Transcript Document

The origin of heavy elements in the solar system
(Pagel, Fig 6.8)
each process contribution is a mix of many events !
1
Abundance pattern: “Finger print” of the r-process ?
Tellurium and
Xenon Peak
Abundance (Si = 106)
Solar abundance of the elements
Platinum Peak
r-process only
(subtract s,p
processes)
Element number (Z)
But: sun formed ~10 billion years after big bang: many stars contributed to elements
 This could be an accidental combination of many different “fingerprints” ?
 Find a star that is much older than the sun to find “fingerprint” of single event
2
Heavy elements in Metal Poor Halo Stars
recall:
[X/Y]=log(X/Y)-log(X/Y)solar
CS22892-052
red (K) giant
located in halo
distance: 4.7 kpc
mass ~0.8 M_sol
[Fe/H]= -3.0
[Dy/Fe]= +1.7
old stars - formed before Galaxy was mixed
they preserve local pollution from individual nucleosynthesis events
3
A single (or a few) r-process event(s)
CS22892-052 (Sneden et al. 2003)
1
solar r
Cosmo
Chronometer
0
-1
-2
40
NEW:
CS31082-001 with U
(Cayrel et al. 2001)
50
other, second
r-process to fill
this up ?
(weak r-process)
60
70
80
element number
main r-process
matches exactly solar r-pattern
conclusions ?
90
Age: 16+- 3 Gyr
(Schatz et al. 2002
ApJ 579, 626)
4
New Observations
r-process elements from single r-process events
in 3 very metal poor stars
Solar r-process
elements from
many events
J. Cowan
 Many more to come from ongoing surveys and followup campaigns (e.g. VLT)
5
Overview heavy element nucleosynthesis
process
conditions
timescale
site
s-process
(n-capture, ...)
T~ 0.1 GK
tn~ 1-1000 yr, nn~107-8/cm3
102 yr
and 105-6 yrs
Massive stars (weak)
Low mass AGB stars (main)
r-process
(n-capture, ...)
T~1-2 GK
tn ~ ms, nn~1024 /cm3
< 1s
Type II Supernovae ?
Neutron Star Mergers ?
p-process
((g,n), ...)
T~2-3 GK
~1s
Type II Supernovae
Light Element
Primary Process
(LEPP) ?
?
(maybe s-process like?)
?
(long if sprocess)
?
6
The r-process
Temperature: ~1-2 GK
Density: 300 g/cm3 (~60% neutrons !)
Rapid neutron
capture
neutron capture timescale: ~ 0.2 ms
b-decay
Seed
(g,n) photodisintegration
Equilibrium favors
“waiting point”
Neutron number
7
Waiting point approximation
Definition: ASSUME (n,g)-(g,n) equilibrium within isotopic chain
How good is the approximation ?
This is a valid assumption during most of the r-process
BUT: freezeout is neglected
Freiburghaus et al. ApJ 516 (2999) 381 showed agreement with dynamical models
Consequences
During (n,g)-(g,n) equilibrium abundances within an isotopic chain are given by:
Y ( Z , A  1)
G ( Z , A  1)  A  1 2 
 nn


Y ( Z , A)
2G ( Z , A)  A mu kT 
2
3/ 2
exp(S n / kT )
• time independent
• can treat whole chain as a single nucleus in network
• only slow beta decays need to be calculated dynamically
• neutron capture rate independent
(therefore: during most of the r-process n-capture rates do not matter !)
8
Pt
Xe
Ni
78Ni, 79Cu
first bottle necks in n-capture flow (80Zn later)
79Cu:
half-life measured 188 ms (Kratz et al, 1991)
78Ni : half-life predicted 130 – 480 ms
2 events @ GSI (Bernas et al. 1997)
9
H. Schatz
Nuclear physics in the r-process
b-delayed n-emission
branchings
(final abundances)
Fission rates and distributions:
• n-induced
• spontaneous
• b-delayed
b-decay half-lives
(abundances and
process speed)
n-capture rates
• in slow freezeout
• maybe in a “weak” r-process ?
n-phyiscs ?
Seed production
rates (aaa,aan, a2n, ..)
Masses (Sn)
(location of the path)
10
Sensitivity of r-process to astro and nuclear physics
Sensitivity to astrophysics
Sensitivity to nuclear physics
1
Abundance
Hot bubble
Classical model
Same nuclear physics
10
ETFSI-Q masses
ETFSI-1 masses
Same r-process model
0
10
-1
10
-2
10
-3
10
Freiburghaus et al. 1999
-4
10
Mass number
Contains information about:
• n-density, T, time
(fission signatures)
• freezeout
• neutrino presence
• which model is correct
100
120
140
160
180
200
220
Mass number
But convoluted with nuclear physics:
• masses (set path)
• T1/2, Pn (Y ~ T1/2(prog),
key waiting points set timescale)
• n-capture rates
• fission barriers and fragments
11
Shell quenching effect on masses/r-process
Ru
Pd
Cd
S2n (MeV)
Mo
Zr
r-process path
ETFSI-1
Neutron number
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Shell quenching effect on masses/r-process
Ru
Pd
Cd
Mo
Zr
S2n (MeV)
distinguish
ETFSI-1
ETFSI-Q
(N=82 quenched)
Neutron number
13
Endpoint of the r-process
r-process ended
by n-induced fission
or spontaneous
fission
(different paths
for different
conditions)
(Goriely & Clerbaux A&A 348 (1999), 798
n-induced fission
(Z,A)
n-capture (DC)
fission
(Z,A+1)
b-delayed fission
(Z,A)
spontaneous fission
bfission
(Z+1,A)
(Z,A+1)
fission barrier
fission
14
Consequences of fission
Fission produces A~Aend/2 ~ 125 nuclei
modification of abundances around A=130 peak
fission products can serve as seed for the r-process
- are processed again into A~250 region via r-process
- fission again
fission cycling !
Note: the exact endpoint of the r-process and the degree and impact of fission
are unknown because:
• Site conditions not known – is n/seed ratio large enough to reach fission ?
(or even large enough for fission cycling ?)
• Fission barriers highly uncertain
• Fission fragment distributions not reliably calculated so far (for fission from
excited states !)
15
Role of beta delayed neutron emission
Neutron rich nuclei can emit one or more neutrons during b-decay if Sn<Qb
(the more neutron rich, the lower Sn and the higher Qb)
(Z,A)
bn
Sn
g
(Z+1,A-1)
(Z+1,A)
If some fraction of decay goes above Sn in daughter nucleus
then some fraction Pn of the decays will emit a neutron (in addition to e- and n)
(generally, neutron emission competes favorably with g-decay - strong interaction !)
16
Effects: during r-process: none as neutrons get recaptured quickly
during freezeout • modification
:
of final abundance
• late time neutron production (those get recaptured)
Calculated r-process production of elements (Kratz et al. ApJ 403 (1993) 216):
before b-decay
after b-decay
smoothing effect from b-delayed n emission !
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Cs (55)
Cs131 Cs132 Cs133 Cs134 Cs135 Cs136 Cs137 Cs138 Cs139 Cs140 Cs141 Cs142 Cs143 Cs144 Cs145 Cs146 Cs147
> 99
Xe (54)
I (53)
Te (52)
Sb (51)
Sn (50)
In (49)
Cd (48)
Ag (47)
> 99
> 99
19.00 > 99
> 99
> 99
63.70 24.94 1.70
1.78
1.01
0.59
0.32
0.23
Xe130 Xe131 Xe132 Xe133 Xe134 Xe135 Xe136 Xe137 Xe138 Xe139 Xe140 Xe141 Xe142 Xe143 Xe144 Xe145 Xe146
> 99
I129
> 99
Pn39.68
=0%
13.60
> 99
> 99
I137
1.73
1.24
0.30
1.15
0.90
I142
I143
I144
I130
I131
I132
I133
I134
I135
I136
I138
I139
I140
I141
> 99
> 99
> 99
> 99
> 99
> 99
83.40 24.50 6.49
2.28
0.86
0.43
I145
r-process
waiting point
Te128 Te129 Te130 Te131 Te132 Te133 Te134 Te135 Te136 Te137 Te138 Te139 Te140 Te141 Te142 Te143 Te144
> 99
> 99
> 99
> 99
> 99
19.00 17.50 2.49
P =99.9%
1.40
n Sb134 Sb135 Sb136 Sb137 Sb138 Sb139 Sb140 Sb141 Sb142 Sb143
Sb127 Sb128 Sb129 Sb130 Sb131 Sb132 Sb133
> 99
> 99
> 99
> 99
> 99
> 99
> 99
1.66
0.82
Sn126 Sn127 Sn128 Sn129 Sn130 Sn131 Sn132 Sn133 Sn134 Sn135 Sn136 Sn137 Sn138 Sn139 Sn140 Sn141 Sn142
> 99
> 99
> 99
> 99
> 99
56.00 39.70 1.20
1.12
In125 In126 In127 In128 In129 In130 In131 In132 In133 In134 In135 In136 In137 In138 In139 In140 In141
2.36
1.60
1.09
0.84
0.61
0.26
0.28
0.20
Example: impact
of Pn for 137Sb
Cd138 Cd139 Cd140
0.18
Cd124 Cd125 Cd126 Cd127 Cd128 Cd129 Cd130 Cd131 Cd132 Cd133 Cd134 Cd135 Cd136
1.24
0.65
0.51
0.43
0.34
0.27
0.20
A=136
( 99%)
Ag123 Ag124 Ag125 Ag126 Ag127 Ag128 Ag129 Ag130 Ag131 Ag132 Ag133
0.29
0.17
0.16
0.10
0.11
0.06
0.05
76
77
78
79
80
81
82
83
84
85
86
A=137
( 0%)
Ag135 Ag136 Ag137 Ag138 Ag139
87
88
89
90
91
92
r-process waiting point
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Summary: Nuclear physics in the r-process
Quantity
Effect
Sn
neutron separation
energy
path
T1/2
b-decay half-lives
• abundance pattern
• timescale
Pn
b-delayed n-emission
branchings
final abundance pattern
• endpoint
• abundance pattern?
• degree of fission cycling
fission (branchings
and products)
G
partition functions
• path (very weakly)
NA<sv>
neutron capture rates
• final abundance pattern
during freezeout ?
• conditions for waiting point
approximation
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National Superconducting Cyclotron Laboratory at
Michigan State University
New Coupled Cyclotron Facility – experiments since mid 2001
Ion Source:
86Kr beam
86Kr
beam
140 MeV/u
86Kr
hits Be
target and
fragments
Tracking
(=Momentum)
TOF start
Separated beam
of r-process
nuclei
TOF stop
dE detector
Implant beam
in detector
and observe decay
Fast beam fragmentation facility – allows event by event particle identification
20
H. Schatz
NSCL Coupled Cyclotron Facility
W. Benenson (NSCL) and B. Richards (WKAR)
21
Installation of D4 steel, Jul/2000
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First r-process experiments at new NSCL CCF facility (June 02)
Measure:
• b-decay half-lives
• Branchings for b-delayed n-emission
Detect:
• Particle type (TOF, dE, p)
• Implantation time and location
• b-emission time and location
• neutron-b coincidences
New NSCL Neutron detector
NERO
3He
+ n -> t + p
neutron
Fast Fragment Beam
Si Stack
(fragment. 140 MeV/u 86Kr)
23
NSCL BCS – Beta Counting System
•
•
•
4 cm x 4 cm active area
1 mm thick
40-strip pitch in x and y
dimensions ->1600 pixels
Si
BCS
Si Si
b
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NERO – Neutron Emission Ratio Observer
3He
Proportional
Counters
BF3 Proportional
Counters
Specifications:
• 60 counters total
(16 3He , 44 BF3)
• 60 cm x 60 cm x 80 cm
polyethylene block
• Extensive exterior
shielding
• 43% total neutron
efficiency (MCNP)
Polyethylene
Moderator
Boron Carbide
Shielding
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NERO Assembly
26
Nero efficiency
NERO Efficiency vs. Neutron Energy
Efficiency (%)
50
13C
40
11B
30
20
51V
10
252Cf
0
0.001
0.01
0.1
1
Neutron Energy (MeV)
10
Scaled MCNP
Curve
27
Energy loss in Si (Z)
Particle Identification
Fast RIB from fragmentation:
• no decay losses
• any beam can be produced
• multiple measurements in one
• high sensitivity
r-process nuclei
78Ni
Doubly
Magic !
78Ni
75Co
77Ni
74Co
73Co
Time of flight (m/q)
28
H. Schatz
Decay data
time (ms)
time (ms)
time (ms)
Fast radioactive beams:
• No decay losses
• Rates as low as 1/day useful !
• Mixed beam experiments easy
29
Results for the main goal: 78Ni (14 neutrons added to stable Ni)
Decay of 78Ni : major bottle-neck for synthesis of heavy elements in the r-process
Managed to create 11 of the doubly magic 78Ni nuclei in ~ 5 days
Time between arrival and decays:
Statistical
Analysis
Result for half-life:
110 +100-60 ms
Compare to theoretical
estimate used:470 ms
time (ms)
 Acceleration of the entire r-process
 Models need to be adjusted to explain observed abundance distribution
30
Neutron Data
With neutron in addition
Nuclei with decay detected
420
420
370
DE (arb units)
DE (arb units)
Nn
320
76Ni
270
220
350
73Co
400
450
500
550
TOF (arb units)
Nn
Pn 
Nb n
370
320
76Ni
270
220
350
73Co
400
450
500
550
TOF (arb units)
neutron detection efficiency
(neutrons seen/neutrons emitted)
31
Results (Hosmer et al.)
DF+CQRPA Borzov et al. 2005,
QRPA: Moller et al. 2003,
Shell model: Lisetzky & Brown 2005
T1/2 (s)
A
A
Preliminary
Pn (%)
Preliminary
A
32
A
H. Schatz
Impact of 78Ni half-life on r-process models
Observed Solar Abundances
1.E+02
Abundance (A.U.)
Model Calculation: Half-Lives from
Moeller, et al. 97
Series4
Same
but with present 78Ni Result
1.E+01
1.E+00
1.E-01
1.E-02
70
120
170
220
Mass (A)
need to readjust r-process model parameters
Can obtain Experimental constraints for r-process models
from observations and solid nuclear physics
 remainig discrepancies – nuclear physics ? Environment ? Neutrinos ?
Need more data
33
NSCL and future facilities reach
Bright future for experiments and observations
 Experimental test of r-process models
is within reach Known half-life
 Vision: r-process as precision probe
NSCL reach
Reach of future facility
(here: ISF - NSCL upgrade
under discussion)
112Mo
Mo
111Nb
Nb
Zr
Y
108Zr
J. Pereira:
(NSCL)
105Y
Sr
Rb
34
Towards an experimental nuclear physics basis for the r-process
Final isotopes, for which >90% of progenitors in the r-process path can be reached
experimentally for at least a half-life measurement
today
Existing facilities
ISF
 These abundances can be compared with observations
to test r-process models
35
H. Schatz
Collaboration
78Ni Collaboration
MSU:
P. Hosmer
R.R.C. Clement
A. Estrade
P.F. Mantica
F. Montes
C. Morton
W.F. Mueller
E. Pellegrini
P. Santi
H. Schatz
M. Steiner
A. Stolz
B.E. Tomlin
M. Ouellette
Mainz:
O. Arndt
K.-L. Kratz
B. Pfeiffer
Pacific Northwest Natl. Lab.
P. Reeder
Notre Dame:
A. Aprahamian
A. Woehr
Maryland:
W.B. Walters
36
Overview of common r process models
• Site independent models:
• nn, T, t parametrization (neutron density, temperature, irradiation time)
• S, Ye, t parametrization (Entropy, electron fraction, expansion timescale)
• Core collapse supernovae
• Neutrino wind
• Jets
• Explosive helium burning
• Neutron star mergers
37
Site independent approach
Goal: Use abundance observations as general constraints on r-process
conditions
BUT: need nuclear physics to do it
nn, T, t parametrization (see Prof. K.-L. Kratz transparencies)
obtain r-process conditions
needed
for which the right N=50 and
N=82 isotopes are
waiting points
(A~80 and 130 respectively)
often in waiting point approximation
Kratz et al. ApJ403(1993)216
38
S, Ye, t parametrization
1.
2.
3.
Consider a blob of matter with entropy S, electron abundance Ye in NSE
Expand adiabatically with expansion timescale t
Calculate abundances - what will happen:
1.
2.
3.
4.
NSE
QSE (2 clusters: p,n,a and heavy nuclei)
a-rich freezeout (for higher S)
(3a and aan reactions slowly move matter from p,n,a cluster
to heavier nuclei – once a heavy nucleus is created it rapidly
captures a-particles
as a result large amounts of A~90-100 nuclei are produce
which serve as seed for the r-process
r-process phase
initially: n,g – g,n equilibrium
later: freezeout
39
Evolution of equilibria:
from Brad Meyers website
cross : most abundant nucleus
colors: degree of equilibrium with that nucleus
(difference in chemical potential)
40
Results for
neutron to seed ratios:
(Meyer & Brown ApJS112(1997)199)
n/seed is higher for
• lower Ye
(more neutrons)
• higher entropy
(more light particles, less
heavy nuclei – less seeds)
(or: low density – low 3a
rate – slow seed assembly)
• faster expansion
(less time to assemble seeds)
2 possible scenarios:
1) high S, moderate Ye
2) low S, low Ye
41
Matter evaporated off the hot neutron star
Neutron
star forms
(size ~ 10 km radius)
r-process site ?
42
How does the r-process work ? Neutron capture !
43
r-process in Supernovae ?
Most favored scenario for high entropy:
Neutrino heated wind evaporating from proto neutron star in core collapse
ne neutrino sphere (ne+p  n+e+ weak opacity
because only few protons present)
ne neutrino sphere (ve+n  p+e+ strong opacity
because many neutrons present)
proto
neutron star
(n-rich)
weak interactions regulate n/p ratio:
ne+p  n+e+
ne+n  p+e-
faster as ne come from deeper
and are therefore hotter !
therefore matter is driven
neutron rich
44
Results for Supernova r-process
Takahashi, Witti, & Janka A&A 286(1994)857
(for latest treatment of this scenario see Thompson, Burrows, Meyer ApJ 562 (2001) 887)
A~90 overproduction
density artificially reduced by
factor 5.5
can’t produce A~195 anymore
density artificially reduced by
factor 5
artificial parameter to get
A~195 peak (need S increase)
other problem: the a effect
45
other problem: the a effect
Recall equilibrium of nucleons in neutrino wind:
ne+p  n+e+
ne+n  p+e-
Maintains a slight neutron excess
np
n p  nn
 0.4
What happens when a-particles form, leaving a mix of a-particles and neutrons ?
46
r-process in neutron star mergers ?
47
Ejection of matter in NS-mergers
Rosswog et al. A&A 341 (1999) 499
Destiny of Matter:
red: ejected
blue: tails
green: disk
black: black hole
(here, neutron stars are
co-rotating – tidally locked)
48
r-process in NS-mergers
large neutron/seed ratios, fission cycling !
But: Ye free parameter …
49
Summary theoretical scenarios
NS-mergers
Supernovae
1e-5 - 1e-4
2.2e-2
Ejected r-process mass
(solar masses)
4e-3 – 4e-2
1e-6 – 1e-5
Summary
less frequent but
more ejection
more frequent and
less ejection
Frequency
(per yr and Galaxy)
50
What does galactic chemical evolution observations tell us ?
Argast et al. A&A 416 (2004) 997
Supernovae
NS mergers
Model star average
with error
observations
Average
ISM
Dots: model stars
 Neutron Star Mergers ruled out as major contributor
51
r- and s-process elements in stars with varying metallicity
(Burris et al. ApJ 544 (2000) 302)
s-process:
• later (lower mass stars)
• gradual onset (range of stars)
s-process
r-process
r-process:
• very early (massive stars)
• sudden onset (no low mass
star contribution)
~age
confirms massive stars
as r-process sites
(but includes SN and
NS-mergers)
52
Multiple “r-processes”
Star to star stability of all elements
(for very r-rich stars)
Star to star scatter of light vs heavy
for all stars [Fe/H]<-2.5, no s-process
(J.J. Cowan)
(Honda et al. 2004)
Additional “light” element primary process (LEPP) exists
(Travaglio et al. 2004 , Montes et al. 2006 to be published)
 It contributes to solar r-process residual abundances
53
Honda et al. 2006
Ivans et al. 2006
54
Honda et al. 2006
55
56
 Disentangling by isotope?
57
58