 MeV)
Experimental Nuclear Astrophysics was born
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Summary stellar burning
gt0.8M0
gt8M0
gt12M0
Why do timescales get smaller ?
Note Kelvin-Helmholtz timescale for red
supergiant 10,000 years, so for massive stars,
no surface temperature - luminosity changefor
C-burning and beyond
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Final composition of a 25 M0 star
up to H-burned
up to Ne-burned
up to Heburned
unburned
up to Siburned
up to Oburned
mass fraction
interior mass (M0)
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Endpoints of stellar evolution
The end of stellar evolution is an inert core of
spent fuel that cannot maintaingas pressure to
balance gravity
Such a core can be balanced against gravitational
collapse by electron degeneracypressure IF the
total mass is less than the Chandrasekhar mass
limit
Chandrasekhar Mass
Only if the mass of a inert core is less than
Chandrasekhar Mass Mch
Electron degeneracy pressure can prevent
gravitational collapse
In more massive cores electrons become
relativistic and gravitationalcollapse occurs
(then Pr4/3 instead of Pr5/3).
For NZ MCh1.46 M0
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Mass and composition of the core depends on the
ZAMS mass and the previous burning stages
(ZAMS zero age main sequence)
MZAMS
Last stage
Core
Mass
Result
lt 0.3 M0 H burning He
0.3- 8 M0 He burning C,O
MltMCh
core survives
8-12 M0 C burning O,Ne,Mg
gt 8-12 M0 Si burning Fe
collapse
MgtMCh
How can 8-12M0 mass star get below Chandrasekhar
limit ?
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Death of a low mass star a Planetary Nebula
image HSTLittle Ghost Nebuladistance 2-5
kLyblue OIIIgreen HII red NII
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core collapse supernova mechanism
Hard to do in simulation (rare in real life?)
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Some facts about Supernovae
1. Luminosity
Supernovae might be the brightest objects in the
universe, and can outshine a whole galaxy (for a
few weeks)
Energy of the visible explosion 1051
ergsLuminosity
109-10 L0
2. Frequency
1-10 per century and galaxy
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Tarantula Nebula in LMC (constellation Dorado,
southern hemisphere) size 2000ly (1ly 6
trillion miles), disctance 180000 ly
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Supernova 1987A seen by Chandra X-ray
observatory, 2000
Shock wave hits inner ring of material and
creates intense X-ray radiation
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Cas A supernova remnant
seen over 17 years
youngest supernova in our galaxy possible
explosion 1680 (new star found in Flamsteeds
catalogue)
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3. Observational classes (types)
no hydrogen lines
Type I
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 II
collapse of Fe core in a normal massive star (H
envelope)
Type I 2 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
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Plateau !
Origin of plateau
later
earlier
As star expands, photospheremoves inward along
theT5000K contour
(H-recombination) T,R stay therefore roughly
fixed Luminosity constant(as long as
photosphere wandersthrough H-envelope)
H-envelope
outer part transparent (H)
inner part opaque (H)
photosphere
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There is another effect that extends SN light
curves Radioactive decay !
(Frank Timmes)

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

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44Ti
59.2-0.6 yr
3.93 h
1157 g-ray
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44Ti decay chain
Distance 10,000 ly
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The mass zones in reality
1170s after explosion, 2.2Mio km width, after
Kifonidis et al. Ap.J.Lett. 531 (2000) 123L
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Type Ia supernovae
white dwarf accreted matter and grows beyond the
Chandrasekhar limit
star explodes no remnant
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Overview heavy element nucleosynthesis
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The r-process
Temperature 1-2 GK Density 300 g/cm3 (60
neutrons !)
neutron capture timescale 0.2 ms
Seed
Proton number
Neutron number
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show movies R-process movie_r2d_self.mov
rp-process movie_rp_self.mov http//groups.nscl.
msu.edu/nero/Web/materials.html
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The r-process path
r-process abundance distribution
r-processpath
RIA Reach
New MSU/NSCL Reach
Known
(Reach for half-life)
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National Superconducting Cyclotron Laboratory
atMichigan State University
New Coupled Cyclotron Facility experiments
since mid 2001
Ion Source86Kr beam
86Kr beam140 MeV/u
Tracking (Momentum)
TOF start
Implant beam in detectorand observe decay
86Kr hits Be target and fragments
TOF stop dE detector
Separated beamof r-processnuclei
Fast beam fragmentation facility allows event
by event particle identification
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Supernova remants neutron stars
Neutron star kicked out with 600 mi/s
SN remnant Puppis A (Rosat)
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Neutron star properties
Mass
Radius
10 km !
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