系外銀河における微細構造輝線の観測とその理解

2014年12月2-3日 @国立天文台

高赤方偏移における星形成銀河の微細構造輝線
観測とその理解 (光赤外からのコメント)
矢部清人 (国立天文台)

系外銀河における微細構造輝線の観測とその理解

2014年12月2-3日 @国立天文台

Cosmic Star-formation history：
Redshift
Hopkins & Beacom 2006

SFR density

z=5

Peak epoch at z~1-3系外銀河における微細構造輝線の観測とその理解

2014年12月2-3日 @国立天文台

Metallicity of galaxies at high redshift：
The Astrophysical Journal Letters, 771:L19 (6pp), 2013 July 10

Zahid et al.

metallicity

Zahid+13, ApJ, 771, L19

stellar mass
Figure 1. MZ relation at five epochs ranging to z ∼ 2.3. The curves are fits to the data defined by Equation (4). The solid curves indicate metallicities determinedusing the KK04 strong-line method and the dashed curves indicate metallicities converted using the formulae of Kewley & Ellison (2008). Data presented in this figure
can be obtained from H.J.Z. upon request.
(A color version of this figure is available in the online journal.)

The mass-metallicity relation and its evolution

系外銀河における微細構造輝線の観測とその理解

2014年12月2-3日 @国立天文台– 9 – redshift：
Metallicity gradient at highpositive gradient

Cresci+2010(220km/s, 13.5G, 0.03)(220km/s, 13.5G, 0.05)(220km/s, 13.5G, 0.07)(220km/s, 3G, 0.03)(220km/s, 3G, 0.05)(220km/s, 3G, 0.07)(150km/s, 3G, 0.03)

Yuan+2011
negative gradient•

Steeper
metallicity
gradient
at Red lines are the measurements for Sp1149 at
Fig. 3.—
Left: Metallicity
vs. galactocentric
radius.
highfrom
redshiftz=1.49this work. The gradient within the central 4.5 kpc is -0.16+/-0.02 dex kpc−1 . Vertical→ Inside-out
red dottedlines show thegrowth?
annulus used to average/sum the spectra. Purple dashed lines show
the typical gradients of local isolated late-type galaxies, using the control sample of Rupke et al.
Positive gradient(2010b). The orange dotted line represents the mean gradient of local early-type galaxies, which→ Infallof pristinegas?
is typically∼ 3 times
shallower thanlocal late-type galaxies (Henry & Worthey 1999). Blue lines•

系外銀河における微細構造輝線の観測とその理解

2014年12月2-3日 @国立天文台

Morphology of galaxies at high redshift：
SXDS/CANDELS

Color composites of galaxies at z~1.4 with HST/ACS+WFC3 images

Large variety of galaxy morphology at z>1

系外銀河における微細構造輝線の観測とその理解

2014年12月2-3日 @国立天文台

Kinematics of galaxies at high redshift：No. 2, 2009SINS SURVEY OF HIGH-REDSHIFT GALAXIES

1403

Forster Schreiber+09
Figure 17. Velocity fields for 30 of the 62 galaxies of the SINS Hα sample. The velocity fields correspond to that derived from the Hα line emission as described in
Section 5.1 (the exception is K20–ID5 for which it was obtained from the [O iii] λ 5007 line instead). The color coding is such that blue to red colors correspond to theblueshifted to redshifted line emission with respect to the systemic velocity. The minimum and maximum relative velocities are labeled for each galaxy (in km s−1 ).
All sources are shown on the same angular scale; the white bars correspond to 1"" , or about 8 kpc at z = 2. The galaxies are approximately sorted from left to rightaccording to whether their kinematics are rotation-dominated or dispersion-dominated, and from top to bottom according to whether they are disk-like or merger-likeas quantified by our kinemetry (Shapiro et al. 2008). Galaxies observed with the aid of adaptive optics (both at the 50 and 125 mas pixel−1 scales) are indicated by theyellow rounded rectangles.

Large variety of galaxy kinematics at z~2系外銀河における微細構造輝線の観測とその理解

2014年12月2-3日 @国立天文台

Connection between galaxies and AGN：
BPT★☆: Stacking analysis

Kewley+01
Yabe et al. 2014

AGN contribution? Different ISM condition?

系外銀河における微細構造輝線の観測とその理解

6
2014年12月2-3日@国立天文台[N II]/Hα and [O III]/Hβ ratios seen at high redshift.
Connection between galaxies
and AGN：
(3)

(4)
AGN

AGN

1.5

0.5
0.0

z~0
HII

-0.5
-1.0

LOG ([OIII]/Hβ)

1.0
0.5

z~0.8

0.0
-0.5

largerq

0.5

LOG ([OIII]/Hβ)

1.0
0.5

z~1.5

0.0

-1.0

highernelarger q-2.0 -1.5 -1.0 -0.5 0.0 0.5
LOG ([NII]/Hα)

-0.5
-1.0
1.0

LOG ([OIII]/Hβ)harder
EUV

0.0
-0.5

-1.0

Kewley et al. 2013

1.0

HII

LOG ([OIII]/Hβ)LOG ([OIII]/Hβ)

1.0

0.5

z~2.5

0.0
-0.5
-1.0
-1.5 -1.0 -0.5 0.0 0.5
LOG ([NII]/Hα)-1.5 -1.0 -0.5 0.0 0.5
LOG ([NII]/Hα)Figure 2. An illustration of the effect of varying different galaxy
parameters on the star-forming galaxy abundance sequence in the[N II]/Hα versus [O III]/Hβ diagnostic diagram. The original SDSS
star-forming galaxy sequence is well-fit by the red theoretical curve.
Raising the hardness of the ionizing radiation field (orange dashedline) moves the abundance sequence towards larger [N II]/Hα and[O III]/Hβ ratios. A similar effect is seen when the electron density of the gas is raised (green dot-dashed line). The relationshipbetween ionization parameter, metallicity and the [N II]/Hα and[O III]/Hβ line ratios is more complex. At high metallicities, raising the ionization parameter causes the [N II]/Hα ratio to become
smaller, while [O III]/Hβ is largely unaffected. At low metallicities,

Different ISM condition
at high redshift?

系外銀河における微細構造輝線の観測とその理解

2014年12月2-3日 @国立天文台

Galaxies at high-redshift is dusty!!：
T. Goto et al.: Infrared luminosity functions with the AKARI

Goto et al. 2010
Fig. 16. Evolution of TIR luminosity density based on TIR LFs (red circles), 8 µm LFs (stars), and 12 µm LFs (filled triangles). The blue open
squares and orange filled squares are for only LIRG and ULIRGs, also based on our LTIR LFs. Overplotted dot-dashed lines are estimates fromthe literature: Le Floc’h et al. (2005), Magnelli et al. (2009), Pérez-González et al. (2005), Caputi et al. (2007), and Babbedge et al. (2006) are incyan, yellow, green, navy, and pink, respectively. The purple dash-dotted line shows UV estimate by Schiminovich et al. (2005). The pink dashedline shows the total estimate of IR (TIR LF) and UV (Schiminovich et al. 2005).

Contribution to SFRD from LIRG/ULIRG increases
with increasing redshift

系外銀河における微細構造輝線の観測とその理解

2014年12月2-3日 @国立天文台

Metallicity of galaxies at high redshift：
The Astrophysical Journal Letters, 771:L19 (6pp), 2013 July 10

Zahid et al.

metallicity

Zahid+13, ApJ, 771, L19

stellar mass
Figure 1. MZ relation at five epochs ranging to z ∼ 2.3. The curves are fits to the data defined by Equation (4). The solid curves indicate metallicities determinedusing the KK04 strong-line method and the dashed curves indicate metallicities converted using the formulae of Kewley & Ellison (2008). Data presented in this figure
can be obtained from H.J.Z. upon request.
(A color version of this figure is available in the online journal.)

Evolution of mass-metallicity relation is real?

系外銀河における微細構造輝線の観測とその理解

high-redshift sample, although without binning them with the rest2014年12月2-3日@国立天文台of the galaxies.

Metallicity of galaxies at
redshift：
2.3 zhigh= 3–4

Abundances of Luminous Infrared Galaxies

A significant sample of 16 galaxies at redshift between 3 and 4 wasobserved by Maiolino et al. (2008) and Mannucci et al. (2009) for the

Mannucci+2010

lower SFRULIRG

Seeing only surface metallicity?
Rupke+2008

metallicity

higherSFR

stellar mass

Left-hand panel:mass–metallicity
of localgalaxies.
g. 11.— Comparison of the mass-metallicity relationFigure
from 1.theFig.the12.—
Differencerelation
betweentheSDSSobserved
a
the
thick
centralline
showingthemedianrelation.
The
coloured
lines
show
the
SS (Tremonti et al. 2004) with LIRG and ULIRG abundancesand ULIRGs and the L − Z and M − Z relavaluesof
SFR.
Right-handpanel: median
metallicity
as a(red)functionof SFRfor galstellar masses. The average LIRG and ULIRG are significantlyof infrared
luminosity.
The
thick
open
d
with
increasing
SFR
atconstant
mass.
er-abundant, as they are when compared to the L − Z relation.
deviations from the M − Z relation. There isdotted lines show 1σ scatter on either side of the mean SDSS
agreement between the two, though comparisotion, which has been shifted upward by 0.1 dex to account forlation yields a slightly larger under-abundance

Low metallicity in dusty and high SFR galaxies?

ratio continues
to increase,
the todecreasing
Charlot& Longhetti(2001). despite
We preferestimate
the ionof these line ratios,
we havesummarized
in Figure(1987).
9 a logical highest masses, despite the fact that both methods are predomP05. electron
semiempiricalmethodof Veilleux& Osterbrocktemperature.
Eventually,explicitly,
at still higher
metallicities,
izationparameter
based
upon nitrothe
process
whereby asofmany
estimatesof boththe chemicalOurcalibration
this
ratio
is
shown
in
Figure7, with inantly based on H ii regions with Te metallicities. At the lowest
methodisavailable
for
only
546/27,730
(2%)
of
The
directT
gen becomes
the dominant coolant in the nebula, and the e abundance and2014年12月2-3日
系外銀河における微細構造輝線の観測とその理解@国立天文台theoreticalmodels.
ionization
parameters
as arecompatible
polynomial
fits asinequation(4)and coefficients
given in stellar masses, this difference disappears. The difference between
electron temperature falls sufficiently to ensurethat thein our
the galaxies
SDSS
sample.
The
[O
iii]
k4363
line
is
weak
with theTable
3. data set can be obtained using the techniques
nitrogen line4.4.
weakensincreasing
metallicity.
Since only
described
inlow
this
section.
Thethisprocess
canscalesbeTheautomated,
and is usuallyobserved
inmetallicity,
metal-poorgalaxies.
SDSS
Diagnostic
Diagram
The S23with
At
very
ratio
simplyasand
the the empirical methods may be attributed to the different H ii rethe [N ii] line is produced in the low-excitation zone ofanIDL
script
to
do
this
is
availableonrequest
from
the
first
contains
very
fewabundance,
metal-poor
theynitrogen
to galaxies
first order.because
However,
it isareknown gion samples used to derive the calibrations. At the highest metOurof theratiois ([S!!6717,catalog
6731
+[S
iii]
the
H calibrationii region, [Nii]/H!
alsoii] sensitivetoionization
author (L. J. K.). However, we will describe a more directin thisand
metallicity
regime,the nitrogen
abundance allicities, the PP04 methods utilize four H ii regions with detailed
compact,
faint (e.g.,
Terlevich
et al.in1991;!!9069,9532)/H" (popularly known as S23intrinsically) is shown in rare, thatparameter.
methodfor the derivationof theseparameters,
x 7 below.
No. 1, 2002
ABUNDANCESIN EXTRAGALACTICHREGIONS
41
Masegosa
et al. 1994;
vanZeeii 2000).
Panel
(10)
ofauthors’
Figuretechni1
From thecomparative
analysiswithother
4.6. The ½N ii"=½O iii" Diagnostic Diagram
ques
which
follows,
weareconfident
that
this
new
techniqueshows that a total of 477 Te metallicities is insufficient to obtain a
1.0
1.5
Figure6. Again,we (7)fit fourth-orderdefined as
q=3e8
(7)
will
providethe most
abundancepolynomials
estimates currentlyThe
advantage ofusing[N ii] and [O iii] lines is that theyq=3e8reliable
1.5e8 (6)clear
M-Z relation.
Becauseweareunableto fitZ an
M-Zrelation
1.5e8 (6)=are
0.5 solarinequation
(4),
and
the
coefficients
given
in
Table
3.
8e7 (5)possible.
are unaffectedbyabsorption
lines
originatingfromunderly8e7 (5)Kewley& theyDopita02 usinglines
4e7 (4)
do notTe ofmethodfurtherin
Te metallicities,
4e7 all(4) suchtheAs
is thecaseconsider
inratiosforbidden
to recombina0 we
2e7 (3)
ing1.0stellar populations,lie close to Balmerthat
2e7 (3)1e7 (2)has a maximum at a certain metaltion lines, the S23
1e7 ratio(2)
can be used to eliminate
errors due to dust reddening,
this work.and5e6 (1)
5. COMPARISONWITH OTHER
5e6 (1)7
0.5
licity,andtherefore
itlargeis twofor
valuedat all other metallicities.
they areboth strong and easily observable in The
the optical.
6
scatterinthe
M-Z
relationisall metallicity
calBRIGHT-LINE
TECHNIQUES
0.5
5
For
this
particular
ratio
the
maximum
occurs at a somewhat
Both empiricalandtheoretical
relationships
for
the
[N
ii]/
4
ibrations;therms
residual
about
the
line
of
best
fit
is
0.08–0.13.
Comparison[O iii] 3ratio as a function of oxygen abundance currently
higher abundance5.1.
than
for the RData
23 ratio; at metallicities of
2-1
The cause of the roughly
scatterinthe
M-Z
relationisunknown.
Our
exist
Considère et al. 2000).
solar
[log
ðO=HÞ
þ
12
$
8:8].
Again, to10.0 (e.g.,
Most of the data sets previously used
to compare
andraisecali- theOurcalibration
of
the
[N
ii]/[O
iii]
ratio
is
shown
in
degeneracy
inthesolutions,aninitialguessofthe metalcomparison between
differentdiagnostics
metallicity
0.0
bratetheabundance
hadcalibrations
been selected showsin different
Figure 8. The fourth-order polynomial fit coefficientsare the andlicitymust firstbe obtainedfromangalaxyalternativethat differing
ionization
parameter
doesdiagnostic.
notH ii
heterogeneousways:amongsomebygalaxies(brightest-0.5 in Table 3.
given
dependent
onionization
parameter
S232 1isorquite
regions,brightest
disk
Hii
regions),someby objective
causeor
contributetothescatter.
The
ionization
parameter
is for all
Because
the
two
ions
have
quite
different
ionization3
metallicities,
therefore
the ionization1
prism
searchesand(whichare biased
towardparameterstrong [Oderived
iii]
..
-2
explicitly
calculated
andintoii]by
account
some
metallicity
potentials,
the [N ii]/[O iii] ratio depends strongly
on the4thetaken
4from[S iii]/[Sdiagnostic
shouldtoirregueliminate""4959,5007),
some
Galaxy in
type(suchbeasused
dwarf
-1.0
-0.5
ionization
parameter. Thus, if this diagnostic
is
to
be
use5
5 a freediagnostics ( KD02;
M91),
but
see ainreduction
thisKK04;
asvariable.
lars).
These
different
datawesetsdoarenotreflected
the differences
ful for6 abundance
it must be used in
q=3e8determinations,(7)
6 the various
betweencalibrationsof abundance.
Care must
1.5e8 (6)inscatterfor
thesemethods.
A
full
investigation
into
thescatter
combination
with8e7an(5) independent ionization parameter-1.5
bewilltaken
therefore
when
comparingdifferent
abundance
4e7 (4)4.5.
The½N
ii&=H#
Diagnostic
Diagram
7
in
the
M-Z
relationbepresented
inS.
L.
Ellisonetal.
(2008,
indiagnostic.
However,
2e7 (3) if the [O ii] or [S ii], and [S iii] lines
7
diagnostics
to take into account biases (if any) introduced-3
1e7 (2)
Primary Nitrogen
Secondary
Nitrogen-1.0
5e6 (1)Intheabsenceofotheremissionlines,
the[Nii]/H#
line
preparation).
by the comparison data.
-2.07.5
8.0
8.5
9.0
9.5
ratio
canthebetouseda M-Z
crude8.5
estimatorof9.0themetallicity.
Note
LOG (O/H)
8.0
8.5 + 12
9.0 We9.5
40 7.5KEWLEY&compare
DOPITA
Vol. 142
directly
best-fit
curvesfornine
7.5 chose
8.0 as
9.5
We
use
observationsofHii
regions
available
LOG (O/H) + 12
+ 12that thethelarge[N ii]/H#
ratioLOG
is (O/H)
particularly
from
and2,
homogeneous
data set
van P05.
Zeeto
et shockal.
strong-line
calibrationsin Figureincludingboth
P01ofsensitive
and
4
1.5
Fig.
6.—Thelog(([Sii]!!6717,6731 +[Satiii]high!!9069,9532)/H") (Sq=3e8 (7)23 )excitation
orauthors
the q=3e8
presenceof a185
hard
ionizing
radiation
field,
very difficult to observe.
Second,metallicity,the(1998).
These
observedHii
regionsin13
spiral(7)
1.5e8 (6)vs. metallicity. Curves for each ionization
Fig.
7.—Thelog([Nii]!6584/H#)diagnostic
forabundancevs.
metaldiagnostic
forabundance
1.5e8in(6) metallicity
The
toptheanrms
scatterthemeanFig. 5.—Thelogtemperature(([O8e7ii](5)!3727 givesþ ½O iii] !!4959,
5007)/H")(R23panel) diagZ =about
0.5on
solar
from
AGN.
The
of between
strongshock
excitation
lower
electronthermalelectronsof shows
galaxieswith
theDoubleSpectrographthePalomar5m
8e7 (5)presence)1
6 to 3 ( 108fewer
6 to 3 (
8 or
Z
=
0.5
solar
cm
s
are
shown.
Filledcircles
10
licity.
Curvesfor
eachionization
parameter
q
¼
5
(
10
parameter
betweenq¼
5(
10
4e7metallicity.
(4)
nostic for abundance vs.
Curves for each ionization parameter
4e7 (4)
)
¼
0:2.
The
major
difference
in
mass
bins
of
width
(M
/M
an
AGN
will
increasethe[Nii]/H#
ratio
and
cause
the
high
energy,leadingtoastrongdecreaseinthenumberof)1!log2e7from(3)!
telescope.
These
data
have
the
additional
advantage
of
cov)1
6
8
2e7
(3)
7
cm
s
are
shown.
Filledcircles
representthedata
pointsfrom
ourmodelsrepresentthedata
pointsourmodelsatmetallicities
from
left
to
right
3(
are shown. Filled circles represent
between q ¼ 5 ( 10 to1e7(2) 10 cm s
1e7 (2)collisional
excitations
theat3.0
blueii] thelines(whichhas aofoftherelabundances
determinedbeartificially
high.
IfThe
the
. [Seeelectronic
the M-Zat.
metallicities
left5e6to(1)metallicity
right toof 0.05,
0.1, 0.2,
0.5,
1.0,
1.5,
2.0,
3.0[N
Z*ii]/
of 0.05,
0.2,
0.5, 1.0,
Z*[O
ering
range
in
and
ionization
parameter,
5e6 1.5,(1)of 2.0,
1.0a large
betweencurvesis from
theirpositionalong
they-axis.
thedata 0.1,
pointsfrom
ourmodelsmetallicities
from
left
toedition
right
0.05,
6
2 0.5,
Journal
for
a1.0,
color
versionof this[Seetheratio
electronic
editionof theet
Journalfor aopticalcolor versionof this figure.]
ativelyhighthresholdenergyforelectronic
excitation)
relativetothefor
is to
be
used,
diagnostic
diagrams[See theeditionof the Journal
0.1,
0.2,
1.5,
2.0, 3.0
Z*.figure.]
asH#
was
shown
in
Dopitaal.photoionization(2000).
curveswiththe largest
y-intercept
are standard
all
model
5
alowercolor version
thisii]figure.]energyof[Nlines.
shouldfirst
applied to rulepossibility
presSince [Siii]bemeasurements
areoutnottheavailable
for ofthethevan
4
based(KK04;
Z94;
KD02;
T04;
M91).
Among
these
photoFor Z < 0:5 Z! [log ðO=HÞ þ 12 < 8:6], the metallicity
Zee
H of
ii regions,also used
two additional
data sets
ence
an AGNweorhave
shockedexcitation.
We recommend
the3
ionization
model formetallicities,
the agreement
is&optical
#0.2
dex.
This anddependence of the [N ii]/[O ii] ratio is lost because
nitrogen
Stasińska(1983)
theuse0.5of StheKewleyet Dennefeldal. (2001b)
diagnostics
2 since23 diagnostic;
factor’’ or ‘‘ excitation
parameter’’ to correct
the observed
1 here(like 0oxygen)is predominantly
a primarynucleosynthesis
Kennicutt&basedGarnett(1996).
These
also
cover
a widerange
these
the sametheoreticalmodelsusedagreement
is within
thearemarginofonerror
typically
cited
for
these
1ionization parameter, such as in Pilyugin(2000);
R
23 for ininand
ionization
parameterandmetallicity.
Using
the previous
ESO
element
this metallicity range. In addition, the nitrogenhave
been
shown
to
be
more
reliablethanthecalibrations(#0.1–0.15
dex for eachcalibration).
SomecalibraCharlot
ion2& Longhetti (2001). We prefer to estimate
3.6
m telescope,method
Dennefeld& Stasińska(1983) observed
to-oxygen
abundance ratio shows large scatter fromthe
objectsemiempiricalof Veilleux& Osterbrock(1987).
0.0
izationparameter
explicitly,
based
upon
the
tions
consistently
agree
towithin0.1
dex(e.g.,
KK04
and
Z94;
#40Our
H iicalibration
regions in the
Galaxy,3 In this regime, nitrogen production increases as ato object.
of Magellanic
this ratio isClouds
shownandin theFigure7, withtheoretical
Kennicutt(1996)
observed
a similar numKD02
and M91).while
Comparisons
betweenmetallicities
calculated
function4 ofmodels.
time since the bulk of the starformation
polynomial
fits&
asGarnett
in equation(4)and coefficients
given in-2
ber
ofmethods,
H3.ii regionsin M101
using andthe 2.1
m telescopeat Kitt
5 (Edmunds & Pagel 1978; Matteucci & Tosi 1985;
occurredTableusing these consistentsuch
as KD02
M91,
are likely-0.5
6 al. 1997).
Peak.
Dopita et
Therefore,
for a sampleof H ii regions,
Diagram
4.4. TheS23 Diagnostic
At
very
low
metallicity,
this
ratio
scalessimplyKewley & Ellison 08
to be reliable to within 0.1 dex. However, comparisons between as the7
the varying
age distribution of the stellar population from
nitrogenabundance,
to first order.
However,
it is known
Our calibration of the ratio ([S ii] !!6717,
6731
+[S
iii]
5.2.
Comparison
Techniquesfor
Deriving
Abundancethat showthat
large
disagreement(such
as KK04 andP05)
object to object will cause scatter in themethods
N/O ratios
in thisPrimary
metallicity
regime, the nitrogen
abundance
Nitrogen
Secondary Nitrogen!!9069,9532)/H"(popularly known as S23) is showninPrimary Nitrogen
Secondary-4
observed.
Good evidence for the primary dependence
of
-1.0
will
beNitrogencontaminatedby
the large
discrepancy
betweenAs
discussed
in xsystematic
1, a wide numberof empiricaland semi7.5
8.0
8.5
9.0
9.5
7.5
8.0
8.5
9.0
9.5
nitrogen at low abundances
4 ofempirical approaches already
LOGcan(O/H) +be
12 found in
LOGexist(O/H) + for
12 the determination of
theFigurecalibrations.
Considèreet al. (2000), in which log ðN=OÞ derived from aabundances in H ii regions. We compare three of the most
1.0
The
lowest curves
in Figure2 q=3e8
are(7)the produced
M-Z relations
derived(7)
large
sampleiiq=3e8
regions
is iii]comparedwith thecommonlyusedcalibrations
by McGaugh(1991,
Fig.
8.—Theof
logH
([N
ii]
"6584/[O"5007) diagnostic
for expected
abundance
1.5e8(6)Fig. 4.—Thelog ([Nii]!6584/[Sii] !!6717,Z6731)
diagnostic for
abunFig. 2.—Robust best-fit M-Z relations calculated using the different metallicity
1.5e8
(6)
6
= 0.5 solar
8e7
(5)
using
the
empirical
methods(i.e.,
P01,
P05,
and
the
two
PP04
vs.
metallicity.
Curvesfor
eachionization
parameter
betweenq¼5%10
relations for &1
a primary,secondary, and primary + seconhereafterM91), Zaritsky,
Kennicutt,& Huchra(1994,
here8e7 (5)
dance
vs.
metallicity.
Curvesfor
eachionization
parameter
between4e7(4)8 cm scalibrations
listed in Table 1, except the Te method. The top panel shows the rms
to
3
%
10
are
shown.
Filledcircles
representthedata
points
from
4e7
(4)
*1
8 cm s0Z94),(3)
andCharlotLonghetti(2001,
C01)
dary origin for 2e7
nitrogenproduction (Vila-Costas&
3 + 10areshown. Filled
circles hereafter
represent thedataqafter¼5+
106 tomethods
methods). Theseempirical
are&calibrated
predominantly
2e7 (3)1e7 (2)from left to right of 0.05, 0.1, 0.2, 0.5, 1.0, 1.5,
our models at metallicities
scatter in metallicity about the best-fit relation for each calibration in 0.1 dex bins ofpointsfrom
ourmodelsatmetallicities
from
left
to
right
of
0.05,
0.1,
0.2,
1e7
(2)
5e6
(1)
with
the
results
producedbyourproposed
theoretical
Edmunds[See the electronic edition of the Journal for a color
2.0, 3.0 Z'. 1993).
via version
fits ofof thisthe relationshipbetweenstrong-line
H ii forre-a color stellar mass. The y-axis offset, shape, and scatter of the M-Z relation differ sub(1)the electronic ratios[See
edition ofand
the Journal
0.5,
1.0, 1.5, 2.0,
3.0 Z!. 5e6
0.5
diagnostics.
We
conclude that [N ii]/[O ii] provides an excellent abunfigure.]versionof
thisfigure.]stantially, depending on which metallicity calibration is used.
metallicities.
There
isconsiderable
variation
among
the
gion
T
e
be
dance diagnostic for Z > 0:5 Z!, but this ratio cannot
used at lower abundances.
-1
For Z > 0:5 Z!, the curves in Figure 3 can be fitted by aprimarily
to the different nucleogenic status of the two elesimple
quadratic, facilitating abundance determination
ments. At low metallicity, both the elements are primary
0.0
using the [N ii]/[O ii] ratio, i.e.,
and the ratio becomes insensitive to metallicity. This diag1
nostic is2not
as useful as [N ii]/[O ii] for the determination oflog ðO=HÞ þ 12 ¼ log ½1:54020 þ 1:26602 R

Metallicity calibration with optical lines：[SII]/[NII]

LOG ( [NII] / H" )[NII]/HαN2

LOG ( [NII] / [SII] )LOG{ ({ [OII]+[OIII]
LOG( [SII]+[SIII]) )/ /H!
H!} }LOG ( [NII] / [OIII] )[NII]/[OIII]([OIII] + [OII])/Hβ

R23

O3N2

N2S2

LOG ( [NII] / H" )G { ( [SII]+[SIII] ) / H! }Empirical/theoretical metallicity calibration by using
strong optical lines (note that systematics exist)

系外銀河における微細構造輝線の観測とその理解

2014年12月2-3日 @国立天文台

Metallicity calibration with FIR FSLs：
A&A 526, A149 (2011)

Fig. 3. Predicted emission-line flux ratio of [O]51.80/[N]57.21 as afunction of gas metallicity. Blue and red lines denote the models with
U = 10−2.5 and 10−1.5 , and solid and dashed lines denote the models
with nH = 101.0 cm−3 and 103.0 cm−3 , respectively. The flux ratio observed in M 82 and in the Antennae galaxy is shown by black horizontal
line (these two galaxies show very similar [O]51.80/[N]57.21 flux
ratios; see Table 5). The x-range of this horizontal line corresponds tothe inferred metallicity range for M 82 and the Antenna galaxy.
Fig. 3. Predicted emission-line flux ratio of [O]51.80/[N]57.21 as afunction of gas metallicity. Blue and red lines denote the models with
U = 10−2.5 and 10−1.5 , and solid and dashed lines denote the models
with nH = 101.0 cm−3 and 103.0 cm−3 , respectively. The flux ratio observed in M 82 and in the Antennae galaxy is shown by black horizontal
line (these two galaxies show very similar [O]51.80/[N]57.21 flux
ratios; see Table 5). The x-range of this horizontal line corresponds tothe inferred metallicity range for M 82 and the Antenna galaxy.

Fig. 5. Same as Fig. 3 but for the flux ratio of ([O]51.80+[O]88.33)/[N]57.21. Note the much lower dependence on the gas
density, which makes this ratio particularly suited to measure the gas
metallicity.

instead of [O]51.80/[O]88.33. Since the [N]205.4 emission is very faint and at a very long wavelength, it is very difficult to study with Herschel and SPICA, though its detection
Fig. 5. Same as Fig. 3 but for the flux ratio of ([O]51.80+
should be feasible with ALMA.
[O]88.33)/[N]57.21. Note the much lower dependence on the gas

Metallicity calibration by using
3.3. Dependenceson the stellar
age
FIR[OIII]52µm,[OIII]88µm,
All emission-line flux ratios shown in Figs. 3−7 are calculated[NII]57µm
by adopting constant star-formation SEDs with an age of 1 Myr.

density, which makes this ratio particularly suited to measure the gas
metallicity.

insteadof it[O]51.80/[O]88.33.
Since[N]205.4
However,shouldbe verified
whetherthethepredicted
flux emisratios
sion
is
very
faint
and
at
a
very
long
wavelength,
it
is
very
difdepend on the age of the stellar population, because the adopted
ficult
to
study
with
Herschel
andSPICA,
thoughits
detectionage (1 Myr) seems too young compared to the typical age ofshouldbe feasible
withinALMA.
star-forming
galaxies
general. Cid Fernandes et al. (2003) investigated stellar populations of nearby starburst galaxies andreported
that the typicalstarburst
age is ∼107.0−7.5 yr. More re3.3.
Dependenceson thestellar age
cently, Rodríguez Zaurín et al. (2010) studied stellar populationsof low-z
ULIRGs,flux
findingsimilarageinranges.
It isaretherefore
imAllemission-line
ratiosshown
Figs. 3−7
calculated
portant
to examine
howthe emission-lineinbyadopting
constantstar-formation
SEDs diagnostics
with an age studiedof 1 Myr.
this work it
dependage ofwhether
the stellar
usedforHowever,shouldonbethe
verifiedthe populations
predicted flux
ratios→ “SPICA” is necessary！How
about ALMA?

Nagao+11
Fig. 4. Same as Fig. 3 but for the flux ratio of [O]88.33/[N]57.21.

系外銀河における微細構造輝線の観測とその理解

2014年12月2-3日 @国立天文台

Metallicity of galaxies at high redshift：
1924

F. Mannucci et al.

metallicity

Mannucci+09, MNRAS, 398, 1915

stellar mass
Figure 5. Evolution of the mass–metallicity relation from z = 0.07 (Kewley & Ellison 2008) to z = 0.7 (Savaglio et al. 2005), z = 2.2 (Erb et al. 2006a)
and z = 3–4 (AMAZE+LSD). All data have been calibrated to the same metallicity scale and IMF (Chabrier 2003) in order to make all the different results
directly comparable. Turquoise empty dots show the AMAZE galaxies, blue solid dots the LSD galaxies. The solid square shows the ‘average’ LSD galaxy,having average mass and composite spectrum (see Fig. 4). The lines show quadratic fits to the data, as described in the text.

Metallicity measurements are limited up to z~3

系外銀河における微細構造輝線の観測とその理解

2014年12月2-3日 @国立天文台

Metallicity of galaxies at high redshift：
1924

F. Mannucci et al.

metallicity

Mannucci+09, MNRAS, 398, 1915z>4
stellar mass

Figure 5. Evolution of the mass–metallicity relation from z = 0.07 (Kewley & Ellison 2008) to z = 0.7 (Savaglio et al. 2005), z = 2.2 (Erb et al. 2006a)
and z = 3–4 (AMAZE+LSD). All data have been calibrated to the same metallicity scale and IMF (Chabrier 2003) in order to make all the different results
directly comparable. Turquoise empty dots show the AMAZE galaxies, blue solid dots the LSD galaxies. The solid square shows the ‘average’ LSD galaxy,having average mass and composite spectrum (see Fig. 4). The lines show quadratic fits to the data, as described in the text.

Metallicity measurements are limited up to z~3

系外銀河における微細構造輝線の観測とその理解

2014年12月2-3日 @国立天文台

Metallicity of galaxies at high redshift：

Strong emission lines
used to calibrate themetallicity comethrough the observable
wavelength (>2.5µm)

系外銀河における微細構造輝線の観測とその理解

2014年12月2-3日 @国立天文台

Metallicity of galaxies at high redshift：
1924

F. Mannucci et al.

Mannucci+09, MNRAS, 398, 1915

metallicity

JWST?
SPICA?
orz>4stellar mass

ALMA?

Figure 5. Evolution of the mass–metallicity relation from z = 0.07 (Kewley & Ellison 2008) to z = 0.7 (Savaglio et al. 2005), z = 2.2 (Erb et al. 2006a)
and z = 3–4 (AMAZE+LSD). All data have been calibrated to the same metallicity scale and IMF (Chabrier 2003) in order to make all the different results
directly comparable. Turquoise empty dots show the AMAZE galaxies, blue solid dots the LSD galaxies. The solid square shows the ‘average’ LSD galaxy,having average mass and composite spectrum (see Fig. 4). The lines show quadratic fits to the data, as described in the text.

Metallicity measurements are limited up to z~3

系外銀河における微細構造輝線の観測とその理解

2014年12月2-3日 @国立天文台

Chemical abundance ratio at high redshift:

J. Köppen and G. Hensler: Chemistry of epi

lg(N/O)

0

-0.5

-1

100

-1.5

-2

1

Mcloud /Mgal

6.5

7

Koppen & Hensler 05

30

7.5

10

3

8

8.5

9
lg(O/H)

9.5

Fig. 2. Tracks in the N/O–O/H-diagram of infall models with various
Effect of gas inflow on chemical evolution
mass ratios of the infalling gas and the galaxy. The infall event starts

Fig. 3.
masses系外銀河における微細構造輝線の観測とその理解

2014年12月2-3日 @国立天文台

multaFigure 4. SED of A383-5.2 and population synthesis models which provide
3
e CIII]
best fit to the continuum SED and CIII] equivalent width. The observed
SED
is denoted by the diamond data points. The two grey data points at
ar for8Stark
et al.
<1µm are not includedin the fit because of the uncertainty
with (arXiv:1408.3649)
CIII]
nsiderStarkassociatedet al. 2014
Ly↵ emission contamination and IGM absorption. The data are best fit by arburst’3000
model with a two component star formation history. The UV continuum andn withz=6−7, This work
20
CIII] equivalent width
arepoweredbyarecentstar
formation
episode(cyan
z=1.5−3.0, Starket al. 2014
A383−5.2
stellar2500
z=2.3,
Erb
et
al.
2010
curve), while the opticalcontinuum
is dominated by an older generationofΔvLyα= 120 km/s
z~3, Shapley et al 2003
th age(composite
spec.)
stars (red curve). The composite SED is shown in black. Yellow
diamonds
2000
shift).
show predicted broadband fluxes from the best-fitting model.
15
match-1500
1909Å
t comst leftmost
pro-1000panel is centred on Ly↵ emission which is detected with the visible
he
10
iddle
(unsmoothed)andrightmost(smoothed)panels.
The[CIII]
1907
emissionhestarModel fit to A383.5.2
500
mationlog U
1.70+0.49
fitting. 0
0.64
5
+0.04
log (M⇤,young /M⇤,tot )2.99 0.03
ing of−500
ware (v.2.2.0) in the Recipe Flexible Execution Workbench+0.27 (RE−200 log (Z/Z
0
200
400
600
800
1000
)
1.33
the al0.20
FLEX) environment to perform
a firstreductionofVel (kms−1) calibration and+0.06
log(C/O)
0.58
mass
0.06routines
each exposure. We then applied standard IDL and IRAF
WCIII],0 (Å)

Fλ (10−20 erg cm−2 s−1 Å−1)

Chemical abundance ratio at high redshift:

+0.10

0

Figure5. Velocity
profile of Ly↵emissionin the zthe
= 6.027
galaxy A383-Speciflog(age/yr)
for
optimally
combining
andextracting
15 8.72
exposures.
0.10
5.2. Ly↵ is shifted to the rest-frame using the systemic redshift provided
by
+0.10
−50
0
50
100
150
ically,we usedthe Lyman-↵emissionline,bywell-detectedin eachlog(M/M) is shifted
9.50
⇤,tot
n from
CIII] 1909. The
peak fluxof Ly↵
emission
v=120 kms 1
0.10
WLy_,0 (Å)+0.08 condiexposure,
to correct
for variationsin1 )seeing and 0.29
atmosphericfrom the systemic
redshift.
log(SFR/M
yrg wide0.08
tions between the different OBs, and applied a scaling
and weight+0.05
Note that they are gravitational
⌧
ˆ
0.05
model
0.05level ofFigure 6. A comparison of the rest-frame equivalent widths of CIII] 1909
ing of the 2D V
spectra according to the flux and detectionÅ and Ly↵ fromlensed
the study of objectsz ' 2 lensed sources by Stark et al (2014,
the
line.
We
also
used
the
spatialposition
of
the
line
edshift
6 Lyman-alphaDISCUSSIONAND
SUMMARY
orangepoints),Erbetal(2010)
Table 1. Results of fitting procedure for A383-5.2. Details are providedandin the z' 3 Shapley et al (2003) composite
measured in the reduced spectrum to precisely computethe off-alongside the two new CIII] detections beyond z ' 6 (this paper, red stars).
nt; 30Considerable
effortandhas beenplacedspectroscopic
6§4.1,
resultsareindiscussed
in study§4.2.of z >⇠sets
betweeneach
OBfor optimal
combination.
We usedgalaxies.
Yet, inspite
of significant
allocationsof telescope
timetheto sameompooffsets,
scaling factors
andspectroscopic
weights to combine
the exposuresseveral researchteams, few
redshifts havebeen con- in theicitiesnear-infraredcorrections
slightly
strengthened
firmed beyondarm.
z 'Applying
7 due to these
the absenceof Ly↵.
AlthoughtheDetection of CIII]λ1909 from galaxies at z~6系外銀河における微細構造輝線の観測とその理解

2014年12月2-3日 @国立天文台

Newman et al.
Morphology and Kinematics at high redshift：

nal, 752:111 (19pp), 2012 June 20

Forster Schreiber+09

Clumpy structure in highredshift galaxies
Strong Galactic winds inindividual clumps of highredshift galaxies

Dust geometry in highredshiftgalaxies
Newman et al. 2012, ApJ, 752, 111
Figure 2. Emission line profiles and best fit to Hα and [N ii]λ6584 features for six regions of ZC406690: clump A

系外銀河における微細構造輝線の観測とその理解

MIPS
J1428 [O I] 63
2014年12月2-3日@国立天文台MIPS J1428
F10214F10214F10214[O III] 88[O IV] 26[S III] 33[O III] 52

205
205
179
185
179

Connection between galaxies and AGN：
BPT★☆: Stacking analysis

1 342 187 779
1 342 187 779
1 342 186 812
1 342 187 021
1 342 186 812

1
7
10
6
s

Notes. (a) Number of line/range repetitions per nod cycle × number
calibration uncertainty of 30% applies; (c) Correction factor, applied
central spatial pixel (see text); (d) The signal in this PACS wavelengtof upper limits and RMS unreliable; (e) assuming a Gaussian profile

Sturm+10, A&A, 518, L36[O IV] (25.91µm) / FIR

10−2

Kewley+01
Yabe et al. 2014

AGN diagnostics byusing FIR fine structurelines ?

NGC 1068
F1021410−3

Mrk 231

10−4

Arp 220

10−5

M82

108

109

1010

1011

1012

1013

LFIR [ L ]Fig. 2. The [O IV]/FIR limit in F10214 compared to template obje

系外銀河における微細構造輝線の観測とその理解

2014年12月2-3日 @国立天文台

What I want to know：• Star formation rate of galaxies• Metallicity of galaxies• Distribution of star-formation• ISM condition / AGN connection... from dust-free observations
spatially resolved (if possible) ...