exotic beam studies in nuclear astrophysics marialuisa aliotta school of physics - university of...

Exotic beam studies in Nuclear Astrophysics

Marialuisa Aliotta

School of Physics - University of Edinburgh

11th Euro Summer School on Exotic Beams – Guilford, August 19-27 2004

introduction the need for RIBs in Astrophysics the tools of the trade

the synthesis of the trans-iron elements the s-process the r-process

explosive hydrogen burning novae and X-ray bursts the rp- and p-processes

experimental investigations techniques and detection systems selected examples

Lectures layout

why does one kilogram of gold cost so much more than one kilogram of iron?

Abundance curve of the elements

Fe

Au

7 orders of magnitude

less abundant!

“The 11 Greatest Unanswered Questions of Physics”based on National Academy of Science Report, 2002

[Committee for the Physics of the Universe (CPU)]

Question 3How were the elements from

iron to uranium made?

WHY?

Burbidge, Burbidge, Fowler & Hoyle (B2FH): Rev. Mod. Phys. 29 (1957) 547

from: M. Wiescher, JINA lectures on Nuclear Astrophysics

nucleosynthesis processes

The Synthesis of the Elements in Stars

charged-particle induced reaction

mainly neutron capture reaction

during quiescent stages of stellar evolution

involve mainly STABLE NUCLEI

mainly during explosivestages of stellar evolution

involve mainly UNSTABLE NUCLEI

nuclear processes

M.S. Smith and K.E. Rehm,Ann. Rev. Nucl. Part. Sci, 51 (2001) 91-130

Overview of main astrophysical processes

the vast majority of reactions encountered in these processes involve UNSTABLE species

hence the need for Radioactive Ion Beams

notations and definitions reaction mechanisms reaction cross sections stellar reaction rates

Lecture 1:

The tools of the trade

LOTS AND

LOTS OF

MATHS

WELCOME TOBASIC

ASTRONOMY.BEFORE WESTART, ARETHERE ANY QUESTIONS?

YEAH, LIKEWHAT MAKESASTRONOMYDIFFERENT

FROM ASTROLOGY?

http://www.univie.ac.at/strv-astronomie/unterhaltung.html

nuclear reaction rates

nuclear reactions in stars: a) produce energyb) synthesise elements

if Q > 0 net production of energy

reaction Q-value: Q=[(mp+mT)-(me+mR)]c2

= rQ/

(energy per single reaction)

reaction rate r

(number of reactions per unit time and volume)

Ni = number density of interacting speciesv = relative velocity(v) = velocity distribution in plasma(v) = reaction cross section

vNN1

1r Tp

pT

typical units: MeV g-1 s-1 energy production rate

vdv)v()v(v

stars = cooking pots of the Universe

for reaction: T(p,e)R (T=target, p=projectile, e=ejectile, R=recoil)

KEY QUANTITY

12C(p,)13N(e+)13C(p,)14N(p,)15O(e+)15N(p,)12C

CNO cycle

C

N

O

13

15

12

13 14 15

6 7 8

CNO isotopes act as catalysts

net result: 4p 4He + 2e+ + 2 + Qeff Qeff = 26.73 MeV

cycle limited by decay of 13N (t ~ 10 min) and 15O (t ~ 2 min)

example of nuclear reactions in stars

nucleosynthesis energy production

changes in stellar conditions changes in energy production and nucleosynthesis

(p,)

(p,)(e

need to know REACTION RATE at all temperatures to determine ENERGY PRODUCTION

abundance changes and lifetimes

consider reaction 1+2 3

where 1 is destroyed through capture of 2 and 3 is produced

reactions are random processes with constant probability (cross section) for given conditions

abundance change is governed by same laws of radioactive decay

define:

lifetime of 1 against destruction by with 2:

vNY

11

A1

vNYn

ndt

dn

A21

11

1 is destroyed 13 n

dt

dn

3 is produced

need to know REACTION RATE at all temperatures to determine NUCLEOSYNTHESIS

C

N

O

13

15

12

13 14 15

6 7 8

example

interacting nuclei in plasma are in thermal equilibrium at temperature T

also assume non-degenerate and non-relativistic plasma Maxwell-Boltzmann velocity distribution

kT2

vexpv

kT24(v)

22

2/3

with Tp

Tp

mm

mm

reduced mass

v = relative velocity

stellar reaction rate vdv)v()v(v

need: a) velocity distributionb) cross section

a) velocity distribution

example: Sun T ~ 15x106 K kT ~ 1 keV

kT ~ 8.6 x 10-8 T[K] keV

b) cross section

no nuclear theory available to determine reaction cross section a priori

depends sensitively on: the properties of the nuclei involved the reaction mechanism

and can vary by orders of magnitude, depending on the interaction

in practice, need experiments AND theory to determine stellar reaction rates

1 barn = 10-24 cm2 = 100 fm2

reaction cross sections

Reaction Force (barn) Eproj (MeV)

15N(p,)12C strong 0.5 2.0

3He(,)7Be electromagnetic 10-6 2.0

p(p,e+)d weak 10-20 2.0

examples:

nuclear properties relevant to reaction rates

recall: nucleons in nuclei arranged in quantised shells of given energy

nucleus’s configuration as a whole corresponds to discrete energy levels

exci

tatio

n en

ergy

00 0+

20Ne

1.63 2+

4.25

4.97

4+

2-

ground state

1st excited state

2nd excited state

3rd excited state

excitationenergy [MeV]

spinparity

J

any nucleus in an excited state will eventually decay either by , p, n or emission with a

characteristic lifetime which corresponds to a width in the excitation energy of the state

Heisenberg’s relationship

the lifetime for each individual exit channel is usually given in terms of partial widths

, p, n and with i

example

reaction mechanisms:

I. direct reactions

II. resonant reactions

reaction mechanisms

I. direct process

example: radiative capture A(x,)B

direct transition into a bound state

A+x

B

Q

Ecm

one-step process

2xAHB

reaction cross section proportional to single matrix element can occur at all projectile energies smooth energy dependence of cross section

H = electromagnetic operator describing the transition

other direct processes: stripping, pickup, charge exchange, Coulomb excitation

II. resonant process

two-step process

A+x

B

Sx

Ecm Er

1. Compound nucleus formation (in an unbound state)

B

2. Compound nucleus decay (to lower excited states)

reaction mechanisms

example: resonant radiative capture A(x,)B

2

Br

2

rf xAHEEHE

compound formationprobability x

compound decayprobability

reaction cross section proportional to two matrix elements

only occurs at energies Ecm ~ Er - Q

strong energy dependence of cross section

Q

N. B. energy in entrance channel (Q+Ecm) has to match excitation energy Er of resonant state,

however all excited states have a width there is always some cross section through tails

reaction mechanisms

example: resonant reaction A(x,)B

A+x

C

Sx

Ecm

1. Compound nucleus formation (in an unbound state)

CB

B+S

2. Compound nucleus decay (by particle emission)

2

xr

2

r xAHEEHB

compound formationprobability x

compound decayprobability

N. B. energy in entrance channel (Sx+Ecm) has to match excitation energy Er of resonant state,

however all excited states have a width there is always some cross section through tails

cross section

cross section expressions

for direct reactions

with neutrons with charged particles

example: direct capture

cross sections for direct reactions

)E(PAxHB22

x

A + x B +

“geometrical factor”de Broglie wavelengthof projectile

mE2

h

p

h

matrix elementcontains nuclear

properties of interaction

penetrability/transmission probability for projectile to reach target for interactiondepends on projectile’s angularmomentum and energy E

)E(S)E(PE

1

need expression for P(E)

factors affecting transmission probability:

Coulomb barrier (for charged particles only) centrifugal barrier (both for neutrons and charged particles)

= (strong energy dependence) x (weak energy dependence)

S(E) = astrophysical factor contains nuclear physics of reaction

+ can be easily: graphed, fitted, extrapolated (if needed)

tunneleffect

Ekin ~ kT (keV)

nuclear well

Coulomb potentialV

rR

for projectile and target charges Z1 and Z2

height ofCoulomb barrier

R

eZZ

4

1V

221

0C

in numerical units:

3/12

3/11

2121C AA

ZZ2.1

]fm[R

ZZ44.1]MeV[V

example: 12C(p,) VC= 3 MeV

average kinetic energies in stellar plasmas: kT ~ 1-100 keV!

fusion reactions between charged particles take place well below Coulomb barrier transmission probability governed by tunnel effect

Coulomb barrier

for E<<VC and zero angular momentum, tunnelling probability given by:

2exp)E(P with Gamow factor2/1

21 EZZ29.312

angular momentum is conserved in central potential

linear momentum p (and hence energy) must increase as distance d decreases

angular momentum barrier

incident particle can have orbital angular momentum

z-axis

incident particle

targetnucleus

p

dp = projectile linear momentumd = impact parameter

classical treatment:

L=pd

quantum-mechanical treatment:

)1(L = 0 s-wave= 1 p-wave = 2 d-wave…

with parity of wave-function: = (-1)

(discrete values only)

angular momentum is conserved in central potential non-zero angular momentum implies “angular momentum energy barrier” V

2

2

r2

)1(V

= reduced mass of projectile-target systemr= radial distance from centre of target nucleus

reactions with neutrons

neutron capture

simplest case: s-wave neutrons V = 0 and also VC = 0

attractive nuclear potential

incident wave transmitted wavereflected wave

-V0

discontinuity in potential gives rise to partial reflection of incident wave

transmission probability:

2/1EP for = 0 and hence:

2/1EP for 0 and hence:

v

1

E

1

2/1E

s-wave neutron capture usually dominates at low energies (except if hindered by selection rules)

V=0

higher neutron capture only plays role at higher energies

(or if =0 capture suppressed)

consequences:

neutron capture

cross section decreases strongly

with decreasing energy (effect of barrier)

dependence of neutron capturecross section

2/1E

2/1EP

note: arbitrary scale between different values

lower values dominate reaction rate

dependence of penetrabilitythrough centrifugal barrier

stellar reaction rates for neutron capture

v

1

vconstv

energy range of interest for astrophysics depends on:

temperature and cross section shape

s-wave neutron capture

energy range of interest E ~ kT

th = measured cross section for thermal neutrons

kT2vT

EdE)kT/Eexp()E(vdv)v()v(v

most probable velocity, corresponding to Ecm = kT

thTvv stellar reaction rate

example: 7Li(n,)8Li

deviation from 1/v trend due to resonant contribution (see later)

thermal cross section = 45.4 mb

s-wave neutron capture

case: = 0

v

1

E

1

reactions with charged particles

charged-particle capture

E

bexp)2exp(P

221 eZZ

E2

probability of tunnelling through Coulomb barrier for charged particle reactions

at energies E << Vcoul

assumes: full ion charges zero orbital angular momentum

)2exp(P

units of S(E): keV barn, MeV barn …

additional angular momentum barrier leads to a roughly constant addition to the S-factor

that strongly decreases with S-factor definition for charged particle reactions is

independent of orbital angular momentum (unlike neutron capture processes!)

Sommerfeld parameter

)E(S)2exp(E

1

stellar reaction rates for charged particle capture


dEE

b

kT

Eexp)E(Sv

maximum reaction rate at E0: 0E

b

kT

Eexp

dE

d

Gamow peak

tunnelling throughCoulomb barrier exp(- )

Maxwell-Boltzmanndistribution exp(-E/kT)

rela

tive

prob

abili

ty

energykT E0

E/EG

E0

E0 = relevant energy for astrophysics >> kT

3/29

3/122

21

2/3

0 TAZZ122.02

bkTE

MeV

6/59

6/122

210 TAZZ237.0kTE

3

4E MeV

Gamow peak

N.B. Gamow energy depends on reaction and temperature

and substituting for

cross section

cross section expressions

for resonant reactions

(neutrons and charged particles)

cross section for resonant reactions

for a single isolated resonance:

resonant cross section given by Breit-Wigner expression

22r

21

T1

2

2/EE1J21J2

1J2E

for reaction: 1 + T C F + 2

geometrical factor 1/E

spin factor J = spin of CN’s state

J1 = spin of projectile

JT = spin of target

strongly energy-dependent term1 = partial width for decay via emission of particle 1

= probability of compound formation via entrance channel

2 = partial width for decay via emission of particle 2

= probability of compound decay via exit channel

= total width of compound’s excited state

= 1 + 2 + + …

Er = resonance energy

what about penetrability considerations? look for energy dependence in partial widths!

partial widths are NOT constant but energy dependent!

particle widths 211 )E(P

R

2

P gives strong energy dependence

= “reduced width” (contains nuclear physics info)

energy dependence of partial widths

example: 16O(p,)17F

energy dependence of proton partial width p as function of

particle partial widths have approximatelysame energy dependence as penetrability function seen in direct reaction processes

integrate over appropriate energy region


here Breit-Wigner cross section

22r

21

T1

2

2/EE1J21J2

1J2E

E ~ kT for neutron induced reactionsE ~ Gamow window for charged particle reactions

if compound nucleus has an exited state (or its wing) in this energy range

RESONANT contribution to reaction rate (if allowed by selection rules)

typically: resonant contribution dominates reaction rate reaction rate critically depends on resonant state properties

reaction rate for resonant processes

two simplifying cases: narrow (isolated) resonances broad resonances

reaction rate for:

narrow resonances

broad resonances

<< ER

narrow resonance case

kT

Eexp

kT

2v R

R2

2/3

1212

resonance must be near relevant energy range E0 to contribute to stellar rate

MB distribution assumed constant over resonance region partial widths also constant, i.e. i(E) i(ER)

reaction rate for a single narrow resonance

NOTE

exponential dependence on energy means: rate strongly dominated by low-energy resonances (ER kT) if any

small uncertainties in ER (even a few keV) imply large uncertainties in reaction rate

resonance strength

21

T1 )1J2)(1J2(

1J2(= integrated cross section over resonant region)

often

21 1

21221

221

112

reaction rate is determined by the smaller width !

i values at resonant energies)

kT

Eexp

kT

2v R

R2

2/3

1212

some considerations…

rate entirely determined by “resonance strength” and energy of the resonance ER

experimental info needed:

partial widths i

spin J energy ER

note: for many unstable nuclei

most of these parameters are

UNKNOWN!

example: 24Mg(p,)25Al

the cross section

almost constant S-factordirect capture contribution

non-constant S-factorresonant contribution

… and the corresponding S-factor

Note varying widths of resonant states !

reaction rate through:

narrow resonances

broad resonances

22R

21

2

)2/()EE()E()E(

assume:

2=const, = const and use simplified

broad resonance case

~ ER

same energy dependenceas in direct process

for E << ER very weak energy dependence

broader than the relevant energy window for the given temperature

resonances outside the energy rangecan also contribute through their wings

Breit-Wigner formula +

energy dependence of partial and total widths

N.B. overlapping broad resonances of same J interference effects

summary

stellar reaction rate of nuclear reaction determined by the sum of contributions due to

direct transitions to the various bound states all narrow resonances in the relevant energy window broad resonances (tails) e.g. from higher lying resonances any interference term

Rolfs & RodneyCauldrons in the Cosmos, 1988

erferenceinttailsi

Rii

DCi vvvvv total rate

the Gamow window moves to higher

energies with increasing temperature

different resonances play a role at different temperatures

some considerations

130-220 keV330-670 keV

500-1100 keV

resonant and non resonant contributions

to stellar reaction rates

example

lower temperature rates by direct reactions higher temperature rates dominated by resonances

typically

note:

level density in nuclei increases

with excitation energy

Gamow region

Lecture 2:

The synthesis of the trans-iron elements

the s-process the r-process their astrophysical sites nuclear data needs


charged-particle induced reaction

mainly neutron capture reaction

during quiescent stages of stellar evolution

involve mainly STABLE NUCLEI

mainly during explosivestages of stellar evolution

involve mainly UNSTABLE NUCLEI

nuclear processes

why neutron capture processes for the synthesis of heavy elements?

exponential abundance decrease up to Fe exponential decrease in tunnelling probability for charged-particle reactions

almost constant abundances beyond Fe non-charged-particle reactions

binding energy curve fusion reactions beyond iron are endothermic

characteristic abundance peaks at magic neutron numbers

neutron capture cross sections for heavy elements increasingly larger

large neutron fluxes can be made available during certain stellar stages

Q > 0

Q < 0

mean lifetime of nucleus X against destruction by neutron capture

if n >> unstable nucleus decays

if n << unstable reacts

if n ~ branchings occur

nucleosynthesis beyond iron

start with Fe seeds for neutron capture

vN

1)X(

nn

whenever an unstable species is produced one of the following can happen:

the unstable nucleus decays (before reacting) the unstable nucleus reacts (before decaying) the two above processes have comparable probabilities

ALSO: can be affected too by physical conditions of stellar plasma!

NOTE: n varies depending upon stellar conditions (T, )

different processes dominate in different environments

with:

)X( mean lifetime of nucleus X against decay

7Be nucleus can only decay by electron capture with a lifetime: EC ~ 77 d

factors influencing the -decay lifetime of an unstable nucleus

both - and + decay are hampered in the presence of electron or positron degeneracy

- and + decays may occur from excited isomeric states maintained in equilibrium

with ground state by radiative transitions

electron-capture rates are affected by temperature and density through population of the K electronic shell

example: 7Be

in the Sun, T ~ 15x106 K kT ~ 1.3 keV low-Z nuclei almost completely ionized

e.g. binding energy of innermost K-shell electrons in 7Be: Eb = 0.22 keV

if no electrons available 7Be becomes essentially STABLE!

in fact free electrons present in the plasma can be captured

for solar conditions: EC ~ 120 d factor 1.6 larger than in terrestrial laboratory

aside

the s-process

FeCoNi

Rb

GaGe

ZnCu

SeBr

As

ZrY

Sr

Kr

(n,)

()

()

r-pro

cess

p-pro

cess

63Ni, t1/2=100 y

64Cu, t1/2=12 h, 40% (), 60% ()

79Se, t1/2=65 ky

80Br, t1/2=17 min, 92% (), 8% ()

85Kr, t1/2=11 y

r-only

p-only

s-only

neutron number

prot

on n

umbe

r(from Rene Reifarth)

s-only, r-only and p-only isotopes help to disentangle the individual contributions

A=140A=208

A=85

abundance peaks at A = 85, 140, and 208

how to explain abundance curve with the s process?

s-process abundances

the s-process

the process its astrophysical site(s) nuclear data needs (experimental equipment and techniques)

unstable nucleus decays before capturing another neutron

s-process (s = slow neutron capture process)

<< n

typical lifetimes for unstable nuclei close to the valley of stability: seconds years

assuming: ~ 0.1 b @ E = 30 keV v = 3x108 cm/s

1317 scm103v 316

nn cm

ns103

v

1N

requiring: n ~ 10 y Nn ~ 108 n/cm3

how many neutrons are needed?

classical approach of the s process

(t)(t)Nσv(t)(t)NNσv(t)(t)NNdt

(t)dNAAnA1An1A

A

Maxwellian averaged cross section

AA1A1AA NσNσ

dτ

dN

ttanconsNσNσ AA1-A1-A

production destruction

T

A

v

σvσ

t

0nT (t)dtNvτ neutron exposure

0dτ

dNA

in steady state condition (so-called local equilibrium approximation):

A

A σ

1N

time dependence of abundance NA given by:

N

assuming:

with:

AZ

T ~ constn >>

A=140

A=208

A=85

A

A σ

1N small capture cross sections at neutron magic numbers

pronounced abundance peaks

A=140A=208

A=85

Rolfs & Rodney: Cauldrons in the Cosmos, 1988

constantNσ AA

A~208

a superposition of many neutron irradiations is needed to correctly reproduce the abundance curve

main component (A=88-209) weak component (A<90)

A~85

A~140

condition fulfilled between

magic numbers of neutrons

sudden drops observed at

neutron magic numbers

NOTE

s-process: best understood nucleosynthesis process from nuclear point of view

what about the astrophysical site?

free neutrons are unstable they must be produced in situ

in principle many (,n) reactions can contribute

in practice, one needs suitable reaction rate & abundant nuclear species

most likely candidates as neutron source are:

s-process site(s) and conditions

13C(,n)16O 22Ne(,n)25Mg

22Ne

18O(,)

+)

14N

18F

(,)

25Mg

(,n)

13N

16O

13C

(,n)+)

12C

(p,)

astrophysical site:

core He burning (and shell C-burning) in massive stars (e.g. 25 solar masses)T8 ~ 2.2 – 3.5

astrophysical site:

He-flashes followed by H mixing into 12C enriched zoneslow-mass (1.5 - 3 Msun) TP-AGB starsT8 ~ 0.9 – 2.7

contribution to weak s-process contribution to main s-process

branching points can be used to determine

in some cases: ~ n a branching occurs in nucleosynthesis path

176Lugs 176Hf ~ 5x1010 y176Lum 176Hf ~ 5.3 h

for Nn ~ 108 cm-3 n ~ 1 y

176Lugs essentially STABLE176Lum quickly decays into 176Hf

significant amount of s-process branching from 176Lum -decay is required

need temperatures T8 > 1 to guarantee that isomeric state is significantly populated

from abundance determinations:

29Hf

Hf174

176

(note: 174Hf = p-only nucleus, i.e. not affected by s-process)

example:

in the star during the s process

about 15-20 branchings relevant to s process

neutron flux temperature density

F. Kaeppeler: Prog. Part. Nucl. Phys. 43 (1999) 419 – 483

stellar enhancement of decay (stellar decay rate/terrestrial rate)for some important branching-point nuclei in s-process path @ kT = 30 keV

experimental requirements

neutron capture cross sections on unstable isotopes 12 ≤ A ≤ 210 and 10 ≤ kT ≤ 100 keV about 1/10 of (n,) on radioactive isotopes measured so far

(n,) rates on long-lived branching point nuclei mass region: 79 ≤ A ≤ 204

prompt detection from (n,) reaction (TOF) activation technique (t½ < 0.5 y)

nuclear data needs

(n,) measurements

motivations improved stellar models for s-process temperature, density and neutron flux conditions galactic nucleosynthesis cosmo-chronometry

nuclear reactions: 7Li(p,n)7Be, 3H(p,n)3He (Karlsruhe, Tokyo)

(,n) reactions from energetic electrons on heavy metal targets (ORELA @ Oak Ridge, Tennessee and GELINA @ Geel, Belgium)

spallation reactions by energetic particle beams (LANSCE @ Los Alamos, n-TOF @ CERN)

neutron production techniques

capture cross-section measurements

Time-Of-Flight technique

applicable to all stable nuclei

need pulsed neutron source for En determination via TOF

signature for neutron capture events

total energy of cascade to ground state

need 4 detector of high efficiency,

good time and energy resolution

Karlsruhe: 42 individual BaF2 crystals

F. Kaeppeler: Rep. Prog. Phys. 52 (1989) 945 – 1013

advantages:

high sensitivity tiny samples are enough for (n,) mesurements good for RIBs

high selectivity samples of natural composition can be used

limitations: (n,) capture must produce unstable species cross section measurements at E= 25 (and 52) keV only

7Li(p,n)7Be (or 3H(p,n)3He)

angle-integrated spectrum closely resembles a MB distribution at kT = 25 keV (52 keV)reaction rate measured in such spectrum gives proper stellar cross section

applications with RIA type facilities: R. Reifarth et al.: NIMA 524 (2004) 215-226

activation technique

capture cross-section measurements

the s-process in a nutshell

temperature 2.2 – 3.5x108 K 0.9x108 Kneutron density 7x105 cm-3 4x108 cm-3

neutron source 22Ne(,n) 13C(,n) & 22Ne(,n)stellar site core helium burning TP-AGB stars

in massive stars

Weak component Main component

synthesis path along valley of -stability up to 209Bi n-source: 13C(,n)16O and/or 22Ne(,n)25Mg quiescent scenarios: e.g. He burning (T8 ~ 1 – 4; E0 ~ 30 keV)

branching points: if ~n several paths possible

data needs: (n,) cross sections on unstable nuclei along stability valleycapture data at branching points

motivation: s-process stellar models; physical conditions of astrophysical site

review: F. Kaeppeler: Prog. Part. Nucl. Phys. 43 (1999) 419 – 483

the r-process

the process its astrophysical site(s) nuclear data needs

r-process abundances Nr can be obtained as the difference between

solar abundances Nsolar and calculated s-process abundances Ns

ssolarr NNN

A=80

A=130A=195

constraints from elemental abundances

example:

Ultra Metal Poor giant halo stars give info on early nucleosynthesis in Galaxy

weak r-process main r-process

for A ≥ 130 solar-like abundances even for stars originating from very different regions of Galactic halo

main r-process independent of astrophysical site

for A ≤ 130 under-abundances

weak r-process

recall:[X/Y]=log(X/Y)-log(X/Y)solar

CS22892-052

red giant located in galactic halo

[Fe/H]= -3.0

[Dy/Fe]= +1.7

unstable nucleus reacts before capturing decay

r-process (r = rapid neutron capture process)

n<<

typical lifetimes for unstable nuclei far from the valley of stability: 10-6 – 10-2 s

requiring: n ~ 10-4 s Nn ~ 1020 n/cm3

termination point: fission of very heavy nuclei

explosive scenarios needed to account for such high neutron fluxes

waiting pointsRolfs & Rodney: Cauldrons in the Cosmos, 1988

classical approach of the r process

assume

(n,) ↔ (,n) equilibrium within isotopic chain, and

-flow equilibrium

-decay of nuclei from each Z-chain to (Z+1) is equal to the flow from (Z+1) to (Z+2)

waiting point approximation

the nucleus with maximum abundance in each isotopic chain must wait for the longer-decay time scales

(,n) photo-disintegration

equilibrium favors“waiting point”

-decay

seed

rapid neutroncapture

N

Z

good approximation for parameter studies, BUT steady-flow approximation is not always valid

)S,T,N(f)A,Z(Y

)1A,Z(Ynn

neutron capture rates do not play relevant role in most of r-process

abundance ratios of neighbouring isotopes only depends on Nn, T and Sn

note:

N

A

Z

A+1

abundance flow from neighbouring isotopic chains is governed by decays

A

A,ZYZYdefine: total abundance in each isotopic chain

N

AZ

Z+1

need nuclear masses (Sn) and lifetimes () together with environment conditions (Nn, T)

Z

1ZY

classical approach of the r process

late neutron captures can modify final abundance distribution mainly in region A > 140

courtesy: H. Schatz

timescale of r process summing up time spent at waiting points: t ~ 0.5 – 10 s

at present very little is known for neutron-rich nuclei very far away from stability must rely on theoretical calculations

@ ISOLDE -decays for ~ 30 neutron-rich nuclei have been determined including N=82 waiting points 130Cd & 129Ag

GSI ~ 70 new masses determined recently in region N=50 & 82

NuPECC Long Range Plan 2004

time-dependent r-process calculations

T = 1.35x109 K

Nn = 1020 gcm-3

t = 1.2 s

Nn = 1022 gcm-3

t = 1.7 s

Nn = 1024 gcm-3

t = 2.1 s a full fit to the solar r-process abundancesrequires a superposition of different stellarconditions (not necessarily different sites)

Pfeiffer et al.: Nucl. Phys. A 693 (2001) 282 – 324

time-dependent r-process calculations

Sn values from four different mass models; constant astrophysical parameters;

t½ for 129Ag and 130Cd calculated according to respective mass values

astrophysical site(s) for the r process

however, possible candidates for r-process site:

type II supernovae merging neutron stars neutrino driven wind of proto-neutron star He shell of exploding massive star merging neutron stars others?...

actual site(s) still unknown

type II supernova merging neutron stars

the r-process in a nutshell

temperature 1-2x109 Ktimescale ~ secondsneutron density 1020-1024 cm-3

neutron source unknownstellar site type II supernovae?

neutron star mergers?

data needs: neutron separation energies Sn (model dependent) nuclear masses far away from stability

-decay lifetimes for neutron rich nucleineutron capture cross sections on key isotopes

motivation: synthesis of heavy elements up to Th, U, Pur-process path(s)abundance patternconditions for waiting point approximation

review: Pfeiffer et al.: Nucl. Phys. A 693 (2001) 282 – 324

synthesis path far from valley of -stability synthesis of n-rich nuclei

waiting points: << n at closed shells abundance peaks (after n 0)

next-generation Radioactive Ion Beam facilities

will allow us to greatly improve our

knowledge of nuclear properties at extreme conditions

and will help our understanding of nuclear processes

in equally extreme astrophysical scenarios





Lectures layout

Lecture 3:

Explosive hydrogen burning

H-burning under extreme conditions astrophysical sites the rp- and p-process(es) nuclear data needs


inside the Sun

sunrise

sunset

hydrogen burning

COLD << part

HOT >> part

pp-chain, CNO, NeNa, MgAl cycles

Hot pp-chain super-massive low-metallicity starsnovae

HCNOHot NeNa (novae), X-ray burstsHot MgAl cycles

rp(p) process X-ray bursts & some SNI

mainly, quiescent stages of stellar evolution

mainly, explosive stages of stellar evolution

PP-I

Qeff= 26.20 MeV

proton-proton chain

p + p d + e+ + p + d 3He +

3He + 3He 4He + 2p

86% 14%

3He + 4He 7Be +

2 4He

7Be + e- 7Li + 7Li + p 2 4He

7Be + p 8B + 8B 8Be + e+ +

99.7% 0.3%

PP-II

Qeff= 25.66 MeV PP-III

Qeff= 19.17 MeV

net result: 4p 4He + 2e+ + 2 + Qeff

proton-proton chain

12C(p,)13N(e+)13C(p,)14N(p,)15O(e+)15N(p,)12C

C

N

O

13

15

12

13 14 15

6 7 8

CNO isotopes act as catalysts and accumulate at largest

net result: 4p 4He + 2e+ + 2 + Qeff Qeff = 26.73 MeV

cycle limited by decay of 13N (t ~ 10 min) and 15O (t ~ 2 min)

CNO cycle

cold CNO

12C(p,)13N(p,)14O(e+,)14N(p,)15O(e+)15N(p,)12C

C

N

O

13

15

12

13 14 15

6 7 8

hot CNO

14

cycle limited by decay of 14O (t ~ 70.6 s) and 15O (t ~ 2 min)

T8 ~ 0.8 – 1

T8 < 0.8

),p(N14v energy production rate:

very hot CNO-cycle

still “beta limited”

T8 ~ 3

3 4 5 6 7 8

9 10

11 12 13

14

C (6) N (7)

O (8) F (9)

Ne (10)Na (11)

M g (12)

T1/2=1.7s

3 flow

breakout

processing beyond CNO cycle

after breakout via:

T8 ≥ 3 15O(,)19Ne

18Ne(,p)21NaT8 ≥ 6

3 4 5 6 7 8

9 10

11 12 13

14

C (6) N (7)

O (8) F (9)

Ne (10)Na (11)

M g (12)

3 flow

explosive hydrogen burning

H burning at temperatures and densities far in excess of those attained in the interiorsof ordinary stars is expected to occur on a variety of astrophysical sites:

hydrogen burning under extreme conditions

hot burning in massive AGB stars (> 4 Msun) (T9 ~ 0.08)

nova explosions on accreting white dwarfs (T9 ~ 0.4)

X-ray bursts on accreting neutron stars (T9 ~ 2)

accretion disks around low mass black holes

neutrino driven wind in core collapse supernovae

explosive scenarios in binary systems

luminosity, time scales & periodicity properties depend on:

nature of accreting object accretion rate

novae

… the appearance of a “new” star…

sudden increase in luminosity by factor 104 – 106

mass ejection (10-3 Minitial)

energy release (1045 erg)

characteristic spectral lines

what causes these outbursts?

main features of novae explosions

how to account for observational features?

very common phenomenon

(about 10 y-1 in our Galaxy)

observational features

NOVAE semi-detached binary system:

White Dwarf and less evolved star (e.g. Red Giant)

assume:

RG completely fills its Roche Lobe while still undergoing expansion;

since no further expansion possible beyond Roche Lobe, matter transfer

occurs through Lagrange point

“dead” star (mainly CO or ONe)radiating away gravitational energymaintained by electron degeneracy

H-shell burningcore contraction (possibly He burning)envelope expansion

the model

n.b. about 50% of all stars are in binary systems

accretion process - artist’s impression

http://hubblesite.org/newscenter/newsdesk/archive/releases/2000/24/video/a

(p,) and (,p) reactions on proton-rich nuclei

nucleosynthesis up to A ~ 40 mass region

T ≤ 3x108 K

~ 103 g cm-3

H-rich matter transfer from RG to WD

temperature and density increase on WD’s surface

H-burning ignition at basis of accreted layer temperature increase due to energy release no expansion on degenerate matter no cooling further temperature/energy increase thermonuclear runaway cataclysmic explosion

(see later for details)nucleosynthesis path: Hot CNO, NeNa & MgAl cycles

matter ejected observations put important constraints on models

the model

12C

13N 14N

15O14O

15N

13C

16O 17O 18O

unstable

stable

HCNO

23Mg21Mg 22Mg

25Al24Al

24Mg 25Mg

26Al

26Mg

27Al

27Si 28Si

rp-processonset

17F 18F

18Ne 20Ne19Ne

21Na 22Na20Na

19F

21Ne 22Ne

23Na

breakoutfrom HCNO

NeNacycle

(,)(p,)

(+)(p,)

(,p)

reaction network for explosive hydrogen burning

(exact path depends on given stellar conditions)

overall mechanism well understood, however open questions still remain:

how and when does mixing at bottom of accreted layer occur?

missing-mass problem (models under-predict mass of ejecta)

do reactions produce more energy than expected?

what contribution to Galactic nucleosynthesis? (13C, 15N, 17O & 27Al)

does a break-out from HCNO cycle (e.g. 15O(,)19Ne) take place?

(models seem to indicate limited leakage out of HCNO)

how to explain observed overproduction (vs. solar) of elements in Ne to Ca region?

nuclear data needs

p-capture reactions on short-lived proton-rich nuclei e.g. 25Al(p,)26Mg and 30P(p,)31S

synthesis of e.g. 18F, 22Na, (26Al) very important for their characteristic -ray emissions

INTEGRAL satellite -ray detection constraint on models

reactions on P and S isotopes for possible flow out of MgAl cycle

considerations

X-ray bursts

main features of X-ray burstsregular recurrent bursts

irregular bursts

1036-1038 erg/s

(stars: 1033 - 1035 erg/s) duration 10 s – 100s recurrence: hours-days regular or irregular

frequent and very bright

phenomenon!

total ~230 X-ray binaries known


1038 - 1039 erg/s fast rise time (1-10 s) duration (~10-100 s) some show double peak at max spectral softening recurrence intervals (several hours)

type I X-ray bursts

type II X-ray bursts

rapid successions of bursts (few minutes interval) sudden flux drop without gradual decay from peak values no spectral softening in decay

figures from Lewin, W.H.G., van Paradijs, J. & Taam, R.E., 1993, Sp Sci Rev, 62, 223


(bursting pulsar:GRO J1744-28)

X-ray pulsarsregular pulses withperiods of 1- 1000 s

X-ray burstersfrequent outbursts of 10-100s durationwith lower, persistent X-ray flux in between

type I burstsburst energy proportionalto duration of preceedinginactivity period

most common type

type II burstsburst energy proportionalto duration of followinginactivity period

“rapid burster”and GRO J1744-28 ?

X-ray binaries

X-RAY BURSTERS & X-RAY PULSARS

semi-detached binary system: Neutron star + less evolved star

T ~ 109 K

~ 106 g cm-3

(,p) and (p,) reactions on proton-rich nuclei

nucleosynthesis up to A ~ 80-100 mass region

explosive ignition of H and/or He (via 3 12C) sudden increase surface temperature emission of strong X-rays cooling down of surface after explosion decay of burst profile if matter transfer continues process may repeat

nucleosynthesis path: breakout from Hot CNO, onset of rp-process

the model

0 1

2

3 4 5

6

7

8

9 10

11 12

13 14

15 16

17 18 19 20

21 22

23 24

25 26 27 28

29 30

31 32

33 34 35 36

37 38 39 40

41 42 43 44

45 46

47 48 49 50

51 52 53 54

n (0)

H (1)

H e (2)

L i (3)

Be (4)

B (5 )

C (6)

N (7)

O (8)

F (9 )

N e (10)

N a (11)

M g (12)

A l (13)

S i (14)

P (15)

S (16)

C l (17)

A r (18)

K (19)

C a (20)

Sc (21)

Ti (22)

V (23)

C r (24)

M n (25)

Fe (26)

C o (27)

N i (28)

C u (29)

Zn (30)

G a (31)

G e (32)

As (33)

Se (34)

B r (35)

K r (36)

R b (37)

S r (38)

Y (39)

Zr (40)

N b (41)

M o (42)

Tc (43)

R u (44)

R h (45)

Pd (46)

Ag (47)

C d (48)

In (49)

Sn (50)

burst ignition

prior to ignition hot CNO cycle

T9 ~ 0.2 ignition 3 12C

Hot CNO cycle II

T9 ~ 0.7 breakout 1: 15O(,)

T9 ~ 0.8 breakout 2: 18Ne(,p)

(~50 ms after breakout 1)leads to rp process and main energy production

courtesy: H. Schatz

0 1 23 4

5 6

7 8

9 10

11 12 13

14

15 16

17 18 19 20

21 22

23 24

25 26 27 28

29 30

31 32

33 34 35 36

37 38 39 4041

42 43 44

45 46 47 48

49 5051 52

5354 55

56

57 58

59

H (1)H e (2)L i (3)

Be (4) B (5) C (6) N (7)

O (8) F (9)

N e (10)N a (11)

M g (12)A l (13)S i (14) P (15)

S (16)C l (17)

A r (18) K (19)

C a (20)Sc (21)

Ti (22) V (23)

C r (24)M n (25)

Fe (26)C o (27)

N i (28)C u (29)

Zn (30)G a (31)

G e (32)As (33)

Se (34)B r (35)K r (36)R b (37)

S r (38) Y (39)

Zr (40)N b (41)

M o (42)Tc (43)

R u (44)R h (45)Pd (46)Ag (47)

C d (48)In (49)

Sn (50)Sb (51)

Te (52) I (53)

Xe (54)

p process: 14O(,p)17F17F(p,)18Ne18Ne(,p)21Na21Na(p,)……

3 reaction++ 12C

thep process

alternating (,p) and (p,) reactions:

for each proton capture there is an

(,p) reaction releasing a proton

net effect: pure He burning

courtesy: H. Schatz

rapid proton captureon proton-rich nuclei

rp process: 41Sc(p,)42Ti42Ti(p,)43V43V(p,)44Cr44Cr(+)44V44V(p,)…

0 1 23 4

5 6

7 8

9 10

11 12 13

14

15 16

17 18 19 20

21 22

23 24

25 26 27 28

29 30

31 32

33 34 35 36

37 38 39 4041

42 43 44

45 46 47 48

49 5051 52

5354 55

56

57 58

59

H (1)H e (2)L i (3)

Be (4) B (5) C (6) N (7)

O (8) F (9)

N e (10)N a (11)

M g (12)A l (13)S i (14) P (15)

S (16)C l (17)

A r (18) K (19)

C a (20)Sc (21)

Ti (22) V (23)

C r (24)M n (25)

Fe (26)C o (27)

N i (28)C u (29)

Zn (30)G a (31)

G e (32)As (33)

Se (34)B r (35)K r (36)R b (37)

S r (38) Y (39)

Zr (40)N b (41)

M o (42)Tc (43)

R u (44)R h (45)Pd (46)Ag (47)

C d (48)In (49)

Sn (50)Sb (51)

Te (52) I (53)

Xe (54)

Schatz et al.: Phys Rev Lett 86 (2001) 3471(4)

3 reaction++ 12C

p process: 14O(,p)17F17F(p,)18Ne18Ne(,p)21Na21Na(p,)……

rp-process end point:the Sn-Sb-Te cycleknown ground state emitters

therp process

endpoint varies with conditions: amount of H available at ignition peak temperature

type I X-ray burst

luminosity

abundances of waiting points

H, He abundance

nuclear energy generation

H. Schatz et al.: Phys Rev Lett 86 (2001) 3471(4)

considerations

overall mechanism seems well understood, however open questions still remain:

why are there bursters and pulsars?

why do burst timescales vary from ~10 s to ~100 s?

what is the origin of superbursts (103 times stronger and longer, t ~ 104-105 s)?

what crust composition of neutron star?

what contribution of X-ray bursts to galactic nucleosynthesis?

some Type I X-ray bursts show double peak in luminosity separated by a few seconds

possible cause: waiting points in thermonuclear reaction flow ?

waiting points: 26Mg, 26Si, 30S, and 34Ar ?

(,p) reactions too weak because of Coulomb barrier and

(p,)-reaction quenched by photo-disintegration

conclusion: effects of the experimentally unknown 30Sp)33Cl and 34Arp)37K

directly visible in the observation of X-ray burst light curves?

observational signatures for X-ray outbursts?

Fisker and Thielemann: arXiv:astro-ph/0312361 v1 13 Dec 2003

nuclear data needs

thermonuclear runway driven by p-process and rp-process

breakout reactions critical constraints on ignition conditions + runaway timescale

proton-capture reactions on short-lived proton-rich nuclei up to proton-drip line

influence on temperature and luminosity

-decay studies of isomeric and/or excited states

masses, lifetimes, level structure, proton-separation energies

0 1 23 4

5 6

7 8

9 10

11 12 13

14

15 16

17 18 19 20

21 22

23 24

25 26 27 28

29 30

31 32

33 34 35 36

37 38 39 4041

42 43 44

45 46 47 48

49 5051 52

5354 55

56

57 58

59

H (1)H e (2)L i (3)

Be (4) B (5) C (6) N (7)

O (8) F (9)

N e (10)N a (11)

M g (12)A l (13)S i (14) P (15)

S (16)C l (17)

A r (18) K (19)

C a (20)Sc (21)

Ti (22) V (23)

C r (24)M n (25)

Fe (26)C o (27)

N i (28)C u (29)

Zn (30)G a (31)

G e (32)As (33)

Se (34)B r (35)K r (36)R b (37)

S r (38) Y (39)

Zr (40)N b (41)

M o (42)Tc (43)

R u (44)R h (45)Pd (46)Ag (47)

C d (48)In (49)

Sn (50)Sb (51)

Te (52) I (53)

Xe (54)

masses (proton separation energies)

-decay rates

direct reaction rate measurementswith radioactive beams have begun(ANL, LLN, ORNL, ISAC)

(p,) and (,p) reaction rates

indirect information about rates

from radioactive and stable beam experiments

(transfer reactions, Coulomb breakup, …)

many lifetime measurements at radioactive beam facilities (LBL, GANIL, GSI, ISOLDE, MSU, ORNL)know all -decay rates (earth)location of drip line known (odd Z)

separation energiesexperimentally known up to here

some recent mass measurenents-endpoint at ISOLDE and ANLIon trap (ISOLTRAP)

from: H. Schatz

nuclear data need & present status





Lectures layout

Lecture 4:

Experimental techniques and equipment

studies with RIBs reaction types direct approach indirect approach

some selected examples


Isotope Separation on Line (ISOL) (CERN, LLN, ORNL, TRIUMF) Projectile Fragmentation (PF) (GANIL, GSI, MSU, RIKEN) in-flight production (ANL, Notre Dame, TAMU) batch mode production (suitable for long-lived species) …

RIB production methods

independent from chemical properties no limitations on t½ (fast separation)

typical beam energies too high for NA poorer beam quality (energy, size) possible beam contaminations

excellent quality high purity high intensities

limited number of species different production for different species limited to nuclei with t½ ≥ 1s (allow for diffusion)

M.S. Smith and K.E. Rehm, Ann. Rev. Nucl. Part. Sci, 51 (2001) 91-130

targets

solid CH2 target (plastic material)

simple to handle hydrogen depletiondx ~ 50 - 1000 g cm-2 non uniformity

melting problemsdeuterium contamination

He targets

H targets

solid implanted target

simple to handle low concentration (n ~ 1015 - 1017 atoms cm-2)

window-confined gas target

windowless gas target

higher concentration background reactions(depending on pressure) (e.g. on window materials)

higher concentration differential-pumping systemalmost background free high pumping speedsno physical degradation

NUCLEAR DATA NEEDS

reactions involving:

A < 30

A > 30

cross-section dependence:

individual resonancesnuclear properties

statistical properties Hauser-Feshbach calculations

excitation energies

spin-parity & widths

decay modes

masses

level densities

part. separation energy

knowledgerequired:

very large amount of reactions in various astrophysical sites not feasible to study all reactions involved experimental constraints wherever possible

reaction types and experimental techniques

experimental approaches

DIRECT APPROACHES• radiative capture reactions• transfer reactions

INDIRECT APPROACHES• resonant elastic scattering• transfer reactions• time-inverse reactions• coulomb dissociation• ...

DIRECT APPROACHES

RADIATIVE CAPTURE X(p,)Y

X(,)Y

among the most common reactions

in nuclear astrophysics

heavy recoil detection

inverse kinematics forward peaked emission ( ~ 1o) detection efficiency ~ 100%

BUT: high suppression factors required (1010-1015)

RIB intensities ~ 107 ions/s

Smith & Rehm (2001)

examples: 21Na(p,)22Mg, 15O(,)19Ne @ DRAGON – TRIUMF

delayed decay measurements

-ray detection low efficiency ~ 4 coverage needed

-ray background induced by + beam decay

example: 22Na(p,)23Mg @ ANL

example: 19Ne(p,)20Na(+)20Ne*()16O @ Louvain-la-Neuve

100 HPGe detector array

absolute efficiency: 9% for 1.33 MeV ray

gammaspherecourtesy: D. Jenkins

TRANSFER REACTIONS mainly X(p,)Y and X(,p)Y


silicon strip detector arrays large solid angle coverage (e.g. LEDA – see later)

light-heavy nuclei coincidence

examples: 18Ne(,p)21Na @ LLN Groombridge et al.: Phys Rev C 66 (2002) 55802(10)

& TRIUMF (planned) TUDA collaboration

18F(p,)15O @ ORNL Bardayan et al.: Phys Rev C 63 (2001) 65802(6)

Bardayan et al.: Phys Rev Lett 89 (2002) 262501(4)

14O(,p)17F @ LLN (planned) Aliotta et al.

INDIRECT APPROACHES

RESONANT ELASTIC SCATTERING typically X(p,p)X

inverse kinematics + suitably thick target (eg (CH2)n) + proton detection

protons undergo little energy loss, little kinematics variation, little straggling

investigate resonance properties of compound nucleus

the method

retain information on resonance shape

proton spectrum

high-energyprotons

low-energyprotons

low beam energy loss

high beam energy loss

excitation function

energy

target thickness E

Coszach et al.: Phys Rev C 50 (1994) 1695

19Ne(p,p)19Ne resonant elastic scattering

a real case

19Ne + p

20Ne

target thickness E

excitation function

requirement: proton width p ≥ 1 keV due to current limits in detection energy resolution


examples: 11C(p,p) & 7Be(p,p) @ LLN

21Na(p,p) @ TUDA - TRIUMF

useful approach when:

resonance energy not known partial widths not known very weak resonance(s) dominate(s) reaction rate

fit to experimental data (e.g. by R-matrix analysis) to determine resonance properties

relevance

limitations

typically enough thanks to high elastic scattering cross sections

advantage

resonant elastic scattering

progressively harder at lower energies due to increase in Rutherford cross sectionand decrease in resonance widths because of penetrability effects

TRANSFER REACTIONS

e.g. X(d,p)Y

investigate (n,) reactions for s- (or r-)processRIB intensities ~ 104 ions/s

TIME-INVERSE REACTIONS

MASS & LIFETIME MEASUREMENTS

e.g. Coulomb dissociation

investigate (x,) radiative capture reactions

… and much more…

SELECTED EXAMPLES

the 21Na(p,p) resonant elastic scattering

the 21Na(p,)22Mg reaction


the 21Na(p,p)21Na reaction

in A = 21 mass region and for nova (T9 ≤ 0.4) [and X-ray burst temperatures (T9 ≤ 2)]

nuclear level densities are low

proton capture reactions dominated by capture into narrow, isolated resonances

21

T1 )1J2)(1J2(

1J2

kT

Eexp

kT

2v R

R2

2/3

1212

recall from Lecture 1:

need: knowledge of resonance energy and properties of state

9

R2/39

11A T

E605.11expT10x54.1vN

with

stellar reaction rate

considerations

(cm3 s-1mol-1)

reaction rate resonant strength

22r

21

T1

2

2/EE1J21J2

1J2E

Breit-Wigner cross section

fold in with Maxwell-Boltzmann distribution and integrate to get:

the 21Na(p,p) reaction

first measurement with RIB @ TRIUMF – Canada

TUDA (TRIUMF – UK Detector Array)

multipurpose scattering chamber

LEDA system

Louvain-Edinburgh Detector Array CD – Compact Detector

large area larger solid angle coverage high segmentation reduced sensitivity to background (i.e. beam -decay) good modularity different configurations possible

8 sectors x 16 strips = 128 individual channels

T. Davinson et al.: NIM A 454 (2000) 350-358 (for details on LEDA detector array)

12C

13N 14N

15O

17F 18F

18Ne 20Ne19Ne

21Na 22Na20Na

15N

23Mg21Mg 22Mg

14O

Hot CNO cycle

rp-process onset

breakout from HCNO

investigate states in 22Mg of astrophysical relevance for 21Na(p,)22Mg reaction

astrophysical motivation

novae explosions synthesis of 22Na onset of rp-process?

21Na = 5x107 pps

Ecm = 0.45 – 1.40 MeV

(CH2)n = 50 g cm-2 & 250 g cm-2

the experiment

6.813C. Ruiz et al: Phys. Rev C65 (2002) 42801 (R)

determined resonance energies, widths and spin-parity assignments

the results

courtesy: A. Murphy

20Na(p,p)20Na

investigate states in 21Mg of astrophysical relevance for

20Na(p,)21Mg reaction

evidence for a new resonance at E = 4.44 MeV

resonant elastic scattering with TUDA @ TRIUMF

more recently…

data analysis still in progress…

PRELIMINARY!

novae explosion rate: 30-40 y-1

Galaxy lifetime ~ 1010 y

mass ejected ~ 2x10-5 Msun / nova

novae scarcely contribute to Galactic abundances BUT produce key isotopes

(e.g. 7Be, 22Na and 26Al) which provide clear signature of novae explosion via

characteristic -ray emission José & Hernanz, ApJ 494 (1998) 680

astrophysical motivation

22Na: the fingerprint of a nova outburst Clayton & Hoyle, ApJ L101 (1974) 187

novae are thought to be principal galactic site for 22Na

22Na decay has conveniently short half-life

hence spatial and temporal limits to its detection

t½ = 2.602 y

22Na

1.275 MeV

+

22Ne

idea: observe -ray signature to test nova models

why 22Na?

from COMPTEL data on 5 ONe-type novae

(Nova Her 1991, Nova Sgr 1991, Nova Sct 1991, Nova Pup 1991, and Nova Cyg 1992)

+ data from CO-type novae

only upper limit 3x10−8 Msun of 22Na ejected by any nova in the Galactic disk

A. F. Iyudin et al., A&A 300 (1995) 422

over-production predicted by models? (e.g. 10-4 Msun for ONe novae)

too low -ray detection sensitivity? unclear (nuclear) inputs?

theoretical models of novae indicate that measurable -ray fluxes should be observed

from novae explosions within few kilo-parsecs from the Sun (1pc = 3.26 ly)

WHY?

reduce uncertainties affecting reaction rates for production and destruction of 22Na

important constraints on models and observations (i.e. distance and epoch)

compare with observations:-ray detection by INTEGRAL mission

20Ne(p21Na(+)21Ne(p,)22Na (cold NeNa)

20Ne(p21Na(p,)22Mg(+)22Na (hot NeNa)

José: arXiv:astro-ph/0407558 v1 27 Jul 2004in Ne-enriched envelopes of ONe novae:

destruction of 22Na

22Na(p,)23Mg (mainly)

production of 22Na

first direct measurement of the 21Na(p22Mg rate

recently performed with DRAGON recoil separator @ TRIUMF

also indirect determination of the 22Na(p23Mg rate

carried out with the Gamma-sphere array @ ANL

22Na production and destruction processes


(the 22Na production process)

the experiment

first direct measurements of 21Na(p,)22Mg @ TRIUMF – Canada

S. Bishop et al.: Phys Rev Lett 90 (2003) 162501(4)

J.M. D’Auria et al.: Phys Rev C 69 (2004) 65803(16)

Detector of Recoils And Gammas Of Nuclear reactions

differentially pumpedwindow-less

gas target (~ 12cm)

30 BGO-ray detector array

~ 90% of 4 two-stagerecoil separator

(magnetic and electric dipoles)

DSSSDdouble sided silicon strip detectors

for recoil detection

see: J.M. D’Auria et al., Nucl Phys A 701 (2002) 625 and D. Hutcheon et al., NIM A 498 (2003) 190 for details

the DRAGON facility

(start)

(stop)

aim of the measurements:

direct determination of the resonant strengths for key resonances in 22Mg

the method:

thick-target yield measurement

effects of target thickness in cross section measurements

consider: resonance with width and target with thickness

if << thin-target case

yield measurement proportional to resonance profile

if >> thick-target case

yield measurement integrated Breit-Wigner cross sectionsmooth step function

arctan1

m

mmY

T

Tp2

max

yield measurement gives directly

max at E= ER

FWHM =

ER

Y/Ymax

1

1/2

1/4

ER

3/4

negligible beam energy spread

Y/Ymax

1

m

mm

2Y

T

Tp2

max

9Be(p,)10B ER = 1.017 MeV = 4 keV

proton energy [keV]

cou

nts

[a

.u.]

rule for thick-target yield applicability: ~ 6

example

see: Rolfs & Rodney, Cauldrons in the Cosmos (1988) for details

measured for state at Ex = 5.714 MeV

= 1.03 ± 0.16stat ± 0.14sys meV


the results

22Mg


Ecm = 202 -221 keV

21Na beam ~ 109 s-1 (t½ = 22.5 s)

P (H2) = 4.6 Torr

uncertaintiesfrom previous

estimates

new rate


reduced uncertainties on previous estimates

higher rate lower 22Na abundance lower detection limits

confirmed maximum -ray detection distance ~ 1 kpc (with INTEGRAL)

astrophysical implications


9

R2/39

11A T

E605.11expT10x54.1vN

the results

measured for seven resonances of astrophysical relevance in Ecm = 200-1103 keV



Ecm = 0.2 – 1.1 MeV

21Na beam ≤ 109 s-1 (t½ = 22.5 s)

P (H2) ~ 5 Torr

note: little agreement on spin-parity values of excited states, except for Ex = 5.714 MeV

also, state at Ex = 5.837 MeV observed once but not confirmed by other experiments



novae X-ray bursts

dominant contribution from state at Ex = 5.714 MeV up to T9 ~ 1.1

novae: same conclusions as in Bishop et al. (2003)

X-ray bursts: two nova models, using present rate and rate = 0 for 21Na(p,) evolution of explosion is the same alternative routes to heavier nuclei are possible 21Na(p,) not critical



(the 22Na destruction process)

D.G. Jenkins et al., Phys Rev Lett 92 (2004) 031101(4)

12C(12C,n)23Mg

heavy-fusion reaction to investigate properties of relevant states in 23Mg

newly discovered level at Ex = 7.769 MeV in 23Mg (189 keV above proton threshold)

may have a non-negligible contribution to stellar rate if spin-parity confirmed

need for new investigation

new states observed accuracy in excitation energies improved several spin-parity uncertainties resolved

12C 10 pnA Elab 22 MeV 12C target: 40 g cm-2

detector Gammasphere

measurement at Argonne National Laboratory

limits detection to within 0.6 kpc from the Sun instead of 1 kpc from previous rate


C. Angulo et al., Nucl Phys A 656 (1999) 3

~ factor 3 less 22Na ejected in novae

previous limits from NACRE new limits

D.G. Jenkins et al., Phys Rev Lett 92 (2004) 031101(4)


22Na(p,)23Mg

summary and outlook

M.S. Smith and K.E. RehmAnn. Rev. Nucl. Part. Sci, 51 (2001) 91-130

overview of main astrophysical processes

nuclear reactions involving UNSTABLE nuclei are take place in a many astrophysical sites

Radioactive Ion Beams (will) play dominant role in understanding these processes

summary and outlook

at present limitations on:

species which can be produced at current facilities

intensities of available RIBs

much has already been achieved since first RIBs available at LLN (~1990)

better knowledge of explosive H-burning processes

better understanding of key reactions constraints on physical conditions

in astrophysical sites

first measurements ever on structure and properties of highly exotic nuclei

(proton-rich and neutron-rich)

future projects (e.g. RIA, EURISOL, upgraded GSI) aimed at overcoming present limitations

will open up new and exciting avenues in experimental Nuclear Astrophysics

for example…

however…

planned future facilities

The EURISOL Report: published by GANIL, 2003

RIA – Physics White Paper

RIA

Rare Isotope Accelerator facility

objectives

nature of nucleonic matter origin of the elements and energy generation in stars tests of the Standard Model and of fundamental conservation laws

the RIA facility

ISOL fragmentation

yield estimates available at RIA

enormous discovery potential !

RIA – Physics White Paper

the EURISOL project

The EURISOL Report: published by GANIL, 2003

1 GeV 5 mA proton beam RFQLINAC

With special thanks to:

Hendrik Schatz

Michael Heil

for permission to use and readapt some of their materialand to:

Rene Reifarth

for unknowingly providing me with one of his slides

W.D. Arnett and J.W. Truran NucleosynthesisThe University of Chicago Press, 1968

J. Audouze and S. Vauclair An introduction to Nuclear AstrophysicsD. Reidel Publishing Company, Dordrecth, 1980

E. Böhm-Vitense Introduction to Stellar Astrophysics, vol. 3Cambridge University Press, 1992

D.D. Clayton Principles of stellar evolution and nucleosynthesisThe University of Chicago Press, 1983

H. Reeves Stellar evolution and NucleosynthesisGordon and Breach Sci. Publ., New York, 1968

C.E. Rolfs and W.S. Rodney Cauldrons in the CosmosThe University of Chicago Press, 1988 (…the “Bible”)

BOOKS

copies of these lectures at: www.ph.ed.ac.uk/~maliotta/teaching

further reading

exotic beam studies in nuclear astrophysics marialuisa aliotta school of physics - university of...

Documents

energy production slide

neutron capture reaction

energy production rate

example of nuclear reactions

nucleosynthesis p

rq energy

random processes

recoil key quantity