2.1 Origin of the Elements
A Heger, Monash University, Monash, VIC, Australia
C Fro¨hlich, North Carolina State University, Raleigh, NC, USA
JW Truran, University of Chicago, Chicago, IL, USA; Argonne National Laboratory,
Argonne, IL, USA ã 2014 Elsevier Ltd. All rights reserved.
This article is a revision of the previous edition article by J.W. Truran and A. Heger, volume 1, pp. 1–15, © 2003, Elsevier Ltd.
2.1.1 Introduction 1 2.1.2 Abundances and Nucleosynthesis 2 2.1.3 IMS: Evolution and Nucleosynthesis 3 2.1.3.1
Shell Helium Burning and 12C Production 4 2.1.3.2 s-Process Synthesis in Red Giants 4 2.1.4 Massive Star Evolution
and Nucleosynthesis 4 2.1.4.1 Nucleosynthesis in Massive Stars 6 2.1.4.1.1 Hydrogen burning 6 2.1.4.1.2 Helium
burning and the s-process 6 2.1.4.1.3 Hydrogen and helium shell burning 7 2.1.4.1.4 Carbon burning 7 2.1.4.1.5 Neon
and oxygen burning 7 2.1.4.1.6 Silicon burning 7 2.1.4.1.7 Explosive nucleosynthesis 7 2.1.4.1.8 The p-process 8
2.1.4.1.9 The r-process 9 2.1.5 Type Ia Supernovae: Progenitors and Nucleosynthesis 9 2.1.6 Nucleosynthesis and
Galactic Chemical Evolution 10 References 12
2.1.1 Introduction
Nucleosynthesis is the study of the nuclear processes that
are responsible for the formation of the elements that
constitute the baryonic matter of the universe. The elements
of which the universe is composed indeed have a quite
complicated nucle osynthesis history, which extends from
the first 3 min of the Big Bang through to the present.
Contemporary nucleosynthe sis theory associates the
production of certain elements/iso topes or groups of
elements with a number of specific astrophysical settings,
the most significant of which are (1) the cosmological Big
Bang, (2) stars, and (3) supernovae.
Cosmological nucleosynthesis studies predict that the
condi tions characterizing the Big Bang are consistent with
the synthesis only of the lightest elements: 1H, 2H, 3He, 4He,
and 7Li (Iocco et al., 2009). These contributions define the
primordial compo sitions both of galaxies and of the first
stars formed therein. Within galaxies, stars and supernovae
play the dominant role both in synthesizing the elements
from carbon to uranium and in returning heavy element
enriched matter to the interstellar gas from which new stars
are formed. The mass fraction of our solar system (formed
some 4.6 billion years ago) in the form of heavy elements is
around 1.4% (Asplund et al., 2009) to 1.53% (Lodders et al.,
2009), and stars formed today in our galaxy can be a factor
2 or 3 more enriched (Edvardsson et al., 1993). It is the
processes of nucleosynthesis that operate in stars and
supernovae that we review here. We confine our attention to
three broad categories of stellar and supernova sites with
which specific nucleosynthesis products are understood to
be identified: (1)
intermediate-mass stars (IMS), (2) massive stars and
associated core collapse supernovae, and (3) Type Ia
supernovae. The first two of these sites are the
straightforward consequence of the evolution of single stars,
whereas Type Ia supernovae are under stood to result from
binary stellar evolution.
Stellar nucleosynthesis resulting from the evolution of sin
gle stars is a strong function of stellar mass (Woosley et al.,
2002). Following phases of hydrogen and helium burning, all
stars consist of a carbon–oxygen core. In the mass range of
the so-called ‘intermediate-mass’ stars (1M ≲M≲10M ), the
temperatures realized in their degenerate cores never reach
levels at which carbon or neon ignition can occur; hence, the
stars do not build up an iron core. Substantial element pro
duction occurs in such stars during the asymptotic giant
branch (AGB) phase of evolution (Herwig, 2005),
accompanied by significant mass loss, and they evolve to
white dwarfs of carbon–oxygen (or, less commonly,
oxygen–neon) composi tion. In contrast, the increased
pressures that are experienced in the cores of stars of
masses M≳10M yield higher core tem peratures that enable
subsequent phases of carbon, neon, oxy gen, and silicon
burning to proceed. Collapse of an iron core devoid of
further nuclear energy then gives rise to a core collapse
supernova, most commonly of Type II or Type Ib/c, and the
formation of a neutron star or black hole remnant (Heger et
al., 2003). The ejecta of core collapse supernovae contain
the ashes of nuclear burning of the entire life of the star,
except for the part that goes into the remnant, but are also
modified by the explosion itself (Rauscher et al., 2002). They
are the source of most material (by mass) heavier than
helium.
Treatise on Geochemistry 2nd Edition http://dx.doi.org/10.1016/B978-0-08-095975-7.00117-0 1
novae (involving hydrogen thermonuclear runaways in
accreted shells on white dwarfs; Gehrz et al., 1998), x-ray
bursts (hydrogen/helium thermonu clear runaways on
Observations reveal that binary stellar systems comprise neutron stars (Strohmayer and Bildsten, 2003)), and x-ray
roughly half of all stars in our galaxy. Single star evolution, binaries (accretion onto black holes). For some range of
as outlined in the preceding text, can leave in its wake conditions, accretion onto carbon–oxygen white dwarfs
compact stellar remnants: white dwarfs, neutron stars, and permits growth of the CO core from an initial mass of ≲1.1M
black holes. But we also have evidence for the occurrence of to the Chandrasekhar limit M ¼1.4M and a ther monuclear
Ch
all three types of condensed remnant in binaries. In close
runaway in the core leads to a Type Ia supernova.
binary systems, mass transfer can take place from an
In this chapter, we review the characteristics of
evolving companion onto a compact object. This naturally
thermonuclear processing in the three environments we
gives rise to a variety of interesting phenomena: classical
2 Origin of the Elements
identified: IMS, massive stars and core collapse
supernovae, and thermonuclear Type Ia supernovae. This
will be followed by a brief discussion of galactic chemical
evolution, which illustrates how the contributions from each
of these environments are first introduced into the inter
stellar media of galaxies. Reviews of nucleosynthesis
processes include those by Arnett (1995), Trimble (1975),
Truran (1984),
stellar and gas components of galaxies. Significant prog ress
has recently been achieved in this regard as a conse quence
of a wealth of new information on cosmic abundances –
spectroscopic properties of stars in our galaxy and of gas
clouds and galaxies at high redshifts – pouring in from new
ground- and space-based observatories. That being said, it
remains true that the most significant clues to
nucleosynthesis are those provided by our detailed
knowledge of the elemental and isotopic compo sition of
solar system matter. The mass fractions of the stable
isotopes in the solar system are displayed in Figure 1. Key
features that reflect the natures of the nuclear processes by
which the heavy elements are formed include (1) the large
abundances of 12C and 16O, the main products of stellar
helium burning; (2) the domi nance of the a-particle nuclei
through calcium (20Ne, 24Mg, 28Si, 32S, 36Ar, and 40Ca); (3)
the ‘nuclear statistical equilibrium’ peak at the position of
56
Fe; and (4) the abundance peaks in the region past iron at
the neutron closed shell positions (zirconium, barium,
Wallerstein et al. (1997), and Woosley et al. (2002). An
overview of galactic chemical evolution is presented by
Matteucci (2008).
2.1.2 Abundances and Nucleosynthesis
The ultimate goal of nucleosynthesis theory is to explain the
composition of the universe, as reflected, for example, in the
Ar
Cr
Fe
0
Mn
H
He
−2 −4
Co
C
Al
N
)n
Na
Mg
it
r
−6
c
o
Se
f
Si S
P
a
Ne
OF
Ti
Ni
Cu
Ca
Cl
Ga
s
Sc Zn
s
a
−8
Zr RbY BrKr
Mo
K Li
As Ge
Sr
Ru Nb
v
m
(
g
B
RhAg
Cd
o
L
Ce
−10
Pr
XeBa Te Cs Sb Gd
Nd
Sn
Eu
Sn
I
Pd
Er
Yb
Hf
Tb
Ho
Pt
Ir
Pb
Os
TI Bi AuHg
Tm
Re
Th
In
U
Be
La
Lu
Ta
−12
−14
W
Dy
0 50 100 Mass number
150 200
Figure 1 The abundances of the isotopes present in solar system matter are plotted as a function of mass number A. The solar
system abundances for the heavy elements are those compiled by Lodders et al. (2009).
Origin of the Elements 3
Se
Sr
ε
g
o
L
3.00 2.50 2.00 1.50
1.00
0.50 0.00 −0.50 −1.00 −1.50
SS s−process
SS r−process
60 80 100 120 140 160 180 200 220
Mass number
Figure 2 The s-process and r-process abundances in solar system matter, based upon the work by Ka¨ppeler et al. (1989). Note the
distinctive s-process signatures at masses A 88, 138, and 208 and the corresponding r-process signatures at A 130 and 195, all
attributable to closed shell effects on neutron-capture cross sections. It is the r-process pattern thus extracted from solar system
abundances that can be compared with the observed heavy element patterns in extremely metal-deficient stars. The total solar
system abundances for the heavy elements are those compiled by Anders and Grevesse (1989).
for the synthesis of a number of stable isotopes of nuclei on
the proton-rich side of the valley of beta stability. Our
subsequent
discussions
identify
the
astrophysical
environments in which these diverse processes are now
understood to occur.
and lead), confirming the occurrence of processes of
neutron capture synthesis. The solar system abundance
patterns associated specifically with the slow (s-process)
and fast (r-process) pro cesses of neutron-capture synthesis
are shown in Figure 2. It is important here to call attention to
the recently revised determi nations of the oxygen and
2.1.3 IMS: Evolution and Nucleosynthesis
carbon abundances in the Sun. Lodders et al. (2009)
derived an updated oxygen abundance for the sun of
Intermediate-mass red giant stars are understood to be the
loge(O)¼8.73 0.07 dex, a value consistent with Allende
primary source both of 12C and of the heavy s-process (slow
Prieto et al. (2001) but approximately a factor 2 below that
neutron-capture) elements, as well as a significant source of
quoted by Anders and Grevesse (1989). Lodders et al.
14
N and other less abundant CNO isotopes. Their contribu
(2009) also determined the solar carbon abundance to be
tions to galactic nucleosynthesis are a consequence of the
loge(C)¼8.39 0.04 dex and the ratio C/O¼0.457. Changes
occurrence of nuclear reactions in helium shell thermal
of similar magnitude were found by Asplund et al. (2009).
pulses on the AGB, the subsequent dredge-up of matter into
The implications of these results for stellar evolution,
the hydrogen-rich envelope by convection, and of mass loss.
nucleosynthesis, the formation of carbon stars, and galactic
This is a very complicated evolution. Current stellar models,
chemical evolution remain to be explored in detail;
recently reviewed by Cristallo et al. (2011), allow the forma
nevertheless, in Figure 5, we show the nucleosynthesis
tion of low-mass ( 1.5Μ ) carbon stars, which represent the
result using this new composition.
main source of s-process nuclei. A detailed review of the
Guided by early compilations (e.g., Suess and Urey,
theo retical models, the nuclear physics inputs, and the
1956) Burbidge et al. (1957) and Cameron (1957) first
observa tional constraints is provided by Ka¨ppeler et al.
identified the nuclear processes by which element formation
(2011).
occurs in stellar and supernova environments: (1) hydrogen
Red giant stars have played a significant role in the histor
burning, which powers stars for approximately 90% of their
ical development of nucleosynthesis theory. Whereas the piv
lifetimes; (2) helium burning, which is responsible for the
otal role played by nuclear reactions in stars in providing an
production of 12C and 16O, the two most abundant elements
energy source sufficient to power stars like the Sun over
heavier than helium; (3) the alpha-process, which we now
billions of years was established in the late 1930s, it
understand as a combination of carbon, neon, and oxygen
remained to be demonstrated that nuclear processes in
burning; (4) the equilibrium process, by which silicon burning
stellar interiors might play a role in the synthesis of heavy
proceeds to the formation of a nuclear statistical equilibrium
nuclei. The recog nition that heavy element synthesis is an
abundance peak centered on mass A¼56; (5) the slow
ongoing process in
(s-process) and rapid (r-process) mechanisms of
4
Origin of the Elements
neutron-capture synthesis of the heaviest elements (A≳60
70); and (6) the p-process. We
now understand the p-process as a combination of the
stellar interiors followed the discovery by Merrill (1952) of
g-process, the n p-process, and the n-process, responsible
the presence of the element technetium in red giant stars.
Since technetium has no stable isotopes and the longest
lived isotope has a half-life t1/2 4.6 106 years, its presence in
red giant atmospheres indicates its formation in these stars.
This confirmed that the products of nuclear reactions operat
ing at high temperatures and densities in the deep interior
can be transported by convection to the outermost regions of
the stellar envelope.
The role of such convective ‘dredge-up’ of matter in the
red giant phase of evolution of 1M 10M stars is now under
stood to be an extremely complex process (Busso et al.,
1999). On the first ascent of the giant branch (prior to helium
ignition), convection can bring the products of CNO cycle
burning (e.g., 13C, 17O, and 14N) to the surface. A second
dredge-up phase occurs following the termination of core
helium burning. The critical third dredge-up, occurring in the
aftermath of thermal pulses in the helium shells of these
AGB stars, is responsible for the transport of both 12C and
s-process nuclei (e.g., technetium) to the surface. The
subsequent loss of this enriched envelope matter by winds
and planetary nebula formation serves to enrich the
interstellar media of galaxies, from which new stars are
born. A brief review of the mechanisms of production of 12C
and the ‘main’ component of the s-process of
neutron-capture nucleo synthesis is presented in the
following sections.
2.1.3.1 Shell Helium Burning and 12C Production
Stellar helium burning proceeds by means of the
‘triple-alpha’ reaction in which three 4He nuclei are
converted into 12C, later followed by the 12C(a,g)16O reaction
that forms 16O at the expense of 12C. Core helium burning in
massive stars (M≳10M ) burns to helium exhaustion, with the
last 10% of helium almost exclusively converting carbon to
oxygen. Additionally, it is followed by carbon and neon
burning phases that convert carbon to oxygen and heavier
elements. Typically, the average 16O/12C ratio in the core of
a massive star prior to collapse is a factor 2 3 higher than
the solar value. Massive stars are thus the major source of
oxygen, whereas low- and IMS dominate the production of
carbon.
The advantage of the helium shells of low- and IMS for
12
C production arises because the 12C/16O ratio is high for
incom plete helium burning. Following a thermal pulse in the
helium shell, convective dredge-up of matter from the helium
shell brings helium, s-process elements, and a significant
mass of 12C to the surface. It is the surface enrichment
associated with this source of 12C that leads to the condition
that the envelope 12C/16O ratio exceeds one such that a
‘carbon star’ is born. Calculations of galactic chemical
evolution indicate that this source of carbon is sufficient to
account for the level of 12C in galactic matter. The levels of
production of 14N, 13C, and other CNO isotopes in this
environment are significantly more difficult to estimate, and
thus, the corresponding contributions of AGB stars to the
galactic abundances of these isotopes remain uncertain.
2.1.3.2 s-Process Synthesis in Red Giants
The formation of most of the heavy elements occurs in one
of two processes of neutron capture: the s-process and the
r-process. These two broad divisions are distinguished on
the basis of the relative lifetimes for neutron captures (tn)
and electron decays (tb). The condition that tn>tb, where tb is
a characteristic lifetime for beta-unstable nuclei near the
valley of beta stability, ensures that as captures proceed, the
neutron capture path itself remains close to the valley of
beta stability. This defines the s-process. In contrast, when
tn<tb, it follows that successive neutron captures proceed
into the neutron-rich regions off the beta-stable valley.
Following the exhaustion of the neutron flux, the capture
products approach the position of the valley of beta stability
by beta decay, forming the r-process nuclei. The s-process
and r-process patterns in solar system matter are those
shown in Figure 2.
The environment provided by thermal pulses in the
helium shells of IMS on the AGB provides conditions
consistent with the synthesis of the bulk of the heavy
s-process isotopes through bismuth. Neutron captures in
AGB stars are driven by a combi nation of neutron sources:
the 13C (a, n) 16O reaction provides the bulk of the neutron
budget at low neutron densities, whereas the
22
Ne (a, n)
25
Mg operates at high temperatures and helps to set the
timescale for critical reaction branches. This s-process site
(the main s-process component) is understood to operate in
low mass AGB stars (1M ≲M≲3M ) and to be responsible for
the synthesis of the s-process nuclei in the mass range
A≳90. Calcu lations reviewed by Busso et al. (1999) indicate
a great sensitivity both to the characteristics of the 13C
‘pocket’ in which neutron production occurs and to the initial
metallicity of the star. In their view, this implies that the solar
system abundances are not the result of a unique s-process
but rather the consequence of a complicated galactic
chemical evolutionary history that wit nessed mixing of the
products of s-processing in stars of different metallicity and a
range of 13C pockets. We can hope that obser vations of the
s-process abundance patterns in stars as a function of
metallicity will ultimately be able to better guide and
constrain such theoretical models.
2.1.4 Massive Star Evolution and Nucleosynthesis
Generally speaking, the evolution of a massive star follows a
well understood path of contraction to increasing central den
sity and temperature. The contraction is caused by the
energy loss of the star due to light radiated from the surface
and neutrino losses (see the succeeding text), and the
released potential energy is in part converted into internal
energy of the gas (virial theorem). This path of contraction is
interrupted by nuclear fusion – first, hydrogen is burned to
helium, then helium to carbon and oxygen. This is followed
by stages of carbon, neon, oxygen, and silicon burning, until
finally a core of iron is produced, from which no more energy
can be extracted by nuclear burning. The onsets of these
burning phases as the star evolves through the
temperature-density plane are shown in Figure 3 for stars of
masses 15M and 25M . Each fuel burns first in the center of
the star, then in one or more shells (Figure 4). Most burning
stages proceed con vectively: the energy production rate by
the burning is so large and centrally concentrated that the
energy cannot be trans ported by radiation (heat diffusion)
alone, and convective motions dominate the heat transport.
The reason for this is the high temperature sensitivity of
nuclear reaction rates: for hydrogen burning in massive
stars, nuclear energy generation has a dependence /T18,
and the dependence is even stronger for later burning
central helium burning is
greatly reduced by thermal
to the efficient mixing caused neutrino losses that carry away
by the convection, the entire energy in situ, instead of
requiring that it be transported
unsta ble region evolves
to the stellar surface by
essentially chemically
homogeneously – replenishing diffusion or convection. These
losses increase with
the fuel at the bottom of the
9
burning region (central or shell temperature as roughly /T to
18
burning) and depleting the fuel /T . When the star has built up
a
large
enough iron core,
elsewhere in that region at the
same time. As a result, shell exceeding
8.0
7.5
)
K(
e
r
u
t
a
r
e
25 M
p
stages. The important consequence is that, due
m
15 M
burning
of this fuel then
commences outside that
He region.
In Table 1, we summarize the
burning stages and their
H
durations for a 20M star. The
timescale for helium burning is
about ten times shorter than
that of hydrogen burning,
0 2 4 6 8 10 Log central density (g
mostly because of the lower
−3
cm )
energy release per unit mass.
The timescale of the burning
Origin of the Elements 5
stages and contraction beyond
e
t
l
a
r
t
n
e
c
g
o
L
10.0 9.5
9.0
8.5
Figure 3 Evolution of the central temperature and density in
stars of 15 and 25 M from birth as hydrogen burning stars until
iron core collapse (Table 1). In general, the trajectories follow a
line of r/T3 but with some deviation downward (toward higher r
at a given T ) due to the decreasing entropy of the core.
Nonmonotonic behavior is observed when nuclear fuels are
ignited, and this is exacerbated in the 15 M model by partial
degeneracy of the gas. Reproduced from Woosley SE, Heger A,
and Weaver TA (2002) The evolution and explosion of massive
stars. Reviews of Modern Physics 74: 1015–1071.
its effective Chandrasekhar mass, the maximum mass for
which such a core can be stable, the core collapses to form
a neutron star or a black hole (see Woosley et al., 2002, for
a more extended review). A core collapse supernova
explosion may result (e.g., Colgate and White, 1966) that
ejects most of the layers outside the iron core, including
many of the ashes from the preceding burning phase. When
the supernova shock
loss (burning + neutrino losses; in
erg g−1 s−1 )
Net nuclear energy
20
p
<−1014 <−1013 <−1012 <−1011 <−1010 <−109 <−108 <−107 <−106 <−105 <−104 <−103 <−102 <−101
<−100 <−10−1 >−10−1
e
)t
l
e
o
n
Net nuclear
v
energy
generation
(burning +
neutrino losses; in
erg g−1 s−1 )
15
(reduces by mass loss)
10
>10−1 >100 >101 >102 >103 >104 >105 >106 >107 >108 >109 >1010
>1011 >1012 >1013 <10−1 >1014
a
M
i
/
g
m
n
Convective envelope
(red supergiant)
e
e
u
l
e
b
v
(
i
Total mass of the star
t
H shell
burning
a
i
d
a
Convection Semiconvection
R
5
g
g
n
He shell burning
i
n
C shell burning
g
n
i
O shell burning
n
r
u
g
ll
e
h
s
i
b
n
r
i
n
r
u
u
H
b
v
i
u
t
e
H
a
e
r
b
d
C
n
b
0
64
O
)
n
i
r
(
Ne
C
burning
O
Si
O
O Si
2 0 −2 −4 −6 −8 Log (time till core collapse per year)
Figure 4 Interior structure of a 22 M star of solar composition as a function of time (logarithm of time till core collapse) and enclosed
mass. Green hatching and red cross hatching indicate convective and semiconvective regions. Convective regions are typically well
mixed and evolve chemically homogeneously. Blue shading indicates energy generation and pink shading energy loss. Both take into
account the sum of nuclear and neutrino loss contributions. The thick black line at the top indicates the total mass of the star, being
reduced by mass loss due to stellar winds. Note that the mass loss rate actually increases at late times of the stellar evolution. The
decreasing slope of the total mass of the star in the figure is due to the logarithmic scale chosen for the time axis.
6 Origin of the Elements
Table 1 Hydrostatic nuclear burning stages in massive stars
Fuel Main products Secondary products T (109 K) Duration (year) Main reaction
H He 14N 0.037 8.1 106 4H! 4He (CNO cycle)
He O, C 18O, 22Ne 0.19 1.2 106 34He! 12C s-process 12Cþ4He! 16O
C Ne, Mg Na 0.87 9.8 102 12Cþ12C!... Ne O, Mg Al, P 1.6 0.60 20Ne! 16Oþ4He 20Neþ4He! 24Mg
O Si, S Cl, Ar, 2.0 1.3 16Oþ16O!... K, Ca
28
24
4
Si Fe Ti, V, Cr, 3.3 0.031 Si! Mgþ He ... Mn, Co, Ni 28Siþ4He! 24Mg...
The table gives burning stages, main and secondary products (ashes), typical temperatures and burning timescales for a 20 M star, and the
main nuclear reactions. An ellipse (...) indicates more than one product of the double carbon and double oxygen reactions, and a chain of
reactions leading to the buildup of iron group elements for silicon burning.
Table 2 Explosive nucleosynthesis in supernovae
Fuel Main products Secondary products T (109 K) Duration (seconds) Main reaction
Innermost ejecta r-process – >10 1 (n, g), b (low Ye)
np-process (high Ye) – >10 1 (n, p), (p, g) Si, O 56Ni Iron group >4 0.1 (a, g) O Si, S Cl, Ar, K, Ca 3–4 1 16Oþ16O Ne, O O, Mg, Ne Na,
Al, P 2–3 5 (g, a) Heavy elements p-process, 2–3 5 (g,n) n-process: – 5 (n, n0),(ne,e )
11
B, 19F, 138La, 180Ta
Similar to Table 1, but for the explosion of the star. The ‘(A,B)’ notation means ‘A’ is on the ingoing channel and ‘B’ is on the outgoing channel.
An ‘a’ is the same as 4He, ‘g’ denotes a photon, that is, a photodisintegration reaction when on the ingoing channel, ‘n’ is a neutron, b shows
b-decay, and n indicates neutrino-induced reactions.
front travels outward, however, for a brief time, peak temper
atures are reached that exceed the maximum temperature
that has been reached in each region in the preceding
hydrostatic burning stages (Table 2). This defines the
transient stage of ‘explosive nucleosynthesis’ that is critical
to the formation of an equilibrium peak dominated by 56Ni
(Truran et al., 1967).
Massive stars build up most of the heavy elements from
oxygen through the iron group from the initial hydrogen and
helium of which they are formed. They also make most of
the s-process heavy elements up to atomic mass number
A¼80 90 from initial iron, converting initial carbon, oxygen,
and nitrogen into 22Ne, thus providing a neutron source for
the s-process. Core collapse supernovae from massive stars
are also the site of the n p-process that synthesizes
neutron-deficient nuclei with atomic mass number A 90.
Massive stars are probably also the source of most of the
proton-rich isotopes of atomic mass number greater than
100 (p-process) and the site of the r-process that is
responsible for many of the neutron rich isotopes from
barium through uranium and thorium.
2.1.4.1 Nucleosynthesis in Massive Stars
The most abundant product of the evolution of massive stars
is oxygen, 16O in particular – the third most abundant
isotope in the universe and the most abundant ‘metal.’
Massive stars are also the main source of most heavy
elements up to atomic mass number A 80, of some of the
rare proton-rich nuclei,
and of the r-process nuclei from barium to uranium. In the
following, we briefly review the burning stages and nuclear
processes that characterize the evolution of massive stars
and the resulting core collapse supernovae.
consumed until the death of the star, some s-processing
may start here as well, as a consequence of the higher
The first stage of stellar burning and energy generation is temperatures characteriz ing the shell burning phase. By the
hydrogen burning, during which hydrogen is transformed to time that the stable burning phases of evolution are
helium by the so-called CNO cycle. An initial abundance of completed and the iron core is on the verge of collapse, this
carbon, nitrogen, and oxygen (of which 16O is the most abun shell consists of a mixture of helium, carbon, and oxygen,
dant) is collected into 14N – the proton capture on this each of which comprises several tens of percent of the
isotope is the slowest reaction in this cycle. The total matter.
number of CNO isotopes remains unchanged – they operate
only as nuclear catalysts in the conversion of hydrogen into 2.1.4.1.4 Carbon burning
helium.
After a contraction phase of several tens of thousands of
years, carbon burning starts in the center. This produces
2.1.4.1.2 Helium burning and the s-process
mostly 20Ne by 12C (12C, a) 20Ne but also some 24Mg by 12C
After hydrogen is depleted, the star contracts toward helium 12
( C, g) 24Mg and 23Na by 12C (12C, p) 23Na. There is also
burning. Before helium burning is ignited, the 14N is
some continuation of the s-process from ‘unburned’ 22Ne
converted into 18O by 14N (a, g) 18F (bþ) 18O and then burned after the end of central helium burning, operating with
to 22Ne by 18O (a, g) 22Ne, where a stands for a 4He nucleus neutrons released during car bon burning via the reaction
and the bracket notation means that the first particle is on 12C (12C, n) 23Mg or from pro tons released in the 23Na
the inward channel and the second on the outward channel. branch that eventually can be converted into neutrons by
The first stage of helium burning involves the interaction of various possible reaction chains.
three alpha particles to form carbon (4He (a, g) 8Be∗ (a, g)
12
C; ‘3 a reaction’). (The asterisk on 8Be∗ indicates that this
2.1.4.1.5 Neon and oxygen burning
is a very short-lived state with a lifetime of only 10 14
Neon burning is induced by photodisintegration of 20Ne into
seconds. In most cases, when formed, it decays back into
an a-particle and 16O. The a-particle is then captured by
two 4He nuclei;
another 20Ne nucleus to make 24Mg or by 24Mg to make 28Si.
Secondary products include 27Al, 31P, and 32S.
2.1.4.1.1 Hydrogen burning
Origin of the Elements 7
comparably rarely, the second reaction occurs. The reaction
chain could hence be written as 2aP8Be∗(a,g)12C.) As
helium becomes depleted, the a-capture on 12C takes over
and pro duces 16O. Indeed, as helium is entirely depleted in
the central regions of the star, most of the ash is 16O, while
12
C only comprises 10 20%, decreasing with increasing
stellar mass.
Toward the end of the central helium burning phase, the
temperature becomes high enough for the 22Ne (a, n) 25Mg
reaction to proceed, which provides a neutron source for the
so-called ‘weak’ component of the s-process. In massive
stars, this process builds up elements of atomic mass
numbers of up to A 80 90, starting with neutron capture on
original 56Fe present in the interstellar medium from which
the star was born. In this environment, the timescale for
neutron capture is slow compared to the weak (b ) decay of
radioactive nuclei produced by the capture. The s-process
path thus proceeds along the neutron-rich side of the valley
of stable nuclei. The main competitor reaction for the
neutron source is the destruc tion channel for 22Ne that does
not produce neutrons, 22Ne (a, g) 26Mg. The main neutron
sink or ‘poison’ for the s-process in this environment is
neutron capture on the progeny of the 22Ne (a, n) reaction,
among which 25Mg is itself a major neutron poison.
2.1.4.1.3 Hydrogen and helium shell burning During
central helium burning, hydrogen continues to burn in a shell
at the outer edge of the helium core, leaving the ashes –
helium – at the base of the shell and increasing the size of
the helium core. This growth of the helium core by the
continued addition of fresh helium fuel contributes to the
destruction of carbon at the end of helium burning.
At the completion of core helium burning, helium con
tinues to burn in a shell overlying the helium-free core. It first
burns radiatively but then becomes convective, especially in
more massive stars. Since helium is not completely
This is closely followed by oxygen burning, for which the
dominant reaction is 16O þ 16O. The significant products are
the a-particle nuclei 28Si, 32S, 36Ar, and 40Ca, together with
isotopes of odd-charge number nuclei, such as 35Cl, 37Cl,
and 39Cl. Early during this phase, usually 24Mg is burned as
well by photodisintegration and a captures similar to neon
burning.
2.1.4.1.6 Silicon burning
Finally,
28
Si (and
32
S) is burned in a similar way as
28
20
Ne:
photodisintegration of
Si and resulting lighter isotopes
accompanies the gradual buildup of iron-peak nuclei with
higher binding energies per nucleon. At the same time, weak
processes – positron decay and electron capture – also
become important, producing a neutron excess and allowing
the for mation of nuclei along the valley of stability. The
ultimate result is the formation of an iron ‘equilibrium’ peak
of nuclei of maximal nuclear binding energies, from which no
further energy can be gained by means of nuclear fusion.
The extent in mass of convective central silicon burning is
typically about 1.05M . Several brief stages of shell silicon
burning follow until the star has built a core of iron group
elements (‘iron core’) large enough to collapse, typically 1.3
to ≳2M . The core had been held up, in part, by electron
degeneracy pressure, but when their Fermi energy becomes
sufficiently large, they can be captured on protons. This
serves both to neutronize the core and to reduce pressure
support, ultimately leading to the collapse of the core to a
neutron star or black hole. Noteworthy, photodisintegration
of iron nuclei plays a role during the collapse itself as well. It
follows that virtually all of the iron-peak products of this
phase of quasihydrostatic silicon burning do not escape from
the star and thus do not contribute to galactic
nucleosynthesis. It is the overlying regions of the core that
will generally be ejected.
may be said of the nuclei in the iron-peak region from
approximately 48Ti to 64Zn. A unique signature of such
Following the collapse of the core, a supernova shock front, massive star and core collapse supernova nucleosynthesis,
driven by energy deposition by neutrinos from the hot proto however, is that nuclei in the 16O to 40Ca region are
neutron star, travels outward through the star, setting the overproduced by a factor 2 3 with respect to nuclei in the
stage for a brief phase of ‘explosive nucleosynthesis.’ In the 48Ti to 64Zn region – relative to solar system abundances.
inner regions, in silicon and oxygen-rich layers, the We will see how this signature, reflected in the compositions
photodisintegra tion of silicon into a-particles and free of the oldest stars in our galaxy, provides impor tant
nucleons is accompanied by the capture of these particles constraints on models of galactic chemical evolution.
onto silicon and heavier nuclei, leading to the formation of a
nuclear statistical equilibrium peak centered on mass A¼56.
2.1.4.1.8 The p-process
Since the material here is characterized by almost equal
The ‘p-process’ generally serves to include all possible
numbers of neutrons and protons, it follows that the final
products of these explosive burning episodes must lie along mecha nisms that can contribute to the formation of
or very near the line of nuclei with equal number of protons proton-rich
2.1.4.1.7 Explosive nucleosynthesis
and neutrons (Z¼N) as well. The dominant species in situ
therefore include the nuclei 44Ti, 48Cr, 52Fe, 56Ni, 60Zn, and
64
Ge. Following decay, these contribute to the production of
the most proton-rich stable isotopes at these mass numbers:
44
Ca, 48Ti, 52Cr, 56Fe, 60Ni, and 64Zn. The neutron
enrichment charac teristic of matter of solar composition also
allows the production of the odd-Z species 51V (formed as
51
Mn), 55Mn (formed as 55Co), and 59Co (formed as 59Cu).
Further out, in the oxygen, neon, and carbon shells,
explo sive burning can act to modify somewhat the
abundance pat terns resulting from earlier hydrostatic
burning phases. The final nucleosynthesis products of the
evolution of massive stars and associated core collapse
supernovae have been
8 Origin of the Elements
discussed by a number of authors (Limongi et al., 2000;
Nomoto et al., 1997; Rauscher et al., 2002; Thielemann et
al., 1996; Woosley and Heger, 2007; Woosley and Weaver,
1995; Woosley et al., 2002). Most of the elements and
isotopes in the region from 16O to 40Ca are formed in relative
proportions consistent with solar system matter. The same
isotopes of heavy nuclei. Three processes making
proton-rich nuclei in massive stars have been identified:
g-process, the n p process, and the n-process.
The g-process involves the photodisintegration of heavy
elements. Obviously, this process is not very efficient, as the
effective ‘seed’ nuclei are the primordial heavy element
constit uents of the star. Most of the heavy p-elements with
atomic mass number greater than 100 that are produced in
massive stars (see Figure 5) are formed by this mechanism.
It is pri marily for this reason that these proton-rich heavy
isotopes are rare in nature.
The n p-process is a primary process occurring in the
early proton-rich ejecta of core collapse supernovae. These
ejecta are subject to a large neutrino and antineutrino flux
from the cooling proto-neutron star that is born in the
explosion. In the proton rich ejecta, charged-particle
reactions (alpha and proton cap tures) produce
neutron-deficient and proton-rich nuclei up to 56Ni and 64Ge.
Due to their long half-life, these two nuclei act as
Ne
P
K
SMg
Fe Co
100
HeB Li
H
10 1
C
Na
N
OF
Al
Si
Cl
V Cr
Ar
CaSc
Mn
Ti
10 20 30 40 50 0
Ni
Fe Co
Cu
)
a
V Cr
Zn
t
c
e
j
e
Mn
Ga Ge
As
Ag Mo Ru Rh Pd Sr
(
Br
Se
Kr
Y
Nb Rb
10
1
100
50 60 70 80 90 100
r
o
t
c
Cs Xe
Ba
I
Sn
Te Sb Cd
Ag Pd
Tb
Nd
Sm Rh
In Gd
Eu Ce Pr
La
a
100
f
n
o
i
t
c
u
10
d
o
r
P
1
100 110 120 130 140 150
Ta
LuHf
Gd
Tb
Dy
HoEr Tm
Yb
Tl
Pb
Bi
Hg
Pt
Au
Ir
Os
Re W
Eu
100
10
1
150 160 170 180 Mass number
190 200
Figure 5 Production factor in the ejecta of a 15 solar mass star based on Lodders et al. (2009) initial abundances with an explosion
energy of 1.2 1051 erg. The other model parameters are similar to the solar survey in Woosley and Heger (2007). The production
factor is defined as the average abundance of the isotope in the stellar ejecta divided by its solar abundance. To guide the eye, we
mark the production factor of 16O by a dashed line and a range of ‘acceptable coproduction,’ that is, within a factor 2 of 16O, by dotted
lines. The absolute value of the production factors does not matter so much, as the ejecta will be diluted with the interstellar material
after the supernova explosion. However, to reproduce a solar abundance pattern for certain isotopes in massive stars, they should be
closely coproduced with the most abundant product of massive stars, 16O.
bottlenecks for the nucleosynthesis unless enough neutrons
are available for (n, p) reactions to occur. Under the
conditions in the supernova ejecta, the (n, p) reactions take
place on much shorter timescales than the b-decays, hence
effectively accelerat ing the nucleosynthesis. The neutrons
necessary for these (n, p) reactions come from antineutrino
absorption reactions on free protons that are very abundant.
This n p-process can produce elements up to atomic
number Z¼60 and may explain the origin of some of the
light p-nuclei.
The rarest of these proton-rich isotopes have a different
and more exotic source: the immense neutrino flux from the
form ing hot neutron star can either convert a neutron into a
proton (making, e.g., 138La and 180Ta) or knock out a
nucleon (e.g., 11B and 19F). Since the neutrino cross
sections are quite small, it is clear that the ‘parent’ nucleus
must have an abundance that is at least several thousand
times greater than that of the neutrino-produced ‘daughter’ if
this process is to contribute significantly to nucleosynthesis.
This is indeed true for the cases of 11B (parent 12C) and 19F
(parent 20N e). The produc tion of 138La and 180Ta, the two
rarest stable isotopes in the universe, is provided by
neutrino interactions with 138Ba and 180Hf, respectively.
2.1.4.1.9 The r-process
It is generally accepted that the r-process synthesis of the
heavy neutron-capture elements in the mass regime A≳130
140 occurs in an environment associated with massive
stars. This results from two factors: (1) the two most
promising mecha nisms for r-process synthesis – a
neutrino-heated ‘hot bubble’ and neutron star mergers – are
both tied to environments associated with core collapse
supernovae, and (2) observations of old stars (discussed in
Section 2.1.6) confirm the early entry of r-process isotopes
into galactic matter.
•
The r-process model that has received the greatest study
in
•
recent years involves a high-entropy (neutrino-driven)
wind from a core collapse supernova (Takahashi et al.,
1994; Woosley et al., 1994). The attractive features of
this model include the facts that it may be a natural
conse quence of the neutrino emission that must
accompany core collapse in collapse events, that it would
appear to be quite robust, and that it is indeed associated
with massive stars of short lifetime. Recent calculations
have, however, called attention to a significant problem
associated with this mechanism: the entropy values
predicted by current core collapse supernova models are
too low to yield the correct levels of production of both
the lighter and heavier r-process nuclei.
The conditions estimated to characterize the
decompressed
ejecta from neutron star mergers
(Lattimer et al., 1977; Rosswog et al., 1999) may also be
compatible with the production of an r-process
abundance pattern generally consistent with solar system
matter. The most recent numerical study of r-process
nucleosynthesis in matter ejected in such mergers
(Freiburghaus et al., 1999) shows specifically that the
r-process heavy nuclei in the mass range A≳130 140 are
produced in solar proportions. Here again, the
association with a massive star/core collapse supernova
environment is consistent with the early
Origin of the Elements 9
appearance of r-process nuclei in the galaxy, and the
mech anism seems quite robust.
production of at least the light (A≲130 140) r-process
isotopes (see also Truran and Cowan, 2000). An analysis of
nucleosynthesis in massive stars by Rauscher et al. (2002)
has found, however, that this yields only a slight
redistribution of heavy mass nuclei at the base of the helium
shell and no significant production of r-process nuclei above
mass A¼100. This environment may, nevertheless, pro vide
a source of r-process-like anomalies in grains.
2.1.5 Type Ia Supernovae: Progenitors
and Nucleosynthesis
Two broad classes of supernovae are observed to occur in
the universe: Type Ia and core collapse. We have learned
from our discussion in the previous section that core
collapse superno vae are products of the evolution of
massive stars (M≳10M ). Observationally, the critical
distinguishing feature of Type I supernovae is the absence of
hydrogen features in their spectra at maximum light.
Theoretical studies focus attention on models for Type I
events involving either exploding white dwarfs (Type Ia) or
the explosions of massive stars (similar to those discussed
previously) that have, via wind driven mass loss or binary
effects, shed virtually their entire hydrogen enve lopes prior
to (iron) core collapse (e.g., Types Ib and Ic). We are
concerned here with supernovae of Type Ia, a subclass of
supernovae Type I, which are understood to be associated
with the explosion of a white dwarf in a close binary system.
The standard model for SNe Ia involves specifically the
growth of a carbon–oxygen white dwarf to the
Chandrasekhar-limit mass in a close binary system and its
subsequent incineration (see, e.g., the review of Type Ia
progenitor models by Livio, 2000).
The iron-peak nuclei observed in nature had their origin
in supernova explosions. Core collapse and Type Ia
supernovae provide the dominant sites in which this
‘explosive nucleosynthesis’ mechanism is known to operate.
These two supernova sites operate on distinctively different
timescales and eject different amounts of iron. As we shall
see, an understanding of the detailed nucleosynthesis
patterns emerging from these two classes of events
provides important insights into the star forma tion histories
of galaxies. Core collapse supernovae produce both the
intermediate-mass nuclei from oxygen to calcium and iron
peak
nuclei.
Calculations
of
charged-particle
nucleosynthesis in massive stars and core collapse
supernovae over the past decade (Limongi et al., 2000;
Nomoto et al., 1997; Rauscher et al., 2002; Thielemann et
al., 1996; Woosley and Heger, 2007; Woosley and Weaver,
1995) reveal one particularly significant distinguishing
feature of the emerging abundance patterns: the elements
from oxygen through calcium (and titanium) are
overproduced relative to iron (peak nuclei) by a factor 2 3.
This means that core
10 Origin of the Elements
Massive stars may also contribute to the abundances of
collapse produces only approximately 1/3 to 1/2 of the iron
the lighter r-process isotopes (A≲130 140). The helium and
in galactic matter. (We note that, whereas all such models
car bon shells of massive stars undergoing supernova
necessarily make use of an artificially induced shock wave
explosions can give rise to neutron production, via reactions
via thermal energy deposition (Thielemann et al., 1996) or a
13
16
18
21
22
25
such as C (a, n) O, O (a, n) Ne, and Ne (a, n) Mg,
piston (Woosley and Weaver, 1995), the general trends
involving residues of hydrostatic burning phases. Early
obtained from such nucleosyn thesis studies are expected to
studies of this problem (Hillebrandt et al., 1981; Truran et al.,
1978) indicated that these conditions might allow the be valid.) As we shall see in our discussion of chemical
evolution, these trends are reflected in the abundance
patterns of metal-deficient stars in the halo of our galaxy
(McWilliam, 1997, 2010; Wheeler et al., 1989). This leaves
to SNe Ia the need to produce the 1/2 to 2/3 of the
iron-peak
nuclei
from
titanium
to
zinc
(Ti–V–Cr–Mn–Fe–Co–Ni–Cu–Zn).
Calculations of explosive nucleosynthesis associated
with carbon deflagration models for Type Ia events
(Domı´nguez et al., 2001; Iwamoto et al., 1999; Maeda et al.,
2010; Thielemann et al., 1986) predict that sufficient
iron-peak nuclei ( 0.6 0.8M of 56Fe in the form of 56Ni) are
synthe sized to explain both the powering of the light curves
due to the decays of 56Ni and 56Co and the observed mass
fraction of 56Fe in galactic matter. Estimates of the timescale
for first entry of the ejecta of SNe Ia into the interstellar
medium of our galaxy yield 2 109 years, at a metallicity
[Fe/H] 1 (Goswami and Prantzos, 2000; Kobayashi and
Nomoto, 2009; Kobayashi et al., 2000).
A representative nucleosynthesis calculation for such a
Type Ia supernova event is shown in Figure 6. Note
particularly the region in mass fraction between 0.2 and
0.8M , which is dominated by the presence of 56Ni in nuclear
statistical equi librium. It is this nickel mass that is
responsible – as a conse quence of the decay of 56Ni
through 56Co to 56Fe – for the bulk of the luminosity of Type
Ia supernovae at maximum light. This is a critical factor in
making SNe Ia the tool of choice for the determination of the
cosmological distance scale – as they are one of the best
calibrated bright stellar objects known. From the point of
view of nucleosynthesis, we then understand that it is the
approximately 0.6M of iron-peak nuclei ejected per event by
Type Ia supernovae that represents the bulk of 56Fe
in galactic matter. The remaining mass that is converted into
nuclei from 16O to 40Ca does not make a significant contribu
tion to the synthesis of these elements.
Chemical Evolution
We have concentrated in this review on three broad
categories of stellar and supernova nucleosynthesis sites:
(1) the mass range 1M ≲M≲10M of ‘intermediate’-mass stars,
for which sub stantial element production occurs during the
AGB phase of their evolution; (2) the mass range M≳10M ,
corresponding to the massive star progenitors of core
collapse supernovae; and (3) Type Ia supernovae, which are
understood to arise as a conse quence of the evolution of
IMS (M≲6 8M ) in close binary systems.
In the context of models of galactic chemical evolution, it
is extremely important to know the production timescales for
each of these sites – that is, the effective timescales for the
return of a star’s nucleosynthesis yields to the interstellar
gas. The lifetime of a 10M star is approximately 5 107 years.
We can thus expect all massive stars M≳10M to evolve on
time scales tcore collapse<108 years and to represent the first
sources of heavy element enrichment of stellar populations
(Heger and Woosley, 2010). In contrast, IMS evolve on
timescales tΙΜS≳108 109 years. Finally, the timescale for SNe
Ia product enrichment is a complicated function of the binary
history of Type Ia progenitors (see, e.g., Livio, 2000).
Observations and theory suggest a timescale tSNIa≳1.5 2 109
years.
Many significant features of the evolution of our galaxy
follow from a knowledge of the nuclear ashes and
evolutionary timescales for the three sites we have
surveyed. The primordial composition of the galaxy was that
which it inherited from cosmological nucleosynthesis –
primarily hydrogen and helium. The characteristics of the
first stellar contributions to nucleosynthesis, whether
associated with Population II or a Population III, reflect (as
might be expected) the nucleosyn thesis products of the
evolution of massive stars (M≳10M )
2.1.6 Nucleosynthesis and Galactic
n
o
it
c
a
r
f
s
s
a
M
100 10−1 10−2 10−3
W7
58
Ni 28Si
56
Fe
12
C
54
Fe
53
Ni
0
0.2 0.4 0.6 0.8 1 1.2 1.4 Interior mass (M )
of short
Spectroscopic
lifetimes (t≲108 abundance studies of the
oldest stars in our
years). The
galaxy, down to
significant
metallicities [Fe/H] 4... 3,
trends in
have been most recently
galactic
reviewed in the context
chemical
of galactic chemical
evolution of
concern here evolution by Kobayashi
et al. (2011). These
are those
involving the studies typically reveal
timing of first abundance trends that
can best be understood
entry of the
products of the as reflect ing the
nucleosynthesis products
other two broad
of the massive stars and
classifications of
asso ciated core collapse
nucleosynthesis
contributors that we have supernovae that can be
identified: low-mass stars expected to evolve and
(s-process) and Type Ia to enrich the interstellar
supernovae (iron-peak media of galaxies on
rapid time
nuclei).
scales (≲108 years). Metal-deficient stars ( 1.5≲[Fe/H]≲ 3) in
our own galaxy’s halo show two significant variations with
Figure 6 The composition of the core of a Type Ia supernova
as a function of interior mass. Note the region of
approximately 0.6 M within which the dominant product is
56
Ni. Reproduced from Timmes FX, Brown EF, and Truran
JW (2003) On variations in the peak luminosity of Type Ia
supernovae. Astrophysical Journal 590: L83–L86.
respect to solar abundances: the elements in the mass
range from oxygen to calcium are overabundant –
relative to iron peak nuclei – by a factor 2 3, and the
abundance pattern in the heavy element region A≳60 70
closely reflects the r-process abundance distribution that
is characteristic of solar
Origin of the Elements 11937
1.4
MP
F/
a
−0.4 −0.6 −0.8
C[
1.2 1.0 0.8 0.6 EMP,RGB EMP,MS
0.4 0.2 0.0 −0.2
]e
Crich,RGB Crich,MS
CEMP,RGB CEMP,MS
−0.228
−6.0 −5.0 −4.0 −3.0 −2.0 −1.0 0.0 1.0
[Fe/H]
Figure 7 Observed evolution of the calcium-to-iron abundance ratio with metallicity. Data collected from the literature using the SAGA
database (Suda et al, 2008, 2011).
2.00
−14.00 −16.00
−18.00
ε
g
o
l
e
v
it
a
le
R
0.00
−2.00
−4.00
−6.00
−8.00
−10.00 −12.00
30 40 50 60 70 80 90
Atomic number
Figure 8 The heavy element abundance patterns in galactic halo stars with the solar system r-only abundance distribution. The solid
blue lines are the scaled r-process-only solar system elemental abundance curves, normalized to the Eu abundance of each star.
Stellar data (from the top) are filled (red) circles, CS 22892–052 (Sneden et al., 2003, 2009); filled (green) squares, HD 115444
(Sneden et al., 2009; Westin et al., 2003); filled (purple) diamonds, BDþ17 3248 (Cowan et al., 2002; Roederer et al., 2010); (black)
stars, CS 31082–001 (Hill et al., 2002; Sneden et al., 2009); solid (turquoise) left
pointing triangles, HD 221170 (Ivans et al., 2006; Sneden et al., 2009); solid (orange) right-pointing triangles, HE 1523–0901 (Frebel
et al., 2007); (green) crosses, CS 22953–003 (Franc¸ois et al., 2007); open (maroon) squares, HE 2327–5642 (Mashonkina et al.,
2010); open (brown) circles, CS 29491–069 (Hayek et al., 2009); and open (magenta) triangles, HE 1219–0312 (Hayek et al., 2009).
All abundances, except those of CS 22892–052, have been vertically displaced downward for display purposes. The solid lines are
solar system r-process-only predictions from Simmerer et al., 2004. Reproduced from Cowan JJ, Roederer IU, Sneden C, and Lawler
JE (2011) r-Process abundance signatures in metal-poor halo stars. In: McWilliam A (ed.) RR Lyrae Stars, Metal-Poor Stars, and the
Galaxy, Carnegie Observatories Astrophysics Series, Vol. 5, p. 223.
system matter, with no evidence for an s-process
nucleosynthe sis contribution. For purposes of illustration,
the trends in [Ca/ Fe] as a function of metallicity [Fe/H] are
shown in Figure 7. Note the factor 2 3 overabundance of
calcium with respect to iron at metallicities below [Fe/H] 1.
We also display in Figure 8 the detailed agreement of
metal-poor star heavy
element abundance pattern with that of solar system
r-process abundances for ten metal-poor but r-process-rich
stars CS 22892–052 (Sneden et al., 2003, 2009), HD
115444 (Sneden et al., 2009; Westin et al., 2003), BDþ17
3248 (Cowan et al., 2002; Roederer et al., 2010), CS
31082–001 (Hill et al., 2002; Sneden et al., 2009), HD
221170 (Ivans et al., 2006; Sneden
12 Origin of the Elements
1.0
0.5
0.0
−0.5
]
u
305
0.600
MP
EMP,RGB
EMP,MS
E/
a
−3.0
B[
−1.0 −1.5
−2.0 −2.5
−4.5 −4.0 −3.5 −3.0 −2.5 −2.0 −1.5 −1.0 −0.5 0.0 0.5
[Fe/H]
Figure 9 The history of the [Ba/Eu] ratio is shown as a function of metallicity [Fe/H]. This ratio reflects to a good approximation the
ratio of s-process to r-process elemental abundances and thus measures the histories of the contributions from these two
nucleosynthesis processes to galactic matter. Data collected from the literature using the SAGA database (Suda et al., 2008, 2011).
a function of [Fe/H] for a large sample of halo and disk stars.
Note that at the lowest metallicities, the Ba/Eu ratio clusters
around the ‘pure’ r-process value ([Ba/Eu]r-process 0.9); at a
metallicity [Fe/H] 2.5, the Ba/Eu ratio shows a gradual
et al., 2009), HE 1523–0901 (Frebel et al., 2007), CS increase. In the context of our earlier review of nucleosyn
22953–003 (Franc¸ois et al., 2007), HE 2327–5642 thesis sites, this provides observational evidence for the first
(Mashonkina et al., 2010), CS 29491–069 (Hayek et al., input from AGB stars, on timescales perhaps approaching
2009), and HE 1219–0312 (Hayek et al., 2009). Both of tIMS≳109 years.
these features are entirely consis tent with nucleosynthesis
In contrast, evidence for entry of the iron-rich ejecta of
expectations for massive stars and associated core collapse SNe Ia is seen first to appear at a metallicity [Fe/H] 1.5...
supernovae.
1.0. This may be seen reflected in the abundance histories
The signatures of an increasing s-process contamination of [Ca/ Fe] with [Fe/H] in Figure 7. This implies input from
first appears at an [Fe/H] 2.5... 2.0. The evidence for this is Type Ia supernovae on timescales tSNIa≳1 2 109 years. It may
provided by an increase in the ratio of barium (the be of interest that this seems to appear at approximately the
abundance of which in galactic matter is dominated by transi tion from halo to (thick) disk stars.
s-process contributions) to europium (almost exclusively an
These observed abundance histories confirm theoretical
r-process product). The ratio [Ba/Eu] is shown in Figure 9 as
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