Sec on 12.
Origin of the elements in the Universe
12.1 Light elements
12.2 Heavy elementstiti

12.3 Chemical evolu on of the Universe

Cosmic abundances (atomic number)

1012

He

10

C

10

10

8

10

6

10

4

10

2

10

0

10-2

O

Fe

0 5 10 15 20 25 30 35 40 45 50 55 60 65 70 75 80 85 90 95
Atomic number Z*Mass ra oti

12Mass fraction (H=10 )H

Cosmic abundances (mass number)
10

0

-1

H He

-2

O

10
10

C

-3

Mass
frac on
質量割合

10

Fe

-4

10

-5

10

s-process
sプロセス

10-7

Sr

10-6

r-processrプロセス

r-processrプロセス
s-processsプロセス
s-process
sプロセス
Ba

-8

10

-9

Au Pb
Pt

10

10

-10

ti

10-11

0

50

100
150
Mass
number
質量数

200

250

tt

h p://www.caltech.edu/news/caltech-led-teams-strike-cosmic-gold-80074

Sec on 12.
Origin of the elements in the Universe
12.1 Light elements
12.2 Heavy elementstiti

12.3 Chemical evolu on of the Universe

Bigbang nucleosynthesis10

η = 6.23x10-10 Nν= 3.0 H0= 70.50 km/s/Mpc Tcmb= 2.725 K

1

-1

10

p

p

n

he4
n

he4

10

d

-3

d

t

Mass Fraction

10

-5

t

he3

he3

10

t

-7

be7

li7

-9

10

n
d

t
li7

-11

10

he4

-13

10

li7 li6
he3

-15

10

1

10

be7

li6

b11

2

10

3

10
Time (s)

4

10

5

10

tt

h p://cococubed.asu.edu/code_pages/net_bigbang.shtml

of し，宇宙の膨張率，及び，それを通しての弱い相互作用の反応率には重力相互
neutrinos with CMB photons andion rate (as measured by the Hubble
重要な役割を果たしている．
couplefrom the CMB photons and~ 2 sec~200 sec ~ 300 sec ~ 900 sec Timeantineutrinos) continue to interact
Fermi 分布を保他の粒子との相互作用が切れたニュートリノは，それ以降もent weakinteractions until the Unihas dropped below ∼0.8 MeV. This宙の膨張とともに温度が低下していく．現在の温度≈
1.9
Kは宇宙背景放射のnot “freeze out” (see Figure 1).
n/p ~ 1/6
n/p ~ 1/7
Nucleo±s, protons,
e
,
and
ν
(
ν̄
)
continueLife me of neutronee
り低い．これは，ニュートリノの脱結合以降に起こる電子と陽電子の対消滅にsynthesis
beit not rapidly enough to allow the(n decay)
トンが多少加熱されるからである．
77
riumvalue of n/p = exp(−#m/T ),= mn − mp = 1.29 MeV. As a result,t > Neutron
1s：中性子の崩壊
Abundance.annualreviews.org • Primordial
Nucleosynthesis
Breakdownof 465(12.52) が右向きに進
中性子の数密度を変化させるのは，中性子のベータ崩壊（
neutron-proton Equilibrium
ratio. Neutrons and protons can be converted
intoThe
rst reac on2= H ）の生成
on:一方，重陽子（(T ~ 0.7D
MeV)
p + e− ↔ n + νe+

p + ν̄e ↔ n + e

n ↔ p + e− + ν̄.

p + n ←→ D + γ
(294)

ti

ti

fi

ft

- A er photon energy decreases
は中性子を原子核の中に取り込むことによって中性子数を保存する．左から右~ 0.1
MeV)we use
he熱反応であり，放出されるエネルギーは重陽子の結合エネルギー
nucleons are in a non-relativistic regime (E = (Tm+p2 /2m),B(D) = 2.
. (276):- Before neutron decay
る．温度が十分低下するまで，上の反応は平衡状態にあり，生成された重陽子
!
"3/2
mp T
%
&3/2
−mp /Teてしまう．右向きの反応が進行するのは，温度が0.1 MeV まで低下して* Binding energy ofkT
D ~≈2 MeV
2π
mp
(mn −mp )/T
Q/T
=
e
≈e
,
(295)
# mn T $3/2
mn
−mn /T
100 秒後である．中性子の平均寿命は τn = 882 s であ
これは宇宙誕生から約e2πAll neutrons go to 4He(n/p ~ 1/7)
(nn /2)(2mp + 2mn )Y =⇠ 0.25np mp + nn mn

Consistent with
Cosmic abundance
h p://hyperphysics.phy-astr.gsu.edu/hbase/Astro/bbnuc.html

No stable nuclei with mass number of 5 and 8titt

Next reac on will be 4He x 3 inside of stars（Not possible in bigbang due to low density）

inoue@tap.scphys.kyoto-u.ac.jp
ese eight stars are determined for the
first
time
by
this
study.
Effective
10 [A/B] = log(N /N ) − log(N /N ) , and
A
B
A
B %
alysis
of
Hα
and
Hβ.
While
seven
stars
have
Li
abundances
as
high
Li
problemelementsA
andB.
Lithiumabundancesarecon
our objects with [Fe/H] < −3 have A(Li) = log(Li/H) + 12 ! 2.0,
instead of log
#Li .
No.
2,
2009
LITHIUMl-poor (EMP) turnoff stars having low Li abundances, asABUNDANCES
reported by OF EX
ances for stars with [Fe/H] < −3 is lower by 0.2 dex than that of theLiabundance
raint on the metallicity dependence or scatter of the Li abundances iss with [Fe/H] < −3. Correlations of theBBN
Li abundance with effectivend Sr, and with the kinematical properties are investigated, but no clear

ti

ti

ti

Destruc on inside
sis, abundances – stars: abundances – stars: Population II plateau??
of stars＋
the existence of poorly understood processes that deplete Li in
Produc onbymetal-poor
turnoff stars (e.g., Korn et al. 2006, 2007), astration
Cosmic raynces haveof Li in the gas that formed these stars (perhaps by massive
spalla once turnoff
zero-metallicity progenitors; Piau et al. 2006), problems inre of bigthe measured abundances of Li in metal-poor stars, systematic
gation, by
uncertainties in the standard BBN model predictions, or exotic
Aoki+09
oor mainprocesses in the early universe arising from nonstandard particle
Figure 3. Li abundances as a function of [Fe/H]. The results of the preThis sophysics (Cyburt
et al. by
2008
references
work are plotted
filledandcircleswith errortherein).
bars. The open triangles
synthesis
Recentmeasurements
of by
very
metal-poor
stars
also
squares
indicate the results
Asplundet al. (2006)turnoff
andBonifacio
et al.
(20
Metallicity

d 6500 K, log g values of 4.0 and 4.5, and a range inalues: +0.10, 0.00, !0.10, !0.20, !0.30, !0.50, !1.00,
2.00, !2.50, and !3.00. We have tested the dependencyof the microturbulence parameter, !, by running mod.0, 1.5, and 1.25 km s!1. Here is a brief summary: an"80 K in TeA gives an error in A( Be) of 0.01 (except
o change from 6000 to 6250 K at log g of 4.5); an error
in [Fe/ H ] gives an error of 0.03–0.04 in A(Be) (except
e/ H ], where a change of 0.5 in [Fe/ H ] gives an error
and an uncertainty in log g of "0.2 produces errors in
m 0.07 (at the high temperature) to 0.10 (at the lower
ures in our models); a change in 0.25 km s!1 in ! is"0.00–0.02 in A( Be). Thus the typical errors in A( Be)
ars are "T # 0:01, "Fe # 0:02, "! # 0:01, but "log g #e accuracy of the spectrum synthesis fits is estimatedlarge a change in A(Be) can still represent the observedwith adjustments in the continuum (which is partly
ed by the S/N of the spectrum) and Gaussian smoothing
thetic spectra. This is typically 0.05–0.06 dex, but it isn G64-37, which has a S/N of only 87. The final errors
quadrature) for each star are given in Table 3.
tion, there may be uncertainties resulting from effects
e-dimensional versus three-dimensional models, as disPrimas et al. (2000a). According to the work of Garcı́aal. (1995) and Kiselman & Carlsson (1996), non-LTE
rays (<=
e foundCosmic
to be insignificant
for SN)Be.

HD 94028 is offset a small amount from the circle below it fromthe Boesgaard et al. (1999a) value. Our result is a confirmationof the high Be abundance found earlier. It also indicates thatthe spread in A(Be) at this metallicity (½Fe/H' # !1:5) is real;
Boesgaard et al. (1999a) remarked in their discussion of this starthat its high Be abundance might indicate an intrinsic dispersion in Be at a given Fe. Two other stars, HD 132475 and theCosmic ray spalla on (Li, Be, B)

Cosmic ray(p, alpha)+ targets (C, N, O)=> Li, Be, B

C, N, O (<= past nucleosynthesis)=> secondary
process(slopeto =those2)abundances
determined here
can be compared

Slope = 1
(primary?)

4. RESULTS AND DISCUSSION

ti

rature. For HR 8888 we find A(Be) ¼ þ1:38 " 0:10, inagreement with 1:44 " 0:13 in Boesgaard et al. (2001)
ada-France-Hawaii Telescope Gecko data. For HD 94028
e a value for A( Be) of +0:51 " 0:12, agreeing with
08 determined by Boesgaard et al. (1999a) from Keckolution Echelle Spectrometer (HIRES) data.

Fig. 7.—Plot of A( Be) vs. [ Fe/ H ] for this sample compared to the literature
Boesgaard+06
sample. Our results are the filled squares superposed on this diagram.
Circlesand triangles (upper limits) are from Boesgaard et al. (1999a) and Boesgaard(2000, 2004). Open squares ( low-metallicity stars) are from Primas et al. (2000a,

Sec on 12.
Origin of the elements in the Universe
12.1 Light elements
12.2 Heavy elementstiti

12.3 Chemical evolu on of the Universe

Stellar life

(C: Essay Web)

PhaseMain reac ons

Products

T

Nuclear sta s cal equilibrium
ti

ti

ti

元素はいかにつくられたか（岩波書店）

Element ejec on from stars
1-6 Msun: AGB mass loss (Karakas 2010, MNRAS, 403, 1413)> 10 Msun: supernovae (Kobayashi et al. 2006, ApJ, 653, 1145)

10

0.01

0.001

ti

Ejected Mass (IMF weighted)

C12
O16
Mg24
1
Si28
S32
0.1 Fe56-2

Rela ve importance(weighted with IMF)
10-3

10-4

10-5
0

5

10 15 20 25 30
Initial Mass (Msun)ti

Ejected Mass (Msun)

10

35

40

0

5

10 15 20 25 30
Initial Mass (Msun)

35

40

Dust produc on in the Universe

2011) 19:43

Page 35 of 67
Page 36 of 67

AGB
d for
rve),
!low (m)andas welln curve).
observed
SNe at
The darko the
12 M! )
noximateP SNe

AGB

Fig. 7 Specific dust
productivity ξd (m) of stellar
masses calculated for "max (m)(dark blue curve), "high (m)(violet curve), "low (m) (cyanfit (m) and " fit (m)curve), "maxlow(black curves) as well as for
AGB stars (green curve). The
dark grey zone corresponds tothe critical mass range(8–12 M! ) and the light grey
region corresponds to theapproximate mass range forType II-P SNe. The solid,
dotted, dashed, dashed-dotted,
and dashed-dot-dotted curves
represent the Salpeter,mass-heavy, top-heavy, Larson 1
and Larson 2 IMFs, respectively

SN

Astron Astrophys Rev (2011) 19:43

IMF weighted

AGB
SN

Gall et al. 2011
Probably
dominated
by
AGB
stars
s. We also plot the highest observationally derivedThe
dusttotalyielddust
for productivity µD of AGB stars and SNe, subdivided into 3 massow (m) are also in good agreement with the observed dust yields

ranges,is presented
nt (see Table 4 and discussion in Sect. 5.3.3). The
inferred
dust in Fig. 8. For a ‘low’ SN efficiency, "low (m), the total amount of
dust producedis almost exclusively manufactured by AGB stars. Dust production bypler remnant may not be representative of dust fromCCSNe (theti（But need SN in the early Universe)

The origin of elements
Low-massStar

Big bang
Core-collapse SNe
Thermonuclear SNe

Nh

Mc

Ts Og

Neutron-capture nucleosynthesiss (slow)-processr (rapid)-process
n
n

n

n
n n

n
n

Decay

Decay

p

p

p
p

p

p

Ba, Pb, …

Au, Pt, U, …

Inside of stars

SN? NS merger?

p p

s-process and r-process

Mass number
A=Z+N

s-processr-process(C) Nobuya Nishimura

Cosmic abundances
10

0

-1

H He

-2

O

10
10

C

-3

Mass
frac on
質量割合

10

Fe

“2nd peak”

-4

10

-5

10

s-process
sプロセス

10-7

Sr

10-6

r-processrプロセス

r-processrプロセス
s-processsプロセス
s-process
sプロセス
Ba

-8

10

“3rd peak”

-9

Au Pb
Pt

10

10

-10

ti

10-11

0

50

100
150
Mass
number
質量数

200

250

94

energy released in the whole chain (6.44). The abov
7 金はどこでつくられたか

which eachrequirep-pneutron
reaction to complete. For
Seedreac oneon of
therefore in that case !nuc = 21 QH rpp /ρ. Expressing rp
13 abundance
the hydrogencan compute the7 （ener
C+4 He →X,16we
O+n
T > 8 x 10 K 7
combination of pp chains as
22
25
Ne+42He→Mg+n（7ρ
#συ$pp ,!pp = ψ qH X
mu
781
1. Shellburning
は約
8×10
K
以上，
（
7.6
）
は約
2.5×10
K 以上で
式（ 7.5 ）
where qH = QH /4mu is the total energy release
per g
(at
the
bo
oms
of
He
H
layers)tween 1 (for the pp1 chain) and 2 (for the pp2 and pp3
る．したがって，中性子捕獲反応プロセスは，恒星の進化
2. HeBothis enrichedthe chains.
ψ and qH therefore depend on the te
けるヘリウム燃焼過程で起きていると考えられる（第 1 章参=> losses.
Shell The
ashoverall temperature depende
ent neutrino
これらのヘリウム捕獲反応により得られる中性子密度は#συ$pp3.
andis shownon
in Fig. 6.5. The pp chain is the le1
Conveccycles withpower-lawν (eq. 6.30)
varyin1011 cm 3 程度であり，このときの核種
（ 鉄など
）
による中
=>amixing
in exponent
the envelope

s-process in AGB stars2+ H is mixed to the He layer
捕獲の寿命は数年〜数千年程度と非常に長い．その結果，
The CNO cycle
4. 12C + H => 13N =>slow
13Cの s をとって「子捕獲反応はゆっくりと進行するので，
If some C,13C
N, andO =>is alreadyin the gas ou+ He
16O +presentnロセス」と呼ばれる．
is sufficiently high, hydrogen fusion can take place v=>
s-processsequence of reactions that typically starts witha proto
第 1 章で述べたように，ヘリウム燃焼過程には2 つの種!
ある．1 つは，恒星の中心で水素が燃え尽きた後に起きる
412 C + 1 H → 13 N + γ
ウムコア燃焼である．10 M 以上の大質量星では，ヘリウ
13 N → 13 C + e+ + ν
ア燃焼の最終段階で中心温度が約 3×108 K に達し ，式（7
元素はいかにつくられたか（岩波書店）
13 C + 1 H → 14 N + γ
tt

fl

ti

ti

3

The CNO cycle

CNO
cycle
If some C, N, and O is already present in the gas out of which a star forms, and if the temperatu

is sufficiently high, hydrogen fusion can take place via the so-called CNO cycle. This is a cyclicsequence of reactions that typically starts with a proton capture by a 12 C nucleus, as follows:!
12 C + 1 H → 13 N + γ

13 N → 13 C + e+ + ν

13 C + 1 H → 14 N + γ" 14

N + 1 H → 15 O + γ

13C should be produced
under H-poor condi on
for s-process

15 O → 15 N + e+ + ν

15 N + 1 H → 12 C + 4 He

→ 16 O + γ

16 O + 1 H → 17 F + γ

17 F → 17 O + e+ + ν

17 O + 1 H → 14 N + 4 He

(6.4

For example, if we denote by rpD the rate of 2 H + 1 H, one has rpD = rpp , etc. Note that describing the p-p reaction
Textbook
by Pols2
1‘slow’ and the H + H as ‘fast’ refers to the difference in cross-section factors #συ$ and not to the number of reactions pti

1

H is provided in the He-burning layer(unique in AGB stars)

Observa onal evidenceFirst evidence
Tc (Z = 43, no stable ist)
(Merrill 1952)

post-AGBti

Lugaro+16

13

s-process in massive stars(weak s-process)22

4
C+SeedHereac→ 16onO+n

Ne+4 He → 25 Mg+n

（ 7.5

（
7.6
8K
T > 2.5 x 10

は約 8×107 K 以上，（ 7.6 ）
は約 2.5×108 K 以上で起式
（ 7.5 ）

る．したがって，中性子捕獲反応プロセスは，恒星の進化にHe burning coreNo. 2, 2010

けるヘリウム燃焼過程で起きていると考えられる（第 1 章参照
14 N(α,γ) 18F(β+ν) 18O (α,γ) 22Ne
これらのヘリウム捕獲反応により得られる中性子密度は 10

10

11

Cu

Ga

WEAK s-PROCESS程度であり，このときの核種（ 鉄など ）による中性of the shell evolution,
捕獲の寿命は数年〜数千年程度と非常に長い．その結果，中growing because of the
1000
Fe
Ni
Zn
Ge
Se
Kr
Sr Zr
Mo
12
+
C reaction (λ12 C12 C ∝Z = 26
28
30
32
34
36
38
子捕獲反応はゆっくりと進行するので，
slowShortlyの s をとって「before the SN
O16x2
O-burning layers, and te100
ロセス」と呼ばれる．
stellar material (Arnett
O16/2
calculations, Limongi e
第 1 章で述べたように，ヘリウム燃焼過程には
2 つの種類10
in the last convective Cfrom T9 = 1.07 to T9 = 1
ある．1 つは，恒星の中心で水素が燃え尽きた後に起きるヘet al. (2000), the C shell l1
is a few times 1012 cm−
ウムコア燃焼である．10 M 以上の大質量星では，ヘリウム
et al. (2004), but these ain theC shell，式
also
th
0.1
ア燃焼の最終段階で中心温度が約 3×108 Kに達し（in7.6
process yields are not af
55 反応が起こりうる．もう
60
65
70
75
80
85 190つは，恒星の中心でヘリウムが燃
95 100convective shell C-burni
Mass Number
SNPignatari+10
explosion. For insta
尽きた後に形成されるコアの周りで起きるヘリウム殻
prohibited by stellar stru
Figure 2. Overabundances
between 57 Fe and 100 Mo atC+O
core He exhaustionfor
Xi /Xsun

Co

ti

cm

3

As

Br

Rb

Y

Nb

tt

h p://www.caltech.edu/news/caltech-led-teams-strike-cosmic-gold-80074

Origin of r-process elements?
Some phenomena related to neutron starSupernova(C) M. Weiss~ 1 event per 100 yr in a galaxy(R ~ 10-2 yr-1)

Neutron star merger(C) NASA~ 1 event per 10,000 yr in a galaxy(R ~ 10-4 yr-1)

Sec on 12.
Origin of the elements in the Universe
12.1 Light elements
12.2 Heavy elementstiti

12.3 Chemical evolu on of the Universe

Stars keep informa on
about nucleosynthesis in the past
“Galac c archeology”

Our Galaxy

ti

ti

tt

h p://astronomy.swin.edu.au/cms/astro/cosmos/T/Thick+Disk

Cosmic abundance
10

0

-1

H He

-2

O

10
10

C

-3

10

X(Fe) ~ 10-3

Fe

-4

質量割合

10

-5

10

s-process

10-7

Sr

sプロセス

10-6

r-processrプロセス
s-processsプロセス
s-process
sプロセス

Ba

-8

10

Mass frac on

r-process
rプロセス

-9

Au Pb
Pt

10

10

-10

ti

10-11

0

50

100
150
200
質量数 Mass number

250

-11

80

100

120

140

160

180

200

Cumulative mass abundance in the sun

3rd peakLanthanide

2nd peak

1st peakSolar r-process abundance

10-10

r-process production rate, sGRBs,
and GW170817
10

220

10-10
10-11

80

100

120

140

160

180

Atomic mass number

200

220

10-7

3rd peakLanthanide

2nd peakBefore discussing details, we describe here the r-p
solar
abundance pattern of r-process elements taken
10-8
three peaks. For the solar abundance pattern, mostaround the first peak. However, the abundance ratio
10-9 poor stars, of which the abundance pattern lmetal
often di↵erent from that of the solar pattern. Some ofthe -10second peak (heavy r-process) that are similar to10
similar amounts
first160
peak
80
100 of120the140180elements
200 220as compar
Atomicpattern (e.g. Sneden
etmassal. number
2008 and references thecontain a substantial amount of the first peak elemen
1st peak

10-9

Cumula ve
mass frac on(from the heavier side)

ti

10-8

Cumulative mass abundance in the sun

-7

ti

10

10-6

Figure 1: The solar abundance pattern of r-process(right). The solar r-process abundance pattern is tak

3rd peak

10-6

2nd peak

10-5Lanthanide

Atomic mass number

1st peakSolar r-process abundance

Our goal in this article is twofold. First we summarize the cumulative evidences supporting
hat r-process nucleosynthesis takes place in rare events in which a significant amount of r-process-5
10ements are produced in each event. Using the rate and mass ejection per event inferred fromhese R-process
measurements, elements
we test the neutron star merger scenario
for the origin of r-process elements in
10-6
he cosmos. This evidence clearly rules out the normal cc-SNe scenario. Moreover, the rate agrees
ith merger estimates from galactic binary neutron stars,
10-7 from sGRBs, and from GW170817. At
he same time the amount of matter is consistent with the kilonova/macronova, AT2017gfo, and-8-7
he candidates
associatedwith cosmological sGRBs. 10Second,
we turn to the Galactic chemicalX(r) ~ 10
volution of r-process elements at later times [Fe/H] & -9 1 and discuss whether the neutron star
10
(Ascenario> 90)can consistently explain the observed distribution
mergerof [Eu/Fe].

[ Mg/F

nature and extent of n-capture nucleosynthesis in the early Galaxy. Abundance signatures of these
0.5
elements can help to identify the parameters of the first stars—long since gone—and the sites forr- and s-process nucleosynthesis (Cowan & Sneden 2006). There are many ways to approach this0
issue. We
begin in Figure 14 by contrasting the abundance behavior as a function of metallicityof the α element Mg and the n-capture element Eu. Only large stellar samples are employed forthese –0.5
correlations; the literature sources are listed in the figure caption.

Abundance ra o in Galac c stars (Mg/Fe)ab[ Mg/Fe][Eu/Fe]

1.5
1.5

Halo

Disk

1.0
1.0
0.5
0.5
00
–0.5
–0.5

b

1.5

–3

–2

–1

0

[Fe/H]

Figure 14
Sneden+08
1.0[Mg/Fe] and [Eu/Fe] abundances as a function of [Fe/H] metallicity for halo and disk stars. For this figure
the data have been taken only from large-sample surveys: Fulbright (2000); Reddy et al. (2003); Cayrel et al.
(2004);0.5
Cohen et al. (2004); Simmerer et al. (2004); Barklem et al. (2005, red points); Reddy, Lambert &Allende Prieto (2006); François et al. (2007). In both panels the dotted lines represent the Solar abundance
ratios. In panel b, the solid red line is a least-square fit to the Eu data, and the two dashed black lines indicate0
the approximate extent of the Eu/Fe data (similar to Cowan & Thielemann 2004).
ti

ti[Eu/Fe]

Time

Time

Time

Time

Time

Time

[Fe/H][Fe/H]

[Mg/Fe]
Core-collapse
SNe
Type Ia(thermonuclear)
SN[Fe/H]

[ Mg/F

nature and extent of n-capture nucleosynthesis in the early Galaxy. Abundance signatures of these
0.5
elements can help to identify the parameters of the first stars—long since gone—and the sites forr- and s-process nucleosynthesis (Cowan & Sneden 2006). There are many ways to approach this0
issue. We
begin in Figure 14 by contrasting the abundance behavior as a function of metallicityof the α element Mg and the n-capture element Eu. Only large stellar samples are employed forthese –0.5
correlations; the literature sources are listed in the figure caption.

Abundance ra o in Galac c stars (Mg/Fe)ab[ Mg/Fe][Eu/Fe]

1.5
1.5

Halo

Disk

1.0
1.0
0.5
0.5
00
–0.5
–0.5

b

1.5

–3

–2

–1

0

[Fe/H]

Figure 14
Sneden+08
1.0[Mg/Fe] and [Eu/Fe] abundances as a function of [Fe/H] metallicity for halo and disk stars. For this figure
the data have been taken only from large-sample surveys: Fulbright (2000); Reddy et al. (2003); Cayrel et al.
(2004);0.5
Cohen et al. (2004); Simmerer et al. (2004); Barklem et al. (2005, red points); Reddy, Lambert &Allende Prieto (2006); François et al. (2007). In both panels the dotted lines represent the Solar abundance
ratios. In panel b, the solid red line is a least-square fit to the Eu data, and the two dashed black lines indicate0
the approximate extent of the Eu/Fe data (similar to Cowan & Thielemann 2004).

[Eu/Fe]

Time

ti

ti

ti

Longer delay me for Type Ia SNeOur understanding about thenucleosynthesis is correct??
(A) Total amount(B) Time scale

ANRV385-AA47-10

ARI

22 July 2009

4:3

Role as a “clock” in galaxy forma on
Sculptor

Fornax

Sagittarius

Carina

MW

1.0

a

0

–0.5

1.0

b

0.5

[Ca/Fe]

0

–0.5

–2

–1

0

[Fe/H]
Figure 11
α-elements, (a) Mg and (b) Ca, in four nearby dwarf spheroidal galaxies: Sgr (orange: McWilliam &Smecker-Hane 2005, Monaco et al. 2005, Sbordone et al. 2007), Fnx (blue: Shetrone et al. 2003, Letarte2007), Scl ( green: V. Hill & DART, in preparation, Shetrone et al. 2003, Geisler et al. 2005), and Carina( purple: Shetrone et al. 2003, Koch et al. 2008a). Open symbols refer to single-slit spectroscopy
measurements, whereas filled circles refer to multiobject spectroscopy. A representative error-bar for thelatter is shown on the left-hand side of the picture. The small gray squares are a compilation of the Milky

Fe in dwarf galaxies were smallerwhen Type Ia SNe begun to operatetis. 2009.47:371-425. Downloaded from www.annualreviews.org
ai Libr. - National Observatory of Japan on 07/24/16. For personal use only.

[Mg/Fe]0.5

Tolstoy 08

Annu. Rev. Astro. Astrophys. 2008.46:241-288. Downloaded from www.annualr
Access provided by National Astronomical Observatory of Japan on 12/15/14. For pelements can help to identify the parameters of the first stars—long since gone—and the sites forr- and s-process nucleosynthesis (Cowan & Sneden 2006). There are many ways to approach thisissue. We begin in Figure 14 by contrasting the abundance behavior as a function of metallicityof the α element Mg and the n-capture element Eu. Only large stellar samples are employed forthese correlations; the literature sources are listed in the figure caption.

Abundance ra o in Galac c stars (r-process/Fe)
1.5

a

Halo

Mg (Z = 12)

Disk[ Mg/Fe]

1.0
0.5
0
–0.5

1.5

r-process

Eu (Z = 63, r-process)b

Larger sca er=> Rare event thannormal core-collapse
supernovae[Eu/Fe]

1.0
0.5
0
–0.5
–3

–2

–1

0

[Fe/H]

Sneden+08

Figure 14

ti

ti

tt

[Mg/Fe] and [Eu/Fe] abundances as a function of [Fe/H] metallicity for halo and disk stars. For this figure
the data have been taken only from large-sample surveys: Fulbright (2000); Reddy et al. (2003); Cayrel et al.

High rate
Low ejec on

Low rate
High ejec on

Smaller sca er
in abundance(e.g., Mg)

Larger sca er
in abundance(e.g., Eu)

tt

tt

ti

ti

ti

Mixing mescale ~ 100 Myr

Assignment 5
Read one paper focusing on chemical elements or metallicity
in your research area and summarize the contents in 2 pages.
(ex.) Measure the metal abundances of galaxy to know XXX.
Phenomena XXX is a ected by metallicity because YYY.
An instrument using the property XXX of the element YYY.
In the report, please include following points:(1) Why you choose that paper(2) Central problem of the eld(3) Method and uniqueness(4) Results

fi

ff

ti

(5) Implica ons

レポート課題 5
自分が研究している(興味のある)現象・対象で「元素」や「金属量」に着目している論文を探し、その内容をA4 2ページ程度にまとめよ。
(例) 銀河の元素量を測って、… を知る
金属量が異なると、…の効果で… はこのように影響を受ける
この装置は…という元素の … という性質を使っている
レポートには以下の点を含めること。
(1) なぜその論文を選んだか(2) 当該研究分野の中心的問題(3) 手法とユニークさ(4) 結果(5) 結果から導かれたこと

Summary: Origin of the elements in the Universe• Origin of the elements• Bigbang nucleosynthesis: H, He, Li• Cosmic-ray spalla on: Li, Be, B• Stellar interior: C-Fe

(AGB stars, core-collapse SNe, thermonuclear SNe)• Neutron capture: > Fes-process: AGB stars
r-process: SN? NS merger?

ti

ti

ti

• Test with stars in our Galaxy and dwarf galaxies• Close rela on with galaxy forma on