Asignatura de Física Nuclear Universidad de Santiago de ... · Física Nuclear, Tema 15. 2. José...
Transcript of Asignatura de Física Nuclear Universidad de Santiago de ... · Física Nuclear, Tema 15. 2. José...
1
Universidad de Santiago de Compostela
Asignatura de Física NuclearCurso académico 2011/2012
Tema 15
Nucleosíntesis primordial y estelar
Física Nuclear, Tema 15 José Benlliure
2Física Nuclear, Tema 15 José Benlliure
Decays and reactions in the stellar plasma
1.1 Gamma-ray transitionsIn a hot plasma, excited states in a given nucleus are thermally populated through photon absorption, Coulombexcitations by surrounding ions, inelastic particle scattering or other mechanisms. The time scale for excitationand de-excitation is much shorter than stellar hydrodynamics time scales. Contrary to lab investigations wheredecaying or reacting nuclei are in their ground state, these excited states will play an important role in stellardecays or reactions. At thermal equilibrium the probability for populating a given state
can be obtained as:
kTE
kTE
eg
egNN
P /
/
The decay of 26Al represents a clear example of the role of excited states in the nuclear media. This nucleus is though to be produced In type II supernovae during the explosive carbon and neon burning phases. This nucleus decays tothe first excited state (1809 keV) in 26Mg. The observation of this gamma ray inseveral -ray telescopes as COMPTEL aboard the Gamma Ray Observatory isa major prove for nucleosynthesis processes in the Universe.
Since the first excited state in 26Al is an isomer decaying to the ground state in26Mg, the observed intensities of 1809 keV -rays can only be transformed in nucleosynthesis rate of 26Al if one takes into consideration the populations of the different excited states in 26Al, and in particular the isomeric state.
g=2j+1 being j the spin of the state
3 José Benlliure
1.2 Weak interactions-decay:In a hot plasma, excited states in a parent nucleus are thermallypopulated and may also undergo -decay transitions to thedaughter nucleus. Even stable nuclei may undergo a -decay in the stellar medium. The total -decay rate will be given by the weighted sum of the individual transition rates ij according to:
i j
ijiP * The sum over i and j runs over the parent and daughter
states and Pi can be obtained from the previous equation.
Under these conditions, -decay becomes temperature dependent by also density dependent at sufficiently large values of density when the electron gas is degenerate, limiting the number of final states available forthe electron emission.
Temperature dependenceof the decay rate of 26Al
Decays and reactions in the stellar plasma
Física Nuclear, Tema 15
4 José Benlliure
1.2 Weak interactionsElectron capture:At the temperature typical of the stellar interior most nuclei posses few, if any, bound electrons. Being theirdecay constant for bound electron capture very small. However, and due to the density of free electrons inthe stellar medium, nuclei can decay through the capture of free electrons. The probability for this process is proportional to the electron density and inversely proportional to the average electron velocity.
At low densities, the kinetic energies of the free electrons are usually small. At very high densities, however, the (Fermi) energy of the degenerate electrons may become sufficiently large to cause nuclei to undergocontinuum capture of energetic electrons, even if they are stable in the laboratory.
Pair production:At high temperatures pair production can become and effective process. Then positron capture should be considered in addition to the continuum electron capture.
Neutrino energy loss mechanism:This mechanism known as Urca process becomes important at high temperatures and densities and consistsof alternate electron captures and --decays involving the same pair of parent and daughter nuclei being the netresult of two subsequent decays of a neutrino anti-neutrino pair.
eXeXXXeX NAZN
AZN
AZN
AZN
AZ ),(),( -
11 L
Decays and reactions in the stellar plasma
Física Nuclear, Tema 15
5 José Benlliure
1.3 Particle induced reactionsIn the stellar medium reactions are produced by collisions of two moving particles therefore the reaction rate (reactions per time and volume unit) will be determined by the number of colliding particles per volume unit, their relative velocity and their interaction cross section.
)(1001 vvNNr
In a stellar plasma at thermodynamic equilibrium the velocity of the constituent particles follows a given distribution. Then we can generalize the expression for the reaction rate taking into account the possibilityof having identical colliding particles as:
0 01
011001101001 )1(
)()(
vNN
vNNdvvvvPNNr
Since nuclei in the plasma move non relativistically their relative velocities can be described by a Maxwell-Boltzmann distribution, then:
Em
mdE
mEe
kTmdEEPdvve
kTmdvvP kTEkTvm
224
2)( 4
2)( 01
0101
/2/3
012)2/(2/3
01 201
)( )(107318.3)()(18 113
0 0
/605.11
10
102/3
9
10/
2/3
2/1
0101
9
smolcmdEeEE
MMMM
TdEeEE
kTmvN TkTE
A
The reaction rate per mol is defined as::
With E given in MeV, T9 =T/109 in kelvin, , M in units of u and the cross section in barn.
Decays and reactions in the stellar plasma
Física Nuclear, Tema 15
6 José Benlliure
1.4 Photon induced reactionsIf one of the colliding particles is a photon the reaction is called photodisintegration (+30+1). Considering that photons move with the speed of light we can write the reaction rate and the corresponding decay constant as:
0
3
3
033 )()()3( )()(
dEEEcNNr
dEEEcNNr
The energy distribution of photons can be obtained from the Planck’s radiation law u(E).
dEe
Ehc
dEEEu
dEEN kTE 1)(8)(
)( /
2
3
Then, the final expression for the decay rate will be:
0 /
2
23 )(1
8)3(
dEEe
Ech kTE
In general photodisintegration are endothermic reactions, then thelower integration limit in the previous equation will be Q3 .
Decays and reactions in the stellar plasma
Física Nuclear, Tema 15
(decay probability per nucleus per second)(reactions per second)
7 José Benlliure
1.5 Abundance evolutionThe rate of change of the abundance of nucleus 0 due to reactions with nucleus 1 can be expressed as:
011001011
001
1
0 )1()0(
)0( vNNrNNdtdN
From these equations we obtain the following relations:
011
1011
11
1
011
1
0110101
01
1
0
0101
0101
)0(1)0(
1)1(
)0(
)0()1(1
)1()0(
vNMXvN
vNMX
vNrN
NNr
A
A
The decay constant of a nucleus for destruction via particle-induced reactions depends explicitly on the stellardensity and implicitly on temperature. If a nucleus 0 can be destroyed by different reactions its total lifetime is:
i i )0(
1)0(
1
Decays and reactions in the stellar plasma
Física Nuclear, Tema 15
Xi =mass fraction
8Física Nuclear, Tema 15 José Benlliure
Cross sections and reaction rates
2.3 Abundance evolutionThe final abundance of a nucleus can be obtained takinginto account all creation and destruction mechanisms (reactions, -decays, photodisintegrations,…)
n o pipiioiniin
kj l mmimlilijkkj
i
NNvNN
NNvNNdtdN
,,
,,,
-
In a realistic situation we should consider the evolution of not just one nuclide, but of several species simultaneously. Such a system of coupled, nonlinear ordinary differential equations is called a nuclearreaction network. Very often, equilibrium conditions together with the reciprocity theorem helps in solvingthese nuclear reaction networks.
9Física Nuclear, Tema 15 José Benlliure
Cross sections and reaction rates
2.4 Forward and reverse reactionsThe cross sections of a forward and reverse reaction are related by the reciprocity theorem. If we considerreactions involving particles with rest mass 0+12+3 we obtain the following relation:
)1()1(
)12)(12()12)(12(
01
23
2323
0101
32
10
2301
0123
EmEm
jjjj
kTQ
kTE
kTE
A
A emm
jjjj
dEeE
dEeE
mm
vNvN /
3/2
23
01
0132
2310
001
/230101
023
/012323
2/1
23
01
2301
0123 2301
01
23
)1)(12)(12()1)(12)(12(
For reactions involving photons, 0+1+3, we obtain these other expressions:
)1(1
)12(2)12)(12(
012
012
01
3
10
301
013
EEcm
jjj
001
/301012/3
2/1
01
03012
012
01
013
10/
2
23
301 01
)(8
)1)(12()12)(12(
18
)3(
dEeEkTN
m
dEEEcm
jjj
e
Ech
vNkTEA
kTE
A
10Física Nuclear, Tema 15 José Benlliure
1.6 Reaction rate equilibriaThe net reaction rate considering forward and reverse reactions is obtained as:
)1()1( 23
012332
01
23011001232301
vNNvNNrrr
being the equilibrium condition (r=0):
kTQnormnorm
normnorme
mm
GGGG
jjjj
vv
NNNN /
3/2
01
23
10
32
10
32
012301
230123
10
32 2301
)12)(12()12)(12(
)1()1(
while for reactions involving photons:
301
30110013301 )3(
)1(N
vNNrrr
kTQnormnorm
norme
GGG
jjj
kTmhv
NNN /
10
3
10
32/3
01
3/22301
0110
3 301
)12)(12()12(
)(1
2)3()1(1
This last expression is known as the Saha statistical equation.
Decays and reactions in the stellar plasma
i
kTEinorm
i g
egG
i
/
11Física Nuclear, Tema 15 José Benlliure
1.7 Nuclear energy generationThe nuclear energy generation in stars is given by the Q value of the thermonuclear reactions taking place in thestellar medium. The energy production per unit time and unit mass is given by::
1
0
01
2301
01
2301102301230123012301 )1()1(
dtdNQvNNQrQ
At higher temperatures also the reverse process must be considered being the net energy production:01232301
The time integrated released energy is obtained from:
finalinitial
N
N
NNN
NQdNQdtfinal
init
0010
1001
230110
01
23012301
)(
)()1(
)()1(
0
0
Decays and reactions in the stellar plasma
12Física Nuclear, Tema 15 José Benlliure
1.8 Nonresonant reactions induced by charged particles
max
02 12
l
lll T
kabs
21
2121
22210
1989534.0 2
2expMMMM
EZZeeZZ
EmT
hl
Absorption cross section of charged particles are dominated bythe 1/k2=1/E dependence term and the transmission coefficient of the Coulomb barrier Tl
Astrophysical S-factorIn order to performed reliable extrapolations of the measuredcross sections at energies of astrophysical interest nuclear Astrophysics introduces the astrophysics S-factor removing the 1/E dependence of the absorption cross sectionsand the s-wave Coulomb barrier probability.
)(1)( 2 ESeE
E
Decays and reactions in the stellar plasma
13Física Nuclear, Tema 15 José Benlliure
1.9 Nonresonant reactions induced by charged particlesWith the definition of the S-factor we can write the nonresonant reaction rate as:
0
/202/3
2/1
010
/22/3
2/1
01 )(8)(
)(8 dEeeS
kTN
mdEeESe
kTN
mvN kTEAkTEA
A
The Gamow peak:The major contribution to the reaction rate will come from energies where the product of the Maxwell-Boltzmann distribution and the Gamow factor is near its maximum. This product is known as the Gamow peak and represents the relatively narrow energy range over which most nuclear reactions occur in a stellar plasma.The location of the maximum of the Gamow peak (E=E0 ) can be obtained from the derivative of the product of these two terms:
01122
22/3
0
01210
210
01
0
kTE
meZZkTEeZZ
Em
dEd
EEhh
3/1
29
10
1021
20
3/120122
10
2
0 122.0)(2
T
MMMMZZkTmeZZE
h
Decays and reactions in the stellar plasma
14Física Nuclear, Tema 15 José Benlliure
1.10 Nonresonant reactions induced by neutronsNeutrons that are produced in a star quickly thermalise and their velocities are given by Maxwell-Boltzmann distributions. Although neutron induced reactions can lead to the emission of charged particles (e.g. n,p) in general these are exothermic reaction being the corresponding barrier transmission coefficients constants, then for s waves:
Ev
Tkabs
1112max
02
l
lll
Considering higher order partial waves the general expression for the reaction rates of reactions induced byneutrons assuming low neutron energies compared to the neutron binding energy is:
dEeEdEeEEkTN
mvN kTEkTEA
A
0
/2/1
0
/2/3
2/1
01)(
)(8 l
The integrand El+1/2e-E/kT represents the stellar energy window in which mostof the nonresonant neutron-induced reactions take place. The maximum of theIntegrand occurs at Emax =(l+1/2)kT
Decays and reactions in the stellar plasma
15Física Nuclear, Tema 15 José Benlliure
1.11 Nonresonant reactions induced by photonsFor nonresonant charged-particle emission reactions induced by photons the decay constant in terms of theastrophysical S-factor is given by the expression:
001
/2/0
001
01
2/
30101301)()()3( dEeeeESdEES
EeeQE kTEkTQkTE
Again, the integrand in this equation represents the Gamow peak centered around. 3010 QEEeff
If we consider now reactions emitting neutrons with an energy smaller than the neutron binding energy weobtain the following equation:
0
2/1301
/
0
2/101
/301)3(
dEQEedEEeQE kTEkTE ll
In this case the energy window for (,n) reactions will be located at
neff QkTE 2/1l
Decays and reactions in the stellar plasma
16Física Nuclear, Tema 15 José Benlliure
1.12 Narrow-resonance reaction ratesPossible resonances is the reaction cross sections could have an important impact on the reaction rates. If we consider narrow resonances (
less than few keV). Isolated resonances can be described by the Breit-Wigner formula:
4/)12)(12()1)(12(
4)( 22
10
012
r
baBW EEjj
JE
where ji are the spins of the target and projectile, J and E are the spin and energy of the resonance, i are the resonance partial widths of entrance andexit channel and is the total resonances width and. Emk 012/2/2 h
The reaction rate for a single narrow resonance can then be obtained as:
dEe
EEw
kTmNdEeEE
kTN
mvN kTE
r
baA
kTEBW
AA
0
/222/3
01
2
0
/2/3
2/1
01 4/)(2)(
)(8 h
where )11)(12(/)1)(12( 1001 jjJw
Decays and reactions in the stellar plasma
17Física Nuclear, Tema 15 José Benlliure
1.13 Total reaction ratesFor the calculation of the total reaction rates, all processes contributing significantly to the reaction mechanism in the effective stellar energy range have to be taken into account.
tnonresonanAk
k
resonancesbroadA
i
i
resonancesnarrowAtotalA vNvNvNvN
Decays and reactions in the stellar plasma
18
- t=10-6 s, T=1013 K (E~1000 MeV)Nucleons and anti-nucleons annihilate but are not created any moreleptons and neutrinos made possible the conversion between protons and neutrons via weakinteractions.
A mechanism should account for the present imbalance between matter and radiation (~10-9)and between matter and anti-matter CP violating decays.
First moments:- t=10-12 s, T=1016 K (E~1000 GeV)
At this temperature matter and radiation is in equilibrium, all species of particles and anti-particles are created but also annihilated.
Física Nuclear, Tema 15 José Benlliure
Primordial nucleosynthesis
2.1 The early Universe: particle physics era
?
,
nnppee
epn
enp
e
e
19
Proton and neutron concentrations are in equilibrium being more abundant protons becausethey are lighter than neutrons.
- t=10-2 s, T=1011 K (E~10 MeV)Electrons are the only remaining leptons, (mc2=105 MeV) and
(mc2=1784 MeV) are no longerproduced and they decay to electrons or annihilate but neutrinos are produced in neutral weak interactions.
Ffísica Nuclear, Tema 15 José Benlliure
2.1 The early Universe: particle physics era
o
oee
o
Zee
Zee
Zee
kTcmm
p
n pneNN /2
- t=1 s, T=1010 K (E~1 MeV)Neutrino interactions are no longer important and they expand freely.
Primordial nucleosynthesis
20Física Nuclear, Tema 15 José Benlliure
2.2 The early Universe: nuclear physics era- t=190 s, T=109 K (E<0.5 MeV)
Positrons were no longer produced and they annihilated, then the proportion of protons and neutrons was fixed.
7/1~87/13/ pn NN
At this moment, protons and neutrons underwent a series of nuclear reactions leading to the transformation of protons and neutrons in 4He with no free neutrons surviving. After this processthe Universe was composed of protons, helium, electrons, photons and neutrinos.
- t~500000 y, T=5 103 K (E~0.43 eV)Under these conditions, ions and electrons combined forming a neutral gas. From this moment onthe Universe changed from being radiation dominated to be matter dominated. The consequence is that the previously opaque Universe became transparent since radiation could travel unscatteredthrough space. Indeed, the background radiation we observe now was produced at that time.
Primordial nucleosynthesis
21
dnp
Física Nuclear, Tema 15 José Benlliure
The first nuclear reaction that took place in the Universe was the production of deuterium in the proton andneutron fusion reaction.
2.3 First nuclear reactions
At high temperatures, the reverse reaction occurs as quickly as deuterium production, and there is no effective accumulation of deuterium nuclei. Taking into account that the photon energy necessary for photo-dissociation is 2.225 MeV one can obtain that deuterium production is effective at temperatures below T=9 108 K and thus, for a time t=250 s.
The final abundance of deuterium will depend on the neutron half life but also on the ratio between photonsand baryons densities at the early Universe. The comparison between calculated and measured abundances allow to determine one of the main cosmological parameters, the baryon density of the Universe.
E
>2.225 MeV
1 2 3 4 5MeV
Primordial nucleosynthesis
22Física Nuclear, Tema 15 José Benlliure
Primordial nucleosynthesis
2.3 First nuclear reactionsd+n 3H+d+d 3H+pd+p 3He+d+d 3He+nn+3He 3H+pd+3H 4He+nd+3He 4He+p
3H+4He 7Li+3He+4He 7Be+n+7Be 7Li+p7Be 7Li+e-+p+7Li 8Be 4He+4He
The main reactions in the Big Bang nucleosynthesis process leads mostly to the transformation of the existing neutrons in the early Universe into 4He nuclei. Additionally, some deuterium 3He, 3H andheavier species as 7Be, 7Li and 6Li are also produced.
23Física Nuclear, Tema 15 José Benlliure
Primordial nucleosynthesis
2.4 Big Bang nucleosynthesis and cosmologyThe final abundances of different nuclear species produced duringthe Big Bang nucleosynthesys process are very much dependenton some fundamental cosmological parameters such as the baryondensity of the Universe.
Nuclear models allow to determine the expected abundances of lightnuclei produced after the Big Bang as a function of the Universebaryon density with high accuracy.
Primordial abundances obtained from astronomic abundances ordetailed characterizations of the temperature fluctuations in theCosmic background radiation (CBR) together with Big Bangnucleosynthesis models help in putting some light in fundamental questions such as:- the baryon density of the Universe- galactic chemical evolution- particle physics standard model
25
Ciclo p-pMeV 7.24224 4
eeHep
Física Nuclear, Tema 14 José Benlliure
BeHeHe 743
Estrellas de la secuencia principalT ~ 8-55 MK
3.1 Combustión hidróstatica de hidrógeno
Nucleosíntesis estelar: A<60
26
Ciclo CNO
HeCpN
eNOOpNNpC
eCNNpC
e
e
41215
1515
1614
1413
1313
1312
MeV 7.24224 4 eeHep
Física Nuclear, Tema 15 José Benlliure
Estrellas de la secuencia principal con una composición inicial que incluye C, N y OT ~ 8-55 MK
La captura de protones es mucho más lenta que las desintegraciones radiactivas
3.1 Combustión hidróstatica de hidrógeno
Nucleosíntesis estelar: A<60
27
Los ciclos en los que participan núcleos con A>20 solo son posibles si existen núcleos semillas de esa masa ya que la tasa de la reacción 19F(p,)20Ne es al menos tres órdenes de magnitud inferior al de la reacción19F(p,)16O
Física Nuclear, Tema 15 José Benlliure
Ciclos superioresNe-NaMg-AlSi-pS-Cl
3.1 Combustión hidróstatica de hidrógeno
Nucleosíntesis estelar: A<60
28
MgNeHeNeOHeOCHe
MeVQCBeHeMeVQBeHeHe
24204
20164
16124
1284
844
) 37.7( ) 091.0(
Física Nuclear, Tema 15 José Benlliure
Escenarios: estrellas AGB, T ~ 0.1-0.4 GK3.2 Combustión hidrostática de helio
Nucleosíntesis estelar: A<60
29
SOOnSOOpPOOHeSiOOHeMgOO
MgCCpNaCCHeNeCC
321616
311616
311616
4281616
4241616
241212
231212
4201212
2
Física Nuclear, Tema 15 José Benlliure
3.3 Combustión hidrostática de elementos pesados
Nucleosíntesis estelar: A<60
30
Escenarios: estrellas AGB masivas o sistemasbinarios (Novas), T ~ 0.1-0.4 GKLa captura de protones compite con la desintegración .
Física Nuclear, Tema 15 José Benlliure
Ciclo CNO caliente
3 4 5 6 7 8
9 10
111213
14
C (6)N (7)
O (8) F (9)
Ne (10)Na (11)
Mg (12)
3 4 5 6 7 8
9 10
111213
14
C (6)N (7)
O (8) F (9)
Ne (10)Na (11)
Mg (12)
3 4 5 6 7 8
9 10
111213
C (6)N (7)
O (8) F (9)
Ne (10)Na (11)
Mg (12)
T1/2 =1.7s
3
flow
T9 < 0.08 GK
T9 ~ 0.08-0.1 GK
T9 ~ 0.3 GK
4.1 Combustión explosiva de hidrógeno
Nucleosíntesis estelar: A<60
31
Escenarios: sistemas binarios (X-ray burst), T > 0.5 GKLa captura de protones compite con la desintegración .
Física Nuclear, Tema 15 José Benlliure
Ruptura del ciclo CNO
3 4 5 6 7 8
9 10
111213
14
C (6)N (7)
O (8) F (9)
Ne (10)Na (11)
Mg (12)
3
flow
T9 >~ 0.3 15O()19Ne18Ne(,p)21NaT9 >~ 0.6
4.2 Combustión explosiva de hidrógeno y helio
Nucleosíntesis estelar: A>60
32
- novas- X-ray bursts (sistemas binarios)
Escenarios astrofísicos
Reacciones- (p,)
Física Nuclear, Tema 15 José Benlliure
4.3 Captura rápida de protones: proceso rp
Nucleosíntesis estelar: A>60
33Física Nuclear, Tema 15 José Benlliure
4.2 Captura rápida de protones: proceso rp
Nucleosíntesis estelar: A>60
34Física Nuclear, Tema 15 José Benlliure
4.3 Procesos de captura de neutrones
Nucleosíntesis estelar: A>60
35Física Nuclear, Tema 15 José Benlliure
s-only
FeCoNi
Rb
GaGe
ZnCu
SeBr
As
ZrY
Sr
Kr(n,)
()
()
r-pro
cess
p-pro
cess
r-only
p-only
neutron number
prot
on n
umbe
r
4.3 Procesos de captura de neutrones
Nucleosíntesis estelar: A>60
36Física Nuclear, Tema 15 José Benlliure
4.3 Procesos de captura de neutrones
Nucleosíntesis estelar: A>60
37
r process
s process
p process
Física Nuclear, Tema 15 José Benlliure
4.3 Procesos de captura de neutrones
Nucleosíntesis estelar: A>60
38
Número de neutrones
Núme
rode
proto
nes
semilla(n,
(,n
n ~ 107-8 cm-3
n
n ~ 1024 cm-3
n
proceso condiciones escala escenarioproceso scaptura de n
T ~ 0.1 GKn ~ 1-103 a, n ~107-8cm-3
~ 1-105-6 a estrellas AGB
proceso rcaptura de n
T ~ 1-2 GKn ~ s, n ~1024-26cm-3
< 1 s supernovas IIestre. Neutrones
proceso pcaptura de p
T ~ 2-3 GK ~ 1 s supernovas II
Física Nuclear, Tema 15 José Benlliure
4.3 Procesos de captura de neutrones
Nucleosíntesis estelar: A>60
39Física Nuclear, Tema 15 José BenlliureFeCoNi
Rb
GaGe
ZnCu
SeBr
As
ZrY
Sr
Kr(n,)
()
()
63Ni, t1/2 =100 y64Cu, t1/2 =12 h, 40% (), 60% ()
79Se, t1/2 =65 ky
80Br, t1/2 =17 min, 92% (), 8% ()
85Kr, t1/2 =11 y
4.4 El proceso s de captura lenta de neutrones
Nucleosíntesis estelar: A>60
40Física Nuclear, Tema 15 José Benlliure
free neutrons are unstable
they must be produced in situ
most likely candidates as neutron source are:
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
4.4 El proceso s de captura lenta de neutrones
Nucleosíntesis estelar: A>60
41
N
Z
capacerrada
Línea deestabilidad
punto deespera
Camino del proceso r: equilibrio(n,)
(,n),masas
kT)(SkTmπ
A+A
A)G(Z,)+AG(Z,=
A)Y(Z,)+AY(Z,
nu
n /exp212
112/32
h
Física Nuclear, Tema 15 José Benlliure
4.5 El proceso r de captura rápida de neutrones
Nucleosíntesis estelar: A>60
42
N
Z
capacerrada
Línea deestabilidad
punto deespera
Camino del proceso r: equilibrio(n,)
(,n),masas
Velocidad del proceso r: desintegración ,vidas medias
Fin y ciclo del proceso r: fisión,barreras y modos de fisiónFísica Nuclear, Tema 15 José Benlliure
4.5 El proceso r de captura rápida de neutrones
Nucleosíntesis estelar: A>60
43
n >1023 cm-3
T~109 K Tiempo de irradiación Semilla inicial
Modelos independientes del escenario: estáticos o dinámicos
n = 1020
n = 1026
Entropía Tiempo de expansión Física nuclear
Física Nuclear, Tema 15 José Benlliure
4.5 El proceso r de captura rápida de neutrones
Nucleosíntesis estelar: A>60
44
Colapso de E.N: Supernovas IIFrecuen./año y galaxia 10-5 – 10-4 2.2 10-2
Masa ejectada (Mo ) 4 10-3 – 4 10-2 10-6 – 10-5
Colapso de un sistema binariode estrellas de neutronesSupernova tipo II
Modelos hidrodinámicos de evolución estelar
Física Nuclear, Tema 15 José Benlliure
4.5 El proceso r de captura rápida de neutrones
Nucleosíntesis estelar: A>60
45
-e epn
enpe
e- npe
,e,,e,-ee
He2p2n
C3He 12
Fe,...Si,Mg,,OXHe
T>1010 K
T~6 109 K
T~4 109 K
T~2 109 Kproceso r
proceso
n,12 C
n,pn,
n,O,Ne,Mg,Si,Fe,…
102 10103104105
R(km)
M(Mo )3
1,4
2
flujo
de ne
utrino
s(~
10 s)
Expansión adiabática Alta entropía Equilibrio nuclear Abundancia de neutrones
Física Nuclear, Tema 15 José Benlliure
4.5 El proceso r de captura rápida de neutrones: explosiones supernova
Nucleosíntesis estelar: A>60
46
Colapso de un sistema binario deestrellas de neutrones
Baja entropía Abundancia de neutrones
Física Nuclear, Tema 15 José Benlliure
4.5 El proceso r de captura rápida de neutrones
Nucleosíntesis estelar: A>60
47
Producción de núcleos muy lejos de la estabilidad
Captura de neutrones: (camino del proceso r)- Secciones eficaces de captura- Masas de los núcleos participantes (n,)
(,n)
Desintegración beta: (velocidad del proceso y curva de abundancias)- vidas medias- energías de la transición Q
(masas)- tipos de transición- emisión retardada de neutrones
Fisión: (fin y ciclo del proceso r)- Probabilidades de fisión (barreras y densidades)- Distribuciones de productos de fisión- fisión inducida por neutrones- fisión inducida por neutrinos- fisión retardada por la desintegración
N
Z
capacerrada
Línea deestabilidad
punto deespera
Física Nuclear, Tema 15 José Benlliure
4.5 El proceso r de captura rápida de neutrones: inputs de física nuclear
Nucleosíntesis estelar: A>60