Submitted to The Astrophysical Journal
Preprint typeset using LATEX style emulateapj v. 08/22/09
MAPPING THE YOUNGEST GALAXIES TO REDSHIFT ONE1,2
Yuko Kakazu,
Lennox L. Cowie,
Esther M. Hu,
Submitted to The Astrophysical Journal
ABSTRACT
We describe the results of a narrow band search for ultra-strong emission line galaxies (USELs)
with EW(Hβ) ≥ 30 Å. 542 candidate galaxies are found in a half square degree survey using two
∼ 100Å filters centered at 8150Å and 9140Å with Subaru/SuprimeCam. Followup spectroscopy has
been obtained for randomly selected objects in the candidate sample with KeckII/DEIMOS and has
shown that they consist of [O III]λ5007, [O II]λ3727, and Hα selected strong-emission line galaxies at
intermediate redshifts (z < 1), and Lyα emitting galaxies at high-redshift (z >> 5). We determine
the Hβ luminosity functions and the star formation density of the USELs, which is 5-10% of the
value found from ultraviolet continuum objects at z = 0 − 1, suggesting that they correspond to
a major epoch in the galaxy formation process at these redshifts. Many of the USELs show the
temperature-sensitive [O III]λ4363 auroral lines and about a dozen have oxygen abundances satisfying
the criteria of eXtremely Metal Poor Galaxies (XMPGs). These XMPGs are the most distant known
today and our high yield rate of XMPGs suggests that narrowband method is powerful in finding
such populations. Moreover, the lowest metallicity measured in our sample is 12+log(O/H)=7.06
(6.78−7.44), which is close to the minimum metallicity found in local galaxies, though we need deeper
spectra to minimize the errors. HST/ACS images of several USELs exhibit widespread morphologies
from relatively compact high surface brightness objects to very diffuse low surface brightness ones. The
luminosities, metallicities and star formation rates of USELs are consistent with the strong emitters
being start-up intermediate mass galaxies which will evolve into more normal galaxies and suggest
that galaxies are still forming in relatively chemically pristine sites at z < 1.
Subject headings: cosmology: observations — galaxies: distances and redshifts — galaxies: abundances
— galaxies: evolution — galaxies: starburst
1. INTRODUCTION
The study of low-metallicity galaxies is of consider-
able interest for the clues that it can provide about the
first stages of galaxy formation and chemical enrichment.
We would also like to know if there are any genuinely
young galaxies undergoing their first episodes of star for-
mation at low redshifts. To date, the most metal-poor
systems studied have been the blue compact emission-
line galaxies found in the local Universe, with systems
such as I Zw 18 and SBS 0335-052W defining the low
metallicity boundary with measured 12+log (O/H) of
∼ 7.1 − 7.2 (Sargent & Searle 1970; Thuan & Izotov
2005; Izotov et al. 2005). More recently, the Sloan Digi-
tal Sky Survey (SDSS) has yielded additional extremely
metal-poor galaxies (XMPGs) (12+log (O/H) < 7.65 or
Z < Z⊙/12; Kniazev et al. 2003; Izotov et al. 2006a).
Despite enormous efforts, only a few dozen such XMPGs
are known, all at redshift z < 0.05 (e.g., Oey 2006; Izotov
2006b).
Historically, objective prism surveys have been used
to select emission-line galaxies for low-metallicity stud-
ies. (e.g. the Hamburg QSO Survey (Popescu et al.
1996) and its HSS sequel (Ugryumov et al. 1999) that
1 Based in part on data obtained at the Subaru Telescope, which
is operated by the National Astronomical Observatory of Japan.
2 Based in part on data obtained at the W. M. Keck Observa-
tory, which is operated as a scientific partnership among the the
California Institute of Technology, the University of California, and
NASA and was made possible by the generous financial support of
the W. M. Keck Foundation.
3 Institute for Astronomy, University of Hawaii, 2680 Woodlawn
Drive, Honolulu, HI 96822.
discovered HS 2134+0400 (Pustilnik et al. 2006) and
the Kitt Peak International Spectroscopy Survey (KISS;
Salzer et al. 2000; Melbourne & Salzer 2002)). The ad-
vantage of using the objective prism technique rather
than the continuum selection, employed with the SDSS
(Kniazev et al. 2003) or DEEP2 surveys (Hoyos et al.
2005), is that they have a higher efficiency and
provide a more uniform selection. By comparison,
continuum/broad-band surveys have a very low yield rate
(8 new XMPGs and 4 recovered XMPGS from an anal-
ysis of 250,000 spectra over ∼3000 deg2 for the SDSS
(Kniazev et al. 2003), since low-metallicity populations
in their first outburst have weak continuua and strong
emission lines.
An alternative method of discovering strong emission-
line, low-metallicity galaxies is to use narrowband
surveys. Strong emission-line galaxies have histori-
cally been picked up in high-z Lyman alpha searches
(e.g., Stockton & Ridgway 1998; Hu et al. 1998, 2004;
Stern et al. 2000; Tran et al. 2004; Ajiki et al. 2006)
where they have been considered contaminants. How-
ever, the low redshift emission line galaxies seen in these
surveys are of great interest in their own right as we shall
show in the present paper. While some spectroscopic
studies have been carried out for low-redshift galaxies se-
lected from narrowband surveys (e.g., Maier et al. 2006;
Ly et al. 2007), the small sample sizes have precluded
any detailed investigation of metallicity and identifica-
tion of a low-metallicity population.
The narrowband method probes to much deeper limits
than the objective prism surveys. This enables prob-
http://arxiv.org/abs/0704.0643v1
2 Kakazu et al.
ing star-forming populations out to near redshift z ∼ 1
where the cosmic star formation rates are considerably
higher. Furthermore the narrow-band emission-line se-
lection can allow us to assemble very large samples of
strong-emission line objects, with a clean selection of dif-
ferent line types for the construction of luminosity func-
tions.
Such a sample allows us to address such questions
as whether there are substantial populations of strong
star-forming galaxies with low metallicities among more
massive galaxies. There has been considerable contro-
versy about the interpretation of the low metallicity mea-
surements in the blue compact galaxy samples where
the ease with which gas may be ejected in these dwarf
galaxies has complicated the picture (e.g., Corbin et al.
2006) or, at least, resulted in identifying low metal-
licity systems which are not forming their first gen-
eration of stars. The identification of low metallicity
galaxies – at the level of the XMPGs – among mas-
sive galaxies can provides less ambiguous examples of
galaxies that are genuinely ‘young’ and caught in the
initial stages of star formation. Current efforts to iden-
tify low metallicity galaxies from continuum selected
surveys (e.g., Kobulnicky et al. 2003; Lilly et al. 2003;
Kobulnicky & Kewley 2004; Hoyos et al. 2005) have low-
metallicity thresholds that are higher than this – about
one-third solar (in O/H). With a narrow-band selec-
tion criterion much larger emission-line samples includ-
ing such low metallicity galaxies can be identified. With
these large samples it is also possible to determine
whether there is an observed lower metallicity threshold
for such galaxies, and to estimate what the contribution
of such strong star-formers might be at these epochs.
In the present work we use a number of deep, narrow-
band images obtained with the SuprimeCam mosaic
CCD camera (Miyazaki et al. 2002) on the Subaru 8.2-
m telescope to find a large sample of extreme emission-
line galaxies. We first (§2) outline the selection criteria
(magnitude and flux thresholds) for the target fields re-
sulting in a sample of 542 galaxies, which we call USELs
(Ultra-Strong Emission Liners). We then describe (§3)
the spectroscopic followups for 161 of these galaxies us-
ing multi-object masks with the DEIMOS spectrograph
(Faber et al. 2003) on the 10-m Keck II telescope. Sam-
ple spectra for each class of object are shown. Flux
calibration and equivalent width distributions are pre-
sented in §4, and the resulting measured line ratios are
discussed. In §5 luminosity functions are constructed and
star formation rates are estimated for the sample. These
galaxies are estimated to contribute roughly 10% to the
measured star-formation rate (without extinction correc-
tions) at this epoch. Analysis of the metallicities is given
in §6. Their morphologies and dynamical masses are dis-
cussed in §7 and a final summary discussion is given in
§8. We use a standard H0 = 70 km s
−1 Mpc−1, Ωb =
0.3, ΩΛ = 0.7 cosmology throughout.
2. THE NARROW BAND SELECTION
The emission line sample was chosen from a set of
narrow band images obtained with the SuprimeCam
camera on the Subaru 8.2-m telescope. The data were
obtained in a number of runs between 2001 and 2005
under photometric or near photometric conditions. We
used two ∼120 Å (FWHM) filters centered at nominal
5000 6000 7000 8000 9000 10000
WAVELENGTH (A)
Fig. 1.— Schematic illustration of the selection process and a
typical spectrum of the galaxies we find. The objects are chosen
based on their excess light in one of two narrow band filters at 8160
Å and 9140Å. The present case corresponds to an Hα emission
line object found in the 9140Å filter (shown with the narrow solid
curve). Also illustrated are the broad band V (dash-dot), R (solid),
I (dashed), and z′ (dotted) filters use to measure the continuum.
The spectrum shown corresponds to object 205 in Table 3 and is an
Hα emitter at z = 0.3983. The easily visible lines are the Balmer
series and the [O III] lines at λλ5007, 4959, and 4363Å.
wavelengths of 8150 Å and 9140 Å in regions of low
sky background between the OH bands. The nominal
specifications for the Subaru filters may be found at
http://www.naoj.org/Observing/Instruments/SCam/sensitivity.html
and are also described in Ajiki et al. (2003). We shall
refer to these filters as NB816 and NB912.
About 5 hour exposures were obtained with NB816 and
∼10 hour exposures with NB912 yielding 5 sigma limits
fainter than 26 mags in both bands. Deep exposures in
B, V , R, I and z′ were also taken for the fields. The data
were taken as a sequence of dithered background-limited
exposures and successive mosaic sequences were rotated
by 90 degrees. Only the central uniformly covered areas
of the images were used. Corresponding continuum expo-
sures were always obtained in the same observing run as
the narrowband exposures to avoid false identifications
of transients such as high-z supernovae, or Kuiper belt
objects, as emission-line candidates. A detailed descrip-
tion of the full reduction procedure for images is given in
Capak et al. (2004). All magnitudes are given in the AB
system (Oke 1990). These were measured in 3′′ diameter
apertures, and had average aperture corrections applied
to give total magnitudes.
The primary purpose of the program was to study Lyα
emitters at redshifts of z ∼ 5.7 and z ∼ 6.6 (Hu et al.
2004, 2007) but the narrow band imaging is also ideal
for selecting lower redshift emission line galaxies and it
is for this purpose that we use these data in the present
paper. The fields which we use and the area covered (ap-
proximately a half square degree in each bandpass) are
summarized in Table 1. These are distributed over the
sky to deal with cosmic variance. We selected galaxies in
the narrowband NB816 filter using the Cousins I band
filter as a reference continuum bandpass and including
all galaxies with NB816 < 25 and I − NB816 greater
than 0.8. We selected galaxies in the NB912 filter with
the z filter as the reference continuum bandpass and in-
http://www.naoj.org/Observing/Instruments/ SCam/sensitivity.html
THE YOUNGEST GALAXIES 3
TABLE 1
Narrowband Survey Area Coverage
Field RA (J2000) Dec (J2000) (lII, bII) EB−V
a NB816 NB912
(arcmin2) (arcmin2)
SSA22 22:17:57.00 +00:14:54.5 ( 63.1,−44.1) 0.07 674 591
SSA22 new 22:18:24.67 +00:36:53.4 ( 63.6,−43.9) 0.06 278 278
A370 new 02:41:16.27 −01:34:25.1 (173.4,−53.3) 0.03 278 278
HDF-N 12:36:49.57 +62:12:54.0 (125.9,+54.8) 0.01 710 528
Total 1940 1675
Note. — An adjacent field to A370 new is a site of a gravitational lensing cluster at
z ∼ 0.375, and was omitted from the suvey.
a estimated using http://irsa.ipac.caltech.edu/applications/DUST/ based on
Schlegel et al. (1998)
20 21 22 23 24
N(AB)
20 21 22 23 24
N(AB)
WITH SPECTRA
Fig. 2.— Continuum – Narrow band magnitude versus narrow
band magitude for all objects with narrow band magnitude brighter
than 24. The diamonds show the narrowband excess emission mag-
nitude of the NB912 sample and the squares the NB816 sample.
Galaxies which would be included in an R < 24 continuum selected
sample are shown with solid symbols. The upper panel shows the
complete sample while the lower panel shows the subsample which
has been spectroscopically identified.
cluded all galaxies with NB912 < 25 and z − NB912
greater than 1. The selection process is illustrated for
a galaxy found in the NB912 filter in Figure 1. Both
selections correspond roughly to choosing objects with
emission lines with rest frame equivalent widths greater
than 100 Å. The exact equivalent width limit depends
on the precise position of the emission line in the filter
and the redshift of the galaxy which in turn depends on
which emission line is producing the excess light in the
narrow band.
The final USEL sample consists of 542 galaxies (267
in the NB816 filter and 275 in the NB912 filter). Tabu-
lated coordinates, multi-color magnitudes, and redshifts
(where measured) for these objects are summarized in
Table 2. Very few of these objects would be included
in continuum-selected spectroscopic samples. Figure 2
shows the narrow band excess as a function of narrow-
band (NAB) magnitude for objects with narrow band
magnitudes brighter than 24. The open symbols show
the present sample while the solid symbols show ob-
jects which would be included in an R < 24 continuum-
selected sample.
3. SPECTRA
Spectroscopic observations were obtained for 161
USELs from the sample using the Deep Extragalac-
tic Imaging Multi-Object Spectrograph (DEIMOS;
Faber et al. 2003) on Keck II in a series of runs between
2003 and 2006. The emission line objects were included
in masks designed to observe high-z Lyα candidates and,
as can be seen in the lower panel of Figure 2, constitute
an essentially random sample of the emission line galax-
The observations were primarily made with the G830
ℓ/mm grating blazed at 8640 Å and used 1′′ wide slitlets.
In this configuration, the resolution is 3.3 Å, which is
sufficient to distinguish the [O II]λ3727 doublet struc-
ture. This allows us to easily identify [O II]λ3727 emit-
ters where often the [O II]λ3727 doublet is the only emis-
sion feature. The spectra cover a wavelength range of
approximately 4000 Å and were centered at an average
wavelength of 7800 Å, though the exact wavelength range
for each spectrum depends on the slit position with re-
spect to the center of the mask along the dispersion di-
rection. The G830 grating used with the OG550 blocker
gives a throughput greater than 20% for most of this
range, and ∼ 28% at 8150 Å. The observations were
not generally taken at the parallactic angle, since this
was determined by the mask orientation, so considerable
care must be taken in measuring line fluxes as we dis-
cuss below. Each ∼ 1 hr exposure was broken into three
subsets, with the objects stepped along the slit by 1.5′′
in each direction. Some USELs were observed multiple
times, resulting in total exposure times for these galaxies
of 2−3 hours. The two-dimensional spectra were reduced
following the procedure described in Cowie et al. (1996)
and the final one-dimensional spectra were extracted us-
ing a profile weighting based on the strongest emission
http://irsa.ipac.caltech.edu/applications/DUST/
4 Kakazu et al.
4000 4500 5000 5500 6000 6500 7000
REST WAVELENGTH (A)
Fig. 3.— Spectrum of an Hα emission galaxy selected in the NB912 filter. In the upper plot we have decreased the scale of the vertical
axis by a factor of 10 to show the continuum and the weaker lines. The more important emission line features are labelled and marked
with the dotted lines.
3500 4000 4500 5000 5500
REST WAVELENGTH (A)
Fig. 4.— Spectrum of an [O III] galaxy in the NB816 selected sample. The lower plot shows the relative strengths of the very strong
emission lines in the spectrum. In the upper plot we have decreased the scale of the vertical axis by a factor of 10 to show the continuum
and the weaker lines. The more important emission line features are labelled and marked with the dotted lines.
THE YOUNGEST GALAXIES 5
3500 4000 4500 5000 5500
REST WAVELENGTH (A)
Fig. 5.— Spectrum of an [O III] galaxy selected in the NB912 filter. The lower plot shows the relative strengths of the very strong
emission lines in the spectrum. In the upper plot we have decreased the scale of the vertical axis by a factor of 10 to show the continuum
and the weaker lines. The more important emission line features are labelled and marked with the dotted lines.
3000 3200 3400 3600 3800 4000 4200 4400
REST WAVELENGTH (A)
Fig. 6.— Spectrum of an [O II] galaxy selected in the NB816 filter. The plot shows the relative strengths of the very strong emission
lines in the spectrum. The more important emission line features are labelled and marked with the dotted lines.
6 Kakazu et al.
0.0 0.5 1.0 1.5
REDSHIFT
[OIII]
Hα [OII]
0.0 0.5 1.0 1.5
REDSHIFT
10-18
10-17
10-16
10-15
Fig. 7.— (a) Distribution of redshifts for the spectroscopically
identified sources. [O III] λ5007 emitters are the most common.
Since the focus of this paper is on intermediate-redshift (z . 1)
strong emission line galaxies, we did not plot high redshift Lyα
galaxies (z >> 5) in our NEO sample. High-z Lyα emitters are
discussed in Hu et al. (2004, 2007). (b) Flux versus redhift for the
spectroscopically identified sample. Squares are Hα, diamonds are
[O III] λ5007, and triangles are [O II] λ3727. The solid line shows
the flux limit corresponding to the narrow band magnitude limit of
N(AB)=25 for an emitter with very large equivalent width. Some
objects with lower equivalent widths fall below this limit.
line in the spectrum. A small number of the spectra
were obtained with the ZD600 ℓ/mm grating giving a
correspondingly lower resolution but a wider wavelength
coverage. These observations were centered at 7200 Å.
Essentially all of the emission line candidates which
were observed were identified, though two of the objects
in the NB816 sample are stars where the absorption line
structure mimics emission in the band. Sample spec-
tra are shown in Figures 3, 4, 5, and 6. The measured
redshifts are given in Tables 2 and 3. The narrow band
emission line selection produces a mixture of objects cor-
responding to Hα, [O III]λ5007, and [O II]λ3727 and, at
the faintest magnitudes (> 24), high redshift Lyα emit-
ters. The number of objects seen in each line and the
redshifts where they are found are shown in Figure 7.
The spectroscopically identified sample from both bands
contains 13 Hα, 92 [O III]λ5007, and 23 [O II]λ3727 emit-
ters. In the remainder of the paper we shall focus primar-
ily on the Hα and [O III]λ5007 selected galaxies which lie
between redshifts zero and one.
Since only 30% of the USELs are spectroscopically
identified we must apply a substantial incompleteness
correction in computing the line luminosity function
and the universal star formation histories. Because
the type mix may vary as a function of magnitude we
have adopted a magnitude dependent weighting for each
galaxy equal to the total number of galaxies at this mag-
nitude divided by the number of spectroscopically iden-
tified galaxies. However, the analysis is not particularly
sensitive to the adopted scheme since the fraction of iden-
tified galaxies is relatively constant with magnitude.
4. FLUX CALIBRATIONS
Generally our spectra were not obtained at the par-
allactic angle since this is determined by the DEIMOS
mask orientation. Therefore flux calibration using stan-
dard stars is problematic due to atmospheric refraction
effects, and special care must be taken for the flux cali-
bration. We thus employed three independent methods
for the flux calibration. In §4.1 we introduce “primary
fluxes” which are absolute fluxes of the emission lines
used to select the galaxies. Primary fluxes are computed
from the SuprimeCam broadband and narrowband mag-
nitudes. We use these fluxes to derive luminosity func-
tions of Hα and [O III]λ5007 emitters at z < 1 (§5.1).
In §4.2 we measure line fluxes from the spectra. Rela-
tive line fluxes can be measured from the spectra without
flux calibration as long as we restrict the line measur-
ments to over short wavelength range where the DEIMOS
response is essentially constant. For example, one can as-
sume the response of neighboring lines (e.g. [O III]λ4949
and [O III]λ5007) are the same and therefore one can
measure the flux ratio without calibration. For bright
galaxies, we can absolutely calibrate the fluxes by inte-
grating spectra and equating them to Subaru broadband
fluxes. These line fluxes derived from the spectra are
used as a check of the primary fluxes. We show that
the ratio of [O III]λ5007/[O III]λ4959 is indeed close to
1/3, and that the fluxes computed from the spectra are
highly consistent with the primary fluxes measured in
§4.1. In §4.3, we show Balmer flux ratios f(Hβ)/f(Hα) of
bright Hα emitters are close to the Case B conditions,
suggesting very little reddening.
Metallicity measurements by the direct method re-
quire four emission lines that are widely displaced
over the spectral wavelength range ([O III]λλ4959, 5007,
[O III]λ4363, and [O II]λ3727). To calibrate these lines,
we used neighboring Balmer lines with the assumption
of Case B conditions, and this is described in §4.4.
4.1. Narrow Band Fluxes − Primary Fluxes
For the emission lines used to select each galaxy we
may compute the equivalent widths and absolute fluxes
directly from the narrow band magnitudes (N) and
the corresponding continuum magnitudes (C) from our
SuprimeCam imaging data. For example, in the case of
the NB816 selected emission-line galaxies, N corresponds
to the NB816 magnitude and C is the I band magnitude.
We shall refer to the values calculated in this way as the
primary fluxes and use this quantity to compute the lu-
minosity functions for each emitter in §5.1.
Defining the quantity
R = 10−0.4∗(N−C)
THE YOUNGEST GALAXIES 7
1.5 2.0 2.5 3.0
LOG REST FRAME EW ([OIII]5007)
1.5 2.0 2.5 3.0
LOG REST FRAME EW (Hα)
Fig. 8.— (a) Distribution of the rest frame equivalent widths
determined from the narrow band magnitudes for the spectroscop-
ically identified [O III] λ5007 sources. (b) Distribution of the rest
frame equivalent widths for the Hα selected sample.
the observed frame equivalent width becomes
where φ is the narrow band filter response normalized
such that the integral over wavelength is unity and ∆λ is
the effective width of the continuum filter. The narrow
band filter is often assumed to be rectangular in which
case φ becomes 1/δλ where δλ is the width of the narrow
band but as can be seen from Figure 1 this is not a very
good approximation in the present case. For very high
equivalent width objects the denominator in this equa-
tion becomes uncertain unless the broad band data are
very deep, and this can result in a large scatter in the
very highest equivalent widths where the continuum is
poorly determined.
In the case of the [O III]λ5007 line we must include the
second member of the doublet which also lies within the
narrow band filter. We have computed these cases as-
suming the flux of the [O III]λ4959 line is 0.34 times that
of the [O III]λ5007 line. Then φ = φ1+0.34×φ2 where φ1
is the filter response at the redshifted 5007 Å wavelength
and φ2 is the filter response at redshifted 4959 Å.
The distribution of the rest frame equivalent widths for
the Hα and [O III]λ5007 samples is shown in Figure 8.
The [O III]λ5007 sample selects objects with rest frame
equivalent widths above about 100Å while the lower red-
shift Hα sample selects objects with rest frame equivalent
widths above about 150Å. Since the [O III]λ5007 lines are
also generally stronger than the Hα lines the [O III] selec-
tion chooses less extreme objects than the Hα selection
and will include a larger fraction of galaxies at the given
10-18 10-17 10-16 10-15
f(OIII 5007) erg cm-2 s-1
Fig. 9.— The ratio of the [O III] λ4959 line to [O III] λ5007. The
errors are plus and minus 1 sigma. The median ratio is 0.338 and
the scatter around this value is consistent with that expected from
the statistical errors.
redshift.
The high observed frame equivalent widths make the
line fluxes insensitive to the continuum determination
and these may simply be found from
f = A
10−0.4N − 10−0.4C
where A is the AB zeropoint at the narrow band wave-
length in units of erg cm−2 s−1 Å−1. The flux depends
on the filter response at the emission line wavelength and
correspondingly is most uncertain at the edges of the fil-
ters where this quantity changes rapidly. Primary fluxes
defined here are measured by using narrowband (N) and
broadband (C) magnitudes from Subaru imaging data
with the object redshift information from Keck spectra
for the filter response at the exact location of emission
line wavelength (φ). We use these primary fluxes to con-
struct the luminosity functions of Hα and [O III]λ5007
selected emitters as we discuss in §5.1.
4.2. Line Fluxes from the Spectra
For the short wavelength range where DEIMOS re-
sponse is essentially constant, we may also compute the
relative line fluxes from the spectra without calibration.
For each spectrum we fitted a standard set of lines. For
the stronger lines we used a full Gaussian fit together
with a linear fit to the continuum baseline. For weaker
lines we held the full width constant using the value mea-
sured in the stronger lines and set the central wavelength
to the nominal redshifted value. We also measured the
noise as a function of wavelength by fitting to random
positions in the spectrum and computing the dispersion
in the results.
These fits should provide accurate relative fluxes over
short wavelength intervals where the DEIMOS response
is similar, but may be expected to be poorer over longer
wavelength intervals where the true calibration can vary
from the adopted value. We tested the relative flux cali-
bration for neighboring lines and the noise measurement
by measuring the ratio of the [O III]λ4959/ [O III]λ5007
lines. This is expected to have a value of approximately
0.34. The ratio is shown as a function of the [O III]λ5007
8 Kakazu et al.
10-17 10-16 10-15
FLUX (erg cm-2 s-1)
Fig. 10.— Ratio of fluxes computed from the spectra and the
broad band magnitudes versus those from the narrow band mag-
nitudes. Hα lines are shown as solid boxes, [O III] λ5007 lines as
diamonds and [O II] λ3727 lines as crosses.
flux in Figure 9. The measured values scatter around the
expected value and the dipsersion is consistent with the
noise determination. This result supports our assump-
tion of [O III]λ4959/[O III]λ5007 = 0.34 in the primary
fluxes measurements described in §4.1.
The brighter objects may be absolutely calibrated us-
ing the continuum magnitudes obtained from our Sub-
aru data. We integrated the spectrum convolved with
each SuprimeCam filter response and set this equal to
the broad band flux to normalize the spectrum in each
of the filters. We then used the Gaussian fits to obtain
the spectral line fluxes for lines lying within that broad
band. This procedure only works for sources with well
determined continuum magnitudes (C < 24.5) where the
sky subtraction can be well determined in the spectra.
For these objects the spectrally determined fluxes are
shown versus the primary fluxes in Figure 10 where we
plot the ratio of the spectral to the primary flux versus
the primary line flux. The values agree extremely well
though the measured spectral line fluxes are on average
about 80 − 90% of the primary flux values. This may
reflect a selection bias in the choice of the objects or the
narrow band filters could be slightly narrower than the
nominal profiles.
4.3. Balmer Ratios
We now measured the Balmer ratios for the sample of
objects selected in Hα and where the continuum magni-
tudes were bright enough to absolutely flux calibrate the
spectra. The ratio of f(Hβ)/f(Hα) is shown in Figure 11.
The values average closely to the Case B ratio which is
shown as the solid line and at brighter fluxes the indi-
vidual values also closely match to this value suggesting
that the galaxies have very little reddening. However,
at fainter fluxes the scatter about the average value is
considerably higher than the statistical errors. At the
faintest fluxes it appears that the systematic uncertainty
can be as high as a multiplicative factor of two.
4.4. Final Flux Calibration for Metallicity Analysis
For the metallicity analysis we adopted the pro-
cedure of normalizing the longer wavelength lines
([O III]λλ4959, 5007, [O III]λ4363) to their nearest
10-18 10-17 10-16 10-15
f(Hα) erg cm-2 s-1
Fig. 11.— The ratio of the Hβ/Hα fluxes versus Hα flux. The
values average to the unreddened Balmer decrement shown by the
solid line but at the lower fluxes the scatter is larger than expected
from the statistical errors reflecting the calibration uncertainties for
the fainter sources. The figure shows the ten objects detected in the
Hα line with continuum magnitudes above 24.5 in the bandpasses
corresponding to the lines.
Balmer line to determine the unreddened fluxes. For ex-
ample, in the case of the Hα emission selected galaxies,
we can measure Hα absolute fluxes by the primary fluxes
method described in §4.1. We can then derive Hβ and Hγ
fluxes from Hα fluxes by assuming Case B recombination
[e.g., f(Hα) = 2.85 × f(Hβ), f(Hγ) = 0.469 × f(Hβ) at T
= 104 K and Ne ∼ 10
2−104cm−3; Osterbrock 1989]. As
Hβ and [O III]λλ4959, 5007 have very similar DEIMOS
response, the relative fluxes should remain the same with
or without the flux calibration and this can be expressed
by a simple equation:
f0(Hβ)
f0([O III]λ4959, λ5007)
f(Hβ)
f([O III]λ4959, λ5007)
where f0(Hβ) and f0([O III]λ4959, λ5007) are the flux
counts in the un-calibrated, reddened DEIMOS spec-
tra, while f(Hβ) and f([O III]λ4959, λ5007) are absolute,
unreddened fluxes. Since we know f(Hβ) from f(Hα)
with the Case B assumption and f0(Hβ)/f0([O III]λ4959,
λ5007) from the DEIMOS spectra, we can derive
f([O III]λ 4959, λ5007) using this simple formula. Sim-
ilary, we can absolutely calibrate [O III]λ4363 lines by
using its neighboring Balmer line, Hγ:
f0(Hγ)
f0([O III]λ4363)
f(Hγ)
f([O III]λ4363)
where f0(Hγ) and f0([O III]λ4363) are again the counts
in flux uncalibrated, reddened DEIMOS spectra, and
f(Hγ) and f([O III]λ4363) are absolute fluxes.
In the case of the [O III] selected emitters, we first de-
rive [O III]λλ4959, 5007 absolute fluxes using the primary
fluxes method (§4.1), and then use the above formula to
get absolute fluxes of Hβ, then Hγ (by the Case B ratio),
and finally [O III]λ4363.
This flux calibration technique using neighboring
Balmer line should work well for the [O III]λλ4959, 5007,
λ4363 lines and the [N II] lines which all lie close to
Balmer lines but may be slightly more approximate for
the [S II] lines. The higher order Balmer lines are too un-
certain to apply this procedure due to inadequate S/N of
THE YOUNGEST GALAXIES 9
40 41 42 43
L(Hα) (erg s-1)
1 z=0.24
40 41 42 43
L(Hα) (erg s-1)
1 z=0.39
Fig. 12.— The luminosity function of Hα at z = 0.24 (top panel)
and at z = 0.39 (bottom panel). In each case the open boxes
show the luminosity functions determined from the spectroscopic
sample alone while the solid boxes show the function corrected for
the incompleteness in the spectroscopic identification. The errors
are plus and minus 1 sigma and at the highest luminosity we show
the 1 sigma upper limit.
the lines, and we have used the continuum flux calibrated
values with no reddening for the [O II]λ3727 and [Ne III]
lines. These values will have correspondingly higher flux
uncertainties. Fortunately the [O II]λ3727 line is very
weak in most of the objects and the uncertainty has lit-
tle effect on the metallicity determinations. However,
ionization analyses based on the [Ne III] line should be
undertaken with caution.
5. STAR FORMATION HISTORY
5.1. Hα and [O III]λ5007 Luminosity Functions
Because of the high observed frame equivalent widths
the primary fluxes are insensitive to the continuum de-
termination. However, they do depend on the filter re-
sponse at the emission line wavelength so we first restrict
ourselves to redshifts where the nominal filter response
is greater than 70% of the peak value. This also has
the advantage of providing a uniform selection and we
assume the window function is flat over the defined red-
shift range. Now the volume is simply defined by the se-
lected redshift range for all objects above the minimum
luminosity which we take as corresponding to a flux of
1.5×10−17 erg cm−2 s−1 (Figure 7). The luminosity func-
tion is now obtained by dividing the number of objects
40 41 42 43
L(OIII) (erg s-1)
1 z=0.63
40 41 42 43
L(OIII) (erg s-1)
1 z=0.83
Fig. 13.— The luminosity function of [O III] λ5007 emitters at
z = 0.63 (top panel) and at z = 0.83 (bottom panel). In each case
the open boxes show the luminosity functions determined from the
spectroscopic sample alone while the solid boxes show the function
corrected for the incompleteness in the spectroscopic identification.
The errors are plus and minus 1 sigma and at the highest luminosity
we show the 1 sigma upper limit.
in each luminosity bin by the volume. The incomplete-
ness corrected luminosity function is obtained from the
sum of the weights in each luminosity bin divided by the
volume. The 1 sigma errors shown are calculated from
the Poissonian errors based on the number of objects in
the bin. The calculated Hα luminosity function is shown
for the two redshift ranges corresponding to the NB816
and NB912 selections in Figure 12 and the corresponding
[O III]λ5007 luminosity functions in Figure 13.
5.2. Star Formation Rates
The individual objects have Hα luminosities stretching
up to about 3×1041 erg s−1 where, at the higher redshifts,
we use the Hβ luminosity to infer the Hα value assuming
there is no reddening. For a steady formation this would
require a star formation rate of a few solar masses per
year if we adopt the Kennicutt (1998) conversion rate for
his Salpeter mass function.
Since the objects are more probably caused by star-
bursts the true star formation rate will depend on the
evolutionary history. However, the Hα luminosity den-
sity should give a reasonable estimate of the universal
star formation density of the objects provided only that
most of the ionizing photons are absorbed in the galax-
10 Kakazu et al.
Fig. 14.— The star formation history inferred from the Hα or
Hβ luminosity density as a function of redshift. The data from our
sample are shown in red. The open squares show the total rate
from the entire sample while the solid squares show the values for
objects with Hα rest frame equivalent widths in excess of 200Å or
Hβ rest frame equivalent widths in excess of 70Å. The diamonds
show the UV star formation rates (uncorrected for extinction) from
the ground based work of Wilson et al. (2002) and the triangles
the Galex results of Wyder et al. (2005) and Schiminovich et al.
(2005). Hα measurements from the literature as summarized in
Ly et al. (2007) are shown with filled circles. In all cases the errors
are ±1σ.
ies. We first formed the total Hα or Hβ luminosity den-
sity of the galaxies by summing over the incompleteness
weighted luminosities in each redshift interval. We only
included detected objects and did not attempt to extrap-
olate to lower luminosities but the result are not partic-
ularly sensitive to this because the luminosity functions
are relatively flat at the lower luminosities (Figures 13
and 14). We then converted these to star formation rates
with the Kennicutt (1998) conversion.
The results are shown in Figure 14. We first plot
the rate for the total samples at each redshifts shown
by the open squares. We have shown star formation
rates for UV continuum samples for comparison and the
present sample of strong emitters gives star formation
rates which are about 10% of the UV values over the
redshift interval. For comparison, we also show the star
formation rates from Hα selected samples reported in the
literature and summarized in Ly et al. (2007). In order
to better understand the evolution we have also restricted
our own sample to objects with rest frame equivalent
widths of Hα in excess of 200Å at low redshifts and Hβ
equivalent widths in excess of 70Å at the higher redshifts.
The star formation rates for this sample are shown with
the solid squares. This provides a more uniform selection
with redshift and gives a slower increase than the total
inferred star formation rate. For this sample the SFR is
evolving roughly as (1+z)3 broadly similar to other UV
and optically determined formation rates in this redshift
interval. These more restricted objects comprise about
5% relative to the UV star formation rates at the higher
redshifts.
6. GALAXY METALLICITIES
6.1. [O III] emitters
The spectra are of variable quality and, in order to
measure the metallicities, we need very high signal to
noise observations. It is also important that Balmer lines
are well detected since our flux calibrations are done us-
ing the neighboring Balmer lines (§4.4). We therefore
restricted ourselves to [O III] emitters whose Hβ fluxes
are detected above 15 sigma. Among 92 [O III] emitters
in our total spectroscopic sample, 8 such [O III] emitters
were chosen in the NB912 sample, and 10 in the NB816.
These emitters have Hγ detected above 4 sigma. Tables 4
and 5 give the oxygen line fluxes normalized by their Hβ
fluxes for the NB816 and NB912 selected emitters, re-
spectively. 1σ upper limits are listed when the measured
flux is below 1σ.
The most direct method to estimate the gas-phase
oxygen abundance is to use the electron temperature of
the HII region. Higher gas metallicity increases nebular
cooling, leading to lower electron temperatures. There-
fore electron temperature is a good indicator of the
gas metallicity. The electron temperature can be de-
rived from the ratio of the [O III]λ4363 auroral line to
[O III]λλ5007,4959. This procedure is often referred to
as the ‘direct’ method or Te method (e.g., Seaton 1975;
Pagel et al. 1992; Pilyugin & Thuan. 2005; Izotov et al.
2006c). One well-known problem with the direct method,
however, is that [O III]λ4363 is generally very weak even
in the low-metallicity galaxies. For higher metallicity
systems, the far-IR lines become the dominant coolant
and therefore the optical auroral line is essentially not
detectable. However, the majority of our sample exhibit
[O III]λ4363, already suggesting that these are metal-
deficient systems. To derive Te[O III] and the oxygen
abundances, we used the Pagel et al. (1992) calibrations
with the Te[O II]−Te[O III] relations derived by Garnett
(1992). The results are shown in Table 4 (for NB816
selected [O III] emitters) and Table 5 (for NB912 se-
lected [O III] emitters). The Izotov et al. (2006c) for-
mula, which was developed with the latest atomic data
and photoionization models, gives consistent abundances
within 0.1 dex. The [S II]λλ6717, 6731 lines that are usu-
ally used for the determination of the electron number
density, are beyond the Keck/DEIMOS wavelength cov-
erage for our [O III] emitters. Therefore we assumed ne
= 100 cm−3. However the choice of electron density has
little effect as electron temperature is insensitive to the
electron density; we get the same results even when we
use ne = 1000 cm
The 1σ upper and lower limits on Te[O III] and the
oxygen abundances are also shown in the tables. Because
the [O III]λ4363 flux is weak (< 4σ), the range of our
metallicity measurements are quite wide. However, of
18 [O III] emitters, even the upper metallicity limits on 6
emitters satisfy the definition of XMPGs [12 + log(O/H)
< 7.65; Kunth & Östlin 2000]. All our emitters, except
the ones that only have lower limits on metallicties such
as ID31 in Table 4, have very low metallicities – even the
upper metallicity limits are about 0.02 - 0.2 Z⊙. A few
emitters may even have metallicities that are comparable
to the currently known most metal-poor galaxies [I Zw
18 and SBS0335−052W; 12 + log(O/H) ∼ 7.1 − 7.2].
However our current metallicity errors are too large to
measure the baseline metallicity accurately and higher
S/N spectra will be necessary for this purpose.
Our discovery rate of XMPGs appers to be signif-
icantly higher than those of other surveys. Only 14
new XMPGs have been discovered from the analysis of
THE YOUNGEST GALAXIES 11
Fig. 15.— [O III]λ4959+λ5007/[O III]λ4363 versus
[O II]λ3727/[O III]λ5007 for the [O III] and Hα selected emitters
in Table 4 and 5. The electron temperature of the HII region is
also shown.
∼530,000 galaxy spectra in the SDSS and they are all
located nearby (z < 0.005) (SDSS DR3: Kniazev et al.
2003, SDSS DR4: Izotov et al. 2006a). At higher red-
shift, 17 metal-poor (7.8 < 12 + log(O/H)< 8.3) galax-
ies have been found at z ∼ 0.7 in the initial phase of the
DEEP2 survey of 3,900 galaxies and the Team Keck Red-
shift Survey of 1,536 galaxies (Hoyos et al. 2005). But
none of these galaxies satisfies the condition of XMPGs.
The present sample may be the first XMPGs at inter-
mediate redshift (z ∼ 1) whose oxygen abundances are
securely measured by the direct method. The narrow-
band method is very powerful for finding not only high-
redshift (z >> 5) galaxies, but also strong emission-line,
extremely metal-deficient galaxies at intermediate red-
shifts (z < 1).
Figure 15 shows the electron temperature sensi-
tive line ratio, [O III](λ4959+λ5007)/[OIII]λ4363 ver-
sus [O II]λ3727/[O III]λ5007. If we have an estimate of
the metallicity, as in the present case, we can use the
[O II]λ3727/[O III]λ5007 ratio to estimate the ionization
parameter q. The ionization parameter q is defined as
the number of hydrogen ionizing photons passing through
a unit area per second per unit hydrogen number den-
sity (Kewley & Dopita 2002). For the low metallicity
systems with strong [O III]λ4363 lines, we can see from
Figure 15 that [O II]λ3727 is very weak compared to
[O III]λ5007 with values ranging downwards from 0.3.
Assuming the metallicity is less than 0.2 Z⊙ this would
place a lower bound of q = 108 cm s−1 on the ioniza-
tion parameter based on the Kewley & Dopita (2002)
model. The higher metallicity objects have stronger
[O II]λ3727/[O III]λ5007 which, while in part due to the
metallicity, also requires these objects to have lower ion-
ization parameters suggesting we are seeing an evolution-
ary sequence.
6.2. Hα emitters
Among 13 Hα emitters in our spectroscopic sample,
only 3 NB912 selected emitters have Hβ fluxes ade-
quate (> 15σ) for the purpose of metallicity measure-
ments. Their Te[O III] and oxygen abundances were mea-
sured using the direct method described above, and are
shown in the Table 5 together with the data for the
[O III] emitters. The [O II]λ3727 line of ID266 is out-
side the Keck/DEIMOS wavelength coverage. In order
to derive an upper-limit on the metallicity, we assumed
[O II]λ3727/[O III]λ5007 = 1.0, which is the maximum
value in our sample (see, Fig. 15). All our Hα emitters
are metal poor (Zupper < 0.45 Z⊙), but none of them are
XMPGs.
6.3. Composite Spectrum
As can be seen in Figure 15 the objects with low
[O II]λ3727/[O III]λ5007 have relatively uniform values
of ([O III]λ5007+λ4959)/[OIII]λ4363 and similar metal-
licities. Given the relatively low signal to noise of the
individual spectra it therefore seems of interest to form a
composite spectrum. Such a spectrum will have weight-
ings on the lines which depend on the individual ion-
ization parameters and metallicity but will give a rough
estimate of the average metallicity and temperature of
the sample.
In Figure 16 we show the composite spectrum of
the 8 objects with [O II]λ3727/[OIII]λ5007 less than
0.1. The [O III]λ4363 is now strongly detected with a
value of 16.7 ± 2.1 or eight sigma. The mean temper-
ature is 19, 500 ± 1, 500 K and the average abundance
12+log(O/H)=7.5± 0.1 and the mean rest frame equiv-
alent width of Hβ is 57Å. The results are similar to the
values obtained by averaging the individual analyses of
the eight objects.
7. MORPHOLOGIES
The morphology of the USELs may give us a clue to the
mechanism of their high star formation activity (SFR ∼
a few M⊙/yr) and star formation history; what has trig-
gered the star formation −mergers, gas infall, or galactic
winds? High resolution HST/ACS images are available
for GOODS-North (GOODS-N) region (Giavalisco et al.
2004) which is one of our survey fields. There are 17
NB816 selected USELs in the GOODS-N, and 16 in the
NB912 sample. Figures 17 and 18 show thumbnails of
the NB816 and NB912 selected USELs in the GOODS-N
field (respectively) with each thumbnail 12.′′5 on a side.
The white dashes point to the largest galaxy. We used
continuum broadband images to show underlying stellar
populations: HST/ACS B, V, z′-band images were used
for NB816 emitters, and B, V, I-band for NB912 emit-
ters. High-redshift Lyα emitters (z >> 5) are very red
because of the continuum absorption below Lyα emission
caused by neutral hydrogen in the intergalactic medium.
We do not have spectra for most of the USELs in the
GOODS-N field yet, and none of the USELs in GOODS-
N have metallicity measurements either due to lack of
spectra or low spectral S/N. But we can qualitatively
argue that the USELs at intermediate redshift (z < 1)
exhibit widespread morphologies from relatively compact
high surface brightness objects to very diffuse low surface
brightness ones.
8. DISCUSSION
The present emitters differ from the local dwarf HII
galaxies in a large number of ways though they appear
much more similar to the XMPGs found in the SDSS
samples. They are much more luminous, have large
[O III]/[O II] ratios, and they are a relatively high frac-
tion (about 10% by number from Figure 13) of other
12 Kakazu et al.
3600 3800 4000 4200 4400 4600 4800 5000 5200
REST WAVELENGTH (A)
Fig. 16.— Composite spectrum of the 8 emitters with [O II]λ3727/[O III]λ5007 less than 0.1. The eight spectra have simply been summed
without weighting. The lower panel shows the stronger lines and the upper the continuum and the weaker lines. The stronger emission
lines are labelled and marked with the vertical dotted lines.
Fig. 19.— The oxygen abundance versus the absolute rest frame B magnitude for the [O III] selected samples (red squares). One sigma
errors are shown for the oxygen abundances and one sigma lower limits are shown with upward pointing arrows. The solid line shows the
(Skillman et al. 1989) relation for the nearby dwarf irregulars. As with the local XMPGs (filled circles, Kniazev et al. 2003; Kewley et al.
2007 and GRB hosts (open squares, Stanek et al. 2006; Kewley et al. 2007, the present galaxies are much more luminous at a given
metallicity than the local irregulars. Metal-poor luminous galaxies (but not XMPGs) at z ∼ 1 from Hoyos et al. 2005. are shown as
triangles. A few of our emitters may have oxygen abundances comparable to the most metal-deficient galaxis, I Zw 18 (12 + log O/H =
7.17± 0.01, Thuan & Izotov 2005) and SBS 0335-052W (12 + log O/H = 7.12 ± 0.03, Izotov et al. 2005).
THE YOUNGEST GALAXIES 13
galaxy populations at these redshifts. Taken together
this suggests that we are seeing much more massive
galaxies in the early stages of formation and, since we
need these to have relatively long lifetimes in order to
understand their frequency, that we are seeing objects
undergoing continuous star formation rather than single
starbursts. For the case of constant star formation with a
standard Salpeter IMF a forming galaxy can have equiv-
alent widths above 30Å for 109 yr (Leitherer et al. 1999)
which would allows us to understand the observed num-
ber density of strong emitters relative to the total galaxy
population.
In this type of model we would expect the metal-
licity to grow with time and that higher metallicity
galaxies would have higher continuum magnitudes and
lower equivalent widths in Hβ. We plot the absolute
rest frame B magnitudes versus the Oxygen abundance
in Figure 19. As with the case for the local XMPGs
found in the SDSS (filled circles, Kniazev et al. 2003;
Kewley et al. 2007) and the metal-poor galaxies (7.8 <
12 + log O/H < 8.3) at z ∼ 1 (triangles, Hoyos et al.
2005), the present emitters (red squares) are much more
luminous at a given metallicity than is found for the
local dIrrs (solid line, Skillman et al. 1989). Further-
more there does indeed seem to be a trend to higher
continuum luminosities at higher metallicity consistent
with ongoing star formation raising the luminosity. Re-
cently Kewley et al. (2007) reported the similarity be-
tween XMPGs and long duration GRB hosts ; they share
similar SFRs, extinction levels, and both lie in a similar
region of the luminosity-metallicity diagram. Our sam-
ple metal-deficient galaxies, which also lie in the region of
GRB hosts, may be additional support of the connection
between XMPGs and GRB hosts.
Of the six galaxies with continuum magnitudes
brighter than -18 all but one have metallicities or
lower limits which would place them near or above
12+log(O/H)=8 while the lower luminosity galaxies pri-
marily have 12+log(O/H) in the range 7.1 − 7.8. If we
assumed that the metallicity were a simple linear func-
tion of the age then the more luminous galaxies would be
several times older than the less luminous ones which is
not quite enough to account for the luminosity increase
(see e.g. Leitherer et al. 1999) suggesting that the en-
richment process may be more complex. However, the
accuracy of our current metallicity measurements may
be inadequate for measuring the lowest metallicities in
the sample and we could be underestimating the amount
of metallicity evolution.
The relation between the metallicity and the Hβ equiv-
alent width is shown in Figure 20. There clearly is a large
scatter in metallicity at all equivalent widths suggesting
that the star formation may be episodic with periods in
which bursts of star formation enhances the Hβ equiva-
lent widths in objects where previous star formation has
raised the metallicity.
With better spectra and more accurate metal estimates
we should be able to refine these tests and also determine
whether the number density of objects versus metallic-
ity is consistent with that expected in a simple growth
model. Perhaps even more importantly as larger spectro-
scopic samples are obtained we should be able to deter-
mine if there is a floor on the metallicity corresponding
10 100 1000
EW (H Beta)
Fig. 20.— The oxygen abundance versus the rest frame Hβ equiv-
alent width for the [O III] selected samples. One sigma errors are
shown for the oxygen abundances and one sigma lower limits are
shown with upward pointing arrows. The dotted line shows the
metallicity of 1Zw-18.
to the enrichment in the intergalactic gas. Within the er-
rors we have yet to find an object with lower metallicity
than the lowest metallicity local galaxies but this could
easily change as the observations are improved.
9. SUMMARY
We have described the results of spectrscopic observa-
tions of a narrowband selected sample of extreme emis-
sion line objects. The results show that such objects are
common in the z = 0 − 1 redshift interval and produce
about 5-10% of the star formation seen in ultraviolet or
emission line measurements at these redshifts. A very
large fraction of the strong emitters are detected in the
[O III]λ4363 line and oxygen abundances can be mea-
sured using the direct method. The abundance primar-
ily lie in the 12+log(O/H) range of 7−8 characteristic of
XMPGs.
The results suggest that we are seeing early chemical
enrichment of startup galaxies at these redshifts which
are forming in relatively chemically regions. As we de-
velop larger samples of these objects and improve the
accuracy of their abundance estimates we should be able
to test this model and to determine if there is a floor on
the metallicity of the galaxies.
We are indebted to the staff of the Subaru and Keck
observatories for their excellent assistance with the obser-
vations. We acknowledge Subaru/SuprimeCam support
astronomer, Hisanori Furusawa, for his help over several
years with the observations. Y. Kakazu acknowledges
invaluable advice from Lisa J. Kewley and Roberto Ter-
levich on metallicity measurements. This work was sup-
ported in part by NSF grants AST04-07374 (LLC) and
AST06-87850 (EMH), and Spitzer grant JPL 1289080
(EMH).
14 Kakazu et al.
REFERENCES
Ajiki, M. et al. 2003. AJ, 126, 2091, astro-ph/0307325
Ajiki, M., et al. 2006, PASJ, 58, 113, astro-ph/0511471
Capak, P. et al. 2004, AJ, 127, 180, astro-ph/0312635
Corbin, M. R., Vacca, W. D., Cid Fernandes, R., Hibbard, J. E.,
Somerville, R. S., & Windhorst, R. A. 2006, ApJ, 651, 861,
astro-ph/0607280
Cowie, L. L., Songaila, A., Hu, E. M., & Cohen, J. G. 1996, AJ,
112, 839, astro-ph/9606079
Faber, S. M. et al. 2003, Proc. SPIE, 4841, 1657
Garnett, D. R. 1992, AJ, 103, 1330
Giavalisco, M. et al. 2004, ApJ, 600, L93, astro-ph/0309105
Hoyos, C., Koo, D. C., Phillips, A. C., Willmer, C. N. A., &
Guhathakurta, P. 2005, ApJ, 635, L21, astro-ph/0511066
Hu, E. M., Cowie, L. L., & McMahon, R. G. 1998, ApJ, 502, L99,
astro-ph/9801003
Hu, E. M., Cowie, L. L., Capak, P., Hayashino, T., & Komiyama,
Y. 2004, AJ, 127, 563, astro-ph/0311528
Hu, E. M., Cowie, L. L., Kakazu, Y., & Capak. P. 2007, in
preparation
Izotov, Y. I., Thuan, T. X., & Guseva, N. G., 2005, ApJ, 415, 87,
astro-ph/0506498
Izotov, Y. I., Papaderos, P., Guseva, N. G., Fricke, K. J., &
Thuan, T. X. 2006a, A&A, 454, 137, astro-ph/0604234
Izotov, Y. I. 2006b, in ASP Conf. Ser. 353, Stellar Evolution at
Low Metallicity: Mass Loss, Explosions, Cosmology, ed.
H. J. G. L. M. Lamers, N. Langer, T. Nugis, & K. Annuk (San
Francisco: ASP), 349
Izotov, Y. I., Stasińska, G., Meynet, G., Guseva, N. G., & Thuan,
T. X. 2006c, A&A, 448, 955, astro-ph/0511644
Kennicutt, R. C., Jr. 1998, ARA&A, 36, 189, astro-ph/9807187
Kewley, L. J., & Dopita, M. A. 2002, ApJS, 142, 35,
astro-ph/0206495
Kewley, L. J., Brown W. R., Geller, M. J., Kenyon, S. J., &
Kurtz, M. J. 2007, ApJ, in press, astro-ph/0609246
Kniazev, A. Y., Grebel, E. K., Hao, L., Strauss, M. A.,
Brinkmann, J., & Fukugita, M. 2003, ApJ, 593, L73,
astro-ph/0307401
Kobulnicky, H. A., & Kewley, L. J. 2004, ApJ, 617, 240,
astro-ph/0408128
Kobulnicky, H. A., et al. 2003, ApJ, 599, 1006, astro-ph/0310346
Kunth, D., & Östlin, G. 2000, A&A Rev., 10, 1, astro-ph/9911094
Leitherer, C. et al. 1999, ApJS, 123, 3, astro-ph/9902334
Lilly, S. J., Carollo, C. M., & Stockton, A. N. 2003, ApJ, 597,
730, astro-ph/0307300
Ly, C. et al. 2007, ApJ, 657, 738, astro-ph/0610846
Maier, C., Lilly, S. J., Carollo, C. M., Meisenheimer, K.,
Hippelein, H., & Stockton, A. 2006, ApJ, 639, 858,
astro-ph/0511255
Melbourne, J., & Salzer, J. J. 2002, AJ, 123, 2302,
astro-ph/0202301
Miyazaki, S., et al. 2002, PASJ, 54, 833, astro-ph/0211006
Oey, M. S. 2006, in ASP Conf. Ser. 353, Stellar Evolution at Low
Metallicity: Mass Loss, Explosions, Cosmology, ed.
H. J. G. L. M. Lamers, N. Langer, T. Nugis, & K. Annuk (San
Francisco: ASP), 253
Oke, J. B. 1990, AJ, 99, 1621
Osterbrock, D. E. 1989, Astrophysics of gaseous nebulae and
active galactic nuclei (Mill Valley, CA: University Science
Books)
Pagel, B.E.J., Simonson, E. A., Terlevich, R. J., & Edmunds, M.
G. 1992, MNRAS, 255, 325
Papovich, C. et al. 2004, ApJ, 600, L111, astro-ph/0310888
Pilyugin, L. S. & Thuan, T. X. 2005, ApJ, 631, 231
Popescu, C. C., Hopp, U., Hagen, H. J., & Elsaesser, H. 1996,
A&AS, 116, 43, astro-ph/9510127
Pustilnik, S. A., Engels, D., Kniazev, A. Y., Pramskij, A. G.,
Ugryumov, A. V., & Hagen, H.-J. 2006, Astronomy Letters, 32,
228, astro-ph/0508255
Salzer, J. J., et al. 2000, AJ, 120, 80, astro-ph/0004074
Sargent, W. L. W., & Searle, L. 1970, ApJ, 162, L155
Schiminovich, D., et al. 2005, ApJ, 619, L47, astro-ph/0411424
Schlegel, D. J., Finkbeiner, D. P., & Davis, M. 1998, ApJ, 500,
525, astro-ph/9710327, used to estimate dust extinctions from
source location by the NASA/IPAC Infrared Science Archive
Dust Reddening and Extinction Tool
(http://irsa.ipac.caltech.edu/applications/DUST/)
Seaton M. J. 1975, MNRAS, 170, 475
Skillman, E. D., Kennicutt, R. C., & Hodge, P. W. 1989, ApJ,
347, 875
Stanek, K. Z. et al. 2006, Acta Astron., 56, 333, astro-ph/0604113
Stasińska, G., Schaerer, D., & Leitherer, C. 2002, Ap&SS, 281,
335 McMahon, R. G. 2003, MNRAS, 342, 439,
astro-ph/0302212
Stern, D., Bunker, A., Spinrad, H., & Dey, A. 2000, ApJ, 537, 73,
astro-ph/0002239
Stockton, A., & Ridgway, S. E. 1998, AJ, 115, 1340,
astro-ph/9801056
Thuan, T. X., Izotov, Y. I. 2005, ApJS, 161, 240,
astro-ph/0507209
Tran, K.-V. H., Lilly, S. J., Crampton, D., & Brodwin, M. 2004,
ApJ, 612, L89, astro-ph/0407648
Turnshek, D. A., Bohlin, R. C., Williamson II, R. L., Lupie,
O. L., Koornneef, J., & Morgan, D.H. 1990, AJ, 99, 1243
Ugryumov, A. V., et al. 1999, A&AS, 135, 511
Wilson, G., Cowie, L. L., Barger, A. J., & Burke, D. J. 2002, AJ,
124, 1258, astro-ph/0203168
Wyder, T. K., et al. 2005, ApJ, 619, L15, astro-ph/0411364
http://arxiv.org/abs/astro-ph/0307325
http://arxiv.org/abs/astro-ph/0511471
http://arxiv.org/abs/astro-ph/0312635
http://arxiv.org/abs/astro-ph/0607280
http://arxiv.org/abs/astro-ph/9606079
http://arxiv.org/abs/astro-ph/0309105
http://arxiv.org/abs/astro-ph/0511066
http://arxiv.org/abs/astro-ph/9801003
http://arxiv.org/abs/astro-ph/0311528
http://arxiv.org/abs/astro-ph/0506498
http://arxiv.org/abs/astro-ph/0604234
http://arxiv.org/abs/astro-ph/0511644
http://arxiv.org/abs/astro-ph/9807187
http://arxiv.org/abs/astro-ph/0206495
http://arxiv.org/abs/astro-ph/0609246
http://arxiv.org/abs/astro-ph/0307401
http://arxiv.org/abs/astro-ph/0408128
http://arxiv.org/abs/astro-ph/0310346
http://arxiv.org/abs/astro-ph/9911094
http://arxiv.org/abs/astro-ph/9902334
http://arxiv.org/abs/astro-ph/0307300
http://arxiv.org/abs/astro-ph/0610846
http://arxiv.org/abs/astro-ph/0511255
http://arxiv.org/abs/astro-ph/0202301
http://arxiv.org/abs/astro-ph/0211006
http://arxiv.org/abs/astro-ph/0310888
http://arxiv.org/abs/astro-ph/9510127
http://arxiv.org/abs/astro-ph/0508255
http://arxiv.org/abs/astro-ph/0004074
http://arxiv.org/abs/astro-ph/0411424
http://arxiv.org/abs/astro-ph/9710327
http://irsa.ipac.caltech.edu/applications/DUST/
http://arxiv.org/abs/astro-ph/0604113
http://arxiv.org/abs/astro-ph/0302212
http://arxiv.org/abs/astro-ph/0002239
http://arxiv.org/abs/astro-ph/9801056
http://arxiv.org/abs/astro-ph/0507209
http://arxiv.org/abs/astro-ph/0407648
http://arxiv.org/abs/astro-ph/0203168
http://arxiv.org/abs/astro-ph/0411364
THE YOUNGEST GALAXIES 15
TABLE 2
NB816 Emission-Line Sample
No. RA(2000) Dec(2000) N(AB) Z(AB) I R V B z
1 40.115555 −1.694722 23.92 25.28 24.84 25.44 −99.00 −99.00 −1.0000
2 40.116665 −1.617361 24.26 25.88 25.52 25.74 −99.00 −99.00 −1.0000
3 40.138332 −1.405639 23.49 24.94 24.70 25.22 −99.00 −99.00 0.6343
4 40.174721 −1.704750 24.27 25.46 25.35 25.82 −99.00 −99.00 0.6355
5 40.183056 −1.495417 24.58 25.28 25.85 26.86 −99.00 −99.00 5.6886
6 40.216946 −1.494805 24.80 26.08 26.31 26.14 −99.00 −99.00 0.2416
7 40.250832 −1.744639 23.51 24.35 24.33 24.54 −99.00 −99.00 −1.0000
8 40.276112 −1.518139 24.72 24.78 25.53 25.75 −99.00 −99.00 −1.0000
9 40.276390 −1.623250 24.55 25.36 25.36 25.76 −99.00 −99.00 −1.0000
10 40.284168 −1.453583 21.82 23.20 23.02 22.56 −99.00 −99.00 0.2480
11 40.298889 −1.447389 24.32 25.95 25.42 25.71 −99.00 −99.00 −1.0000
12 40.304165 −1.391694 20.88 22.35 22.06 22.32 −99.00 −99.00 −1.0000
13 40.306946 −1.638500 24.91 25.67 25.72 26.02 −99.00 −99.00 −1.0000
14 40.311111 −1.535111 24.08 25.93 25.49 25.71 −99.00 −99.00 0.6240
15 40.318333 −1.548222 24.60 25.03 25.50 25.96 −99.00 −99.00 −1.0000
16 40.318890 −1.430889 23.37 23.77 24.32 24.42 −99.00 −99.00 1.1804
17 40.319168 −1.446333 23.60 23.85 24.43 24.56 −99.00 −99.00 1.1873
18 40.320835 −1.778028 20.70 21.57 21.60 21.86 −99.00 −99.00 −1.0000
19 40.324165 −1.409972 24.29 24.38 25.18 25.45 −99.00 −99.00 −1.0000
20 40.326111 −1.709805 23.24 24.75 24.62 24.96 −99.00 −99.00 −1.0000
21 40.336388 −1.570194 24.26 24.73 25.06 25.33 −99.00 −99.00 −1.0000
22 40.337223 −1.388194 24.81 27.37 26.71 26.77 −99.00 −99.00 −1.0000
23 40.337502 −1.658306 24.89 25.20 25.76 26.05 −99.00 −99.00 −1.0000
24 40.340279 −1.689472 24.55 26.37 26.07 26.51 −99.00 −99.00 −1.0000
25 40.340557 −1.551889 24.99 25.70 25.93 26.35 −99.00 −99.00 −1.0000
26 40.340832 −1.371222 22.42 23.43 23.29 23.16 −99.00 −99.00 −1.0000
27 40.340832 −1.516250 24.89 25.06 25.69 25.78 −99.00 −99.00 −1.0000
28 40.341110 −1.493500 23.37 24.62 24.55 24.23 −99.00 −99.00 −1.0000
29 40.341389 −1.484139 24.48 25.46 25.49 25.73 −99.00 −99.00 −1.0000
30 40.342777 −1.599528 24.65 25.31 25.50 25.85 −99.00 −99.00 −1.0000
31 40.343056 −1.438833 23.17 24.60 24.54 24.54 −99.00 −99.00 0.6226
32 40.347500 −1.403833 24.86 27.45 26.16 26.35 −99.00 −99.00 0.6324
33 40.349724 −1.598472 23.11 23.87 24.11 24.46 −99.00 −99.00 1.1956
34 40.356388 −1.515722 24.94 26.19 26.18 26.42 −99.00 −99.00 −5.0000
35 40.372223 −1.390361 23.85 24.68 24.72 24.62 −99.00 −99.00 −1.0000
36 40.373611 −1.722528 24.93 25.85 25.74 25.88 −99.00 −99.00 −1.0000
37 40.377777 −1.701889 23.45 23.82 24.26 24.70 −99.00 −99.00 −1.0000
38 40.388054 −1.697361 23.96 24.32 24.79 24.90 −99.00 −99.00 −1.0000
39 40.388889 −1.573361 24.79 25.13 25.65 26.07 −99.00 −99.00 −1.0000
40 40.394444 −1.521389 22.73 23.69 23.84 24.06 −99.00 −99.00 0.6292
Note. — Magnitudes are measured in 3′′ diameter apertures. An entry of ‘−99’ indicates that
no excess flux was measured. −1.0000 in the redshift column means no spectroscopic data were
obtained for the object. This is a sample table showing the first entries of the electronic version
of the table that will accompany the published paper.
16 Kakazu et al.
TABLE 3
NB912 Emission-Line Sample
No. R.A. (J2000.0) Decl. (J2000.0) N(AB) z′(AB) I R V B zspec
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10)
1 40.131668 −1.408361 23.86 25.08 25.89 25.97 −99.00 −99.00 0.8371
2 40.133888 −1.575222 24.76 25.92 26.13 25.84 −99.00 −99.00 1.4498
3 40.148056 −1.593555 23.21 24.62 25.53 26.11 −99.00 −99.00 0.8207
4 40.148335 −1.725417 24.76 25.97 26.98 26.39 −99.00 −99.00 0.8111
5 40.150833 −1.737556 23.31 24.35 24.88 25.32 −99.00 −99.00 0.8269
6 40.153332 −1.536833 23.64 25.00 26.18 26.87 −99.00 −99.00 0.8301
7 40.156387 −1.765833 21.93 23.05 23.51 23.74 −99.00 −99.00 −1.0000
8 40.165833 −1.580056 24.94 26.20 25.92 25.84 −99.00 −99.00 1.4482
9 40.183334 −1.389583 21.77 22.79 23.33 23.52 −99.00 −99.00 0.8325
10 40.184444 −1.596444 23.09 24.50 25.30 25.55 −99.00 −99.00 0.8293
11 40.193611 −1.690083 24.32 25.40 25.83 25.69 −99.00 −99.00 0.8266
12 40.194168 −1.373722 23.93 24.93 25.21 25.35 −99.00 −99.00 −1.0000
13 40.194721 −1.373917 24.87 25.90 26.30 26.47 −99.00 −99.00 −1.0000
14 40.196667 −1.378333 24.07 25.42 26.05 26.18 −99.00 −99.00 −1.0000
15 40.202221 −1.584472 24.22 25.30 25.81 26.19 −99.00 −99.00 0.8289
16 40.203335 −1.471861 24.77 26.20 25.78 25.54 −99.00 −99.00 0.3965
17 40.214722 −1.519917 23.14 24.40 25.38 25.94 −99.00 −99.00 0.8288
18 40.220276 −1.753778 24.34 25.41 26.40 26.29 −99.00 −99.00 −1.0000
19 40.220833 −1.388556 23.24 24.36 25.13 24.99 −99.00 −99.00 −1.0000
20 40.226944 −1.542111 23.02 24.47 25.77 25.48 −99.00 −99.00 0.8208
21 40.229168 −1.720889 24.99 27.85 27.11 · · · −99.00 −99.00 6.4800
22 40.229443 −1.376472 23.75 24.98 25.72 24.66 −99.00 −99.00 −1.0000
23 40.245834 −1.578972 24.61 25.82 27.22 26.91 −99.00 −99.00 0.8285
24 40.280556 −1.421056 24.82 25.86 26.13 25.91 −99.00 −99.00 −1.0000
25 40.290833 −1.746361 23.89 25.16 25.81 25.03 −99.00 −99.00 0.0000
26 40.323891 −1.697667 24.73 25.80 26.47 25.71 −99.00 −99.00 −1.0000
27 40.330833 −1.612389 23.03 24.68 26.06 25.47 −99.00 −99.00 0.0000
28 40.339722 −1.395361 23.87 25.54 26.98 25.45 −99.00 −99.00 0.3889
29 40.346668 −1.448305 23.93 24.98 25.37 25.48 −99.00 −99.00 0.8274
30 40.382500 −1.554056 23.52 24.87 25.60 25.08 −99.00 −99.00 0.3930
31 40.393055 −1.713694 24.94 26.83 26.70 26.50 −99.00 −99.00 −1.0000
32 40.398888 −1.466417 23.87 24.91 24.83 24.64 −99.00 −99.00 1.4590
33 40.403889 −1.530167 24.88 26.08 27.44 27.30 −99.00 −99.00 0.8223
34 40.409443 −1.369222 21.52 23.44 24.47 24.04 −99.00 −99.00 −1.0000
35 40.411667 −1.691417 24.41 25.50 25.68 25.55 −99.00 −99.00 −1.0000
36 40.424999 −1.454028 21.91 22.93 23.42 23.41 −99.00 −99.00 −1.0000
37 40.430000 −1.501111 23.13 24.21 24.76 24.92 −99.00 −99.00 0.8267
38 40.446110 −1.676139 24.91 26.09 26.37 26.04 −99.00 −99.00 −1.0000
39 40.478889 −1.534278 24.87 26.06 25.79 25.56 −99.00 −99.00 0.3861
40 40.506111 −1.755111 22.73 24.27 25.24 25.19 −99.00 −99.00 −1.0000
41 40.511665 −1.596944 24.73 25.96 25.55 25.67 −99.00 −99.00 −1.0000
42 40.518055 −1.666139 22.47 23.73 24.37 24.45 −99.00 −99.00 −1.0000
Note. — Magnitudes are measured in 3′′ diameter apertures. An entry of ‘−99’ indicates that no
excess flux was measured. −1.0000 in the redshift column means no spectroscopic data were obtained for
the object. This is a sample table showing the first entries of the electronic version of the table that will
accompany the published paper.
TABLE 4
Line fluxes and Oxygen Abundance for L816 selected emitters
Object f([OIII]5007) f([OIII]4959) f([OIII]4363) f([OII]3727) Te[OIII] 12+log(O/H)
[OIII] emitters
31 513.6 ± 24.0 222.3 ± 11.4 < 6.60 54.4 ± 4.62 < 1.19 > 8.09
40 577.9 ± 21.6 191.3 ± 8.05 9.40 ± 3.40 140.9 ± 6.22 1.14 < 1.34 < 1.53 7.86 < 8.03 < 8.25
51 401.5 ± 12.6 146.9 ± 5.52 9.40 ± 4.50 < 2.39 1.19 < 1.55 < 1.90 7.51 < 7.62 < 7.94
76 464.4 ± 10.5 191.3 ± 4.86 < 2.90 344.5 ± 8.07 < 0.95 > 8.55
118 492.6 ± 29.7 193.9 ± 12.9 34.4 ± 12.8 11.3 ± 2.61 2.16 < 3.08 < 4.32 6.93 < 7.16 < 7.44
195 335.0 ± 21.4 129.5 ± 10.2 24.0 ± 10.9 97.1 ± 9.71 2.02 < 3.17 < 4.86 6.78 < 7.06 < 7.44
206 597.1 ± 19.5 204.1 ± 7.41 21.6 ± 9.20 < 1.72 1.48 < 1.97 < 2.48 7.42 < 7.55 < 7.84
208 658.0 ± 30.9 249.8 ± 12.3 15.6 ± 8.00 < 22.1 1.16 < 1.56 < 1.93 7.67 < 7.85 < 8.22
223 242.9 ± 15.3 83.3 ± 7.53 23.7 ± 18.9 < 10.1 1.45 < 4.64 < 19.62 6.14 < 6.61 < 7.53
252 466.8 ± 9.32 157.9 ± 3.57 9.20 ± 4.00 139.0 ± 3.71 1.16 < 1.45 < 1.72 7.68 < 7.87 < 8.14
Note. — Only emitters with > 15σ Hβ fluxes are listed. All fluxes are normalized by their f(Hβ) and multiplied by
100. 1σ upper limits are listed for [OII]3727 flux below 3σ and [OIII]4363 below 1σ. The units of Te[OIII] are 10
−4 [K].
THE YOUNGEST GALAXIES 17
Fig. 17.— HST/ACS (B, V, z’) composite images of NB816 emitters in the GOODS-N field with overlaid object IDs from Table 2 and
redshifts, where known. Fields are 12.′′5 on a side.
TABLE 5
Line Fluxes and Oxygen Abundances for L912 selected emitters
Object f([OIII]5007) f([OIII]4959) f([OIII]4363) f([OII]3727) Te[OIII] 12+log(O/H)
[OIII] emitters
3 550.9 ± 12.9 187.9 ± 4.91 23.9 ± 7.90 8.6 ± 2.5 1.74 < 2.20 < 2.71 7.26 < 7.43 < 7.65
6 588.1 ± 35.1 216.0 ± 14.2 18.4 ± 11.0 52.0 ± 8.6 1.20 < 1.79 < 2.39 7.40 < 7.68 < 8.14
9 442.3 ± 15.3 154.7 ± 6.42 < 12.3 157.2 ± 6.88 < 1.70 > 7.70
10 490.0 ± 11.9 178.7 ± 5.14 13.7 ± 4.40 < 2.95 1.42 < 1.69 < 1.97 7.55 < 7.61 < 7.82
17 342.5 ± 20.0 129.7 ± 9.29 15.9 ± 9.40 < 7.72 1.41 < 2.26 < 3.28 7.03 < 7.22 < 7.70
20 418.7 ± 17.4 135.1 ± 6.96 16.8 ± 5.50 24.5 ± 2.45 1.69 < 2.11 < 2.57 7.18 < 7.36 < 7.58
239 202.4 ± 10.2 75.6 ± 6.40 < 8.20 190.6 ± 10.2 < 2.08 > 7.34
270 351.7 ± 15.1 149.7 ± 7.73 12.4 ± 3.40 30.7 ± 2.72 1.59 < 1.87 < 2.16 7.28 < 7.43 < 7.61
Hα emitters
52 589.1 ± 10.0 179.2 ± 3.42 18.3 ± 1.59 < 1.56 1.75 < 1.83 < 1.92 7.62 < 7.67 < 7.72
60 619.1 ± 33.5 206.7 ± 12.5 10.7 ± 7.77 48.4 ± 8.7 0.90 < 1.37 < 1.77 7.67 < 7.96 < 8.57
266 682.8 ± 10.3 217.7 ± 3.57 14.7 ± 2.42 ... 1.40 < 1.52 < 1.63 < 8.3
Note. — Same as Table 4 but for NB912 emitters. The [OII]λ3727 of ID266 is beyond the DEIMOS wavelength
coverage and thus was not being measured.
18 Kakazu et al.
Fig. 18.— HST/ACS (B, V, I) composite images of NB912 emitters in the GOODS-N field with overlaid object IDs from Table 3 and
redshifts, where known. Fields are 12.′′5 on a side.

  We describe results of a narrow band search for ultra-strong emission line
galaxies (USELs) with EW(H beta) > 30 A. 542 candidate galaxies are found in a
half square degree survey using two ~100 Angstrom 8150 A and 9140 A filters
with Subaru/SuprimeCam. Followup spectroscopy for randomly selected objects in
the sample with KeckII/DEIMOS shows they consist of [OIII] 5007, [OII] 3727,
and H alpha selected strong-emission line galaxies at intermediate redshifts (z
< 1), and Ly alpha emitting galaxies at high-redshift (z >> 5). We determine
the H beta luminosity functions and the star formation density of the USELs,
which is 5-10% of the value found from ultraviolet continuum objects at z=0-1,
suggesting they correspond to a major epoch in galaxy formation at these
redshifts. Many USELs show the temperature-sensitive [OIII] 4363 auroral lines
and about a dozen have oxygen abundances characteristic of eXtremely Metal Poor
Galaxies (XMPGs). These XMPGs are the most distant known today. Our high yield
rate of XMPGs suggests this is a powerful method for finding such populations.
The lowest metallicity measured in our sample is 12+log(O/H) = 7.06
(6.78-7.44), close to the minimum metallicity found in local galaxies. The
luminosities, metallicities and star formation rates of USELs are consistent
with the strong emitters being start-up intermediate mass galaxies and suggest
that galaxies are still forming in relatively chemically pristine sites at z <
1.

Submitted to The Astrophysical Journal
Preprint typeset using LATEX style emulateapj v. 08/22/09
MAPPING THE YOUNGEST GALAXIES TO REDSHIFT ONE1,2
Yuko Kakazu,
Lennox L. Cowie,
Esther M. Hu,
Submitted to The Astrophysical Journal
ABSTRACT
We describe the results of a narrow band search for ultra-strong emission line galaxies (USELs)
with EW(Hβ) ≥ 30 Å. 542 candidate galaxies are found in a half square degree survey using two
∼ 100Å filters centered at 8150Å and 9140Å with Subaru/SuprimeCam. Followup spectroscopy has
been obtained for randomly selected objects in the candidate sample with KeckII/DEIMOS and has
shown that they consist of [O III]λ5007, [O II]λ3727, and Hα selected strong-emission line galaxies at
intermediate redshifts (z < 1), and Lyα emitting galaxies at high-redshift (z >> 5). We determine
the Hβ luminosity functions and the star formation density of the USELs, which is 5-10% of the
value found from ultraviolet continuum objects at z = 0 − 1, suggesting that they correspond to
a major epoch in the galaxy formation process at these redshifts. Many of the USELs show the
temperature-sensitive [O III]λ4363 auroral lines and about a dozen have oxygen abundances satisfying
the criteria of eXtremely Metal Poor Galaxies (XMPGs). These XMPGs are the most distant known
today and our high yield rate of XMPGs suggests that narrowband method is powerful in finding
such populations. Moreover, the lowest metallicity measured in our sample is 12+log(O/H)=7.06
(6.78−7.44), which is close to the minimum metallicity found in local galaxies, though we need deeper
spectra to minimize the errors. HST/ACS images of several USELs exhibit widespread morphologies
from relatively compact high surface brightness objects to very diffuse low surface brightness ones. The
luminosities, metallicities and star formation rates of USELs are consistent with the strong emitters
being start-up intermediate mass galaxies which will evolve into more normal galaxies and suggest
that galaxies are still forming in relatively chemically pristine sites at z < 1.
Subject headings: cosmology: observations — galaxies: distances and redshifts — galaxies: abundances
— galaxies: evolution — galaxies: starburst
1. INTRODUCTION
The study of low-metallicity galaxies is of consider-
able interest for the clues that it can provide about the
first stages of galaxy formation and chemical enrichment.
We would also like to know if there are any genuinely
young galaxies undergoing their first episodes of star for-
mation at low redshifts. To date, the most metal-poor
systems studied have been the blue compact emission-
line galaxies found in the local Universe, with systems
such as I Zw 18 and SBS 0335-052W defining the low
metallicity boundary with measured 12+log (O/H) of
∼ 7.1 − 7.2 (Sargent & Searle 1970; Thuan & Izotov
2005; Izotov et al. 2005). More recently, the Sloan Digi-
tal Sky Survey (SDSS) has yielded additional extremely
metal-poor galaxies (XMPGs) (12+log (O/H) < 7.65 or
Z < Z⊙/12; Kniazev et al. 2003; Izotov et al. 2006a).
Despite enormous efforts, only a few dozen such XMPGs
are known, all at redshift z < 0.05 (e.g., Oey 2006; Izotov
2006b).
Historically, objective prism surveys have been used
to select emission-line galaxies for low-metallicity stud-
ies. (e.g. the Hamburg QSO Survey (Popescu et al.
1996) and its HSS sequel (Ugryumov et al. 1999) that
1 Based in part on data obtained at the Subaru Telescope, which
is operated by the National Astronomical Observatory of Japan.
2 Based in part on data obtained at the W. M. Keck Observa-
tory, which is operated as a scientific partnership among the the
California Institute of Technology, the University of California, and
NASA and was made possible by the generous financial support of
the W. M. Keck Foundation.
3 Institute for Astronomy, University of Hawaii, 2680 Woodlawn
Drive, Honolulu, HI 96822.
discovered HS 2134+0400 (Pustilnik et al. 2006) and
the Kitt Peak International Spectroscopy Survey (KISS;
Salzer et al. 2000; Melbourne & Salzer 2002)). The ad-
vantage of using the objective prism technique rather
than the continuum selection, employed with the SDSS
(Kniazev et al. 2003) or DEEP2 surveys (Hoyos et al.
2005), is that they have a higher efficiency and
provide a more uniform selection. By comparison,
continuum/broad-band surveys have a very low yield rate
(8 new XMPGs and 4 recovered XMPGS from an anal-
ysis of 250,000 spectra over ∼3000 deg2 for the SDSS
(Kniazev et al. 2003), since low-metallicity populations
in their first outburst have weak continuua and strong
emission lines.
An alternative method of discovering strong emission-
line, low-metallicity galaxies is to use narrowband
surveys. Strong emission-line galaxies have histori-
cally been picked up in high-z Lyman alpha searches
(e.g., Stockton & Ridgway 1998; Hu et al. 1998, 2004;
Stern et al. 2000; Tran et al. 2004; Ajiki et al. 2006)
where they have been considered contaminants. How-
ever, the low redshift emission line galaxies seen in these
surveys are of great interest in their own right as we shall
show in the present paper. While some spectroscopic
studies have been carried out for low-redshift galaxies se-
lected from narrowband surveys (e.g., Maier et al. 2006;
Ly et al. 2007), the small sample sizes have precluded
any detailed investigation of metallicity and identifica-
tion of a low-metallicity population.
The narrowband method probes to much deeper limits
than the objective prism surveys. This enables prob-
http://arxiv.org/abs/0704.0643v1
2 Kakazu et al.
ing star-forming populations out to near redshift z ∼ 1
where the cosmic star formation rates are considerably
higher. Furthermore the narrow-band emission-line se-
lection can allow us to assemble very large samples of
strong-emission line objects, with a clean selection of dif-
ferent line types for the construction of luminosity func-
tions.
Such a sample allows us to address such questions
as whether there are substantial populations of strong
star-forming galaxies with low metallicities among more
massive galaxies. There has been considerable contro-
versy about the interpretation of the low metallicity mea-
surements in the blue compact galaxy samples where
the ease with which gas may be ejected in these dwarf
galaxies has complicated the picture (e.g., Corbin et al.
2006) or, at least, resulted in identifying low metal-
licity systems which are not forming their first gen-
eration of stars. The identification of low metallicity
galaxies – at the level of the XMPGs – among mas-
sive galaxies can provides less ambiguous examples of
galaxies that are genuinely ‘young’ and caught in the
initial stages of star formation. Current efforts to iden-
tify low metallicity galaxies from continuum selected
surveys (e.g., Kobulnicky et al. 2003; Lilly et al. 2003;
Kobulnicky & Kewley 2004; Hoyos et al. 2005) have low-
metallicity thresholds that are higher than this – about
one-third solar (in O/H). With a narrow-band selec-
tion criterion much larger emission-line samples includ-
ing such low metallicity galaxies can be identified. With
these large samples it is also possible to determine
whether there is an observed lower metallicity threshold
for such galaxies, and to estimate what the contribution
of such strong star-formers might be at these epochs.
In the present work we use a number of deep, narrow-
band images obtained with the SuprimeCam mosaic
CCD camera (Miyazaki et al. 2002) on the Subaru 8.2-
m telescope to find a large sample of extreme emission-
line galaxies. We first (§2) outline the selection criteria
(magnitude and flux thresholds) for the target fields re-
sulting in a sample of 542 galaxies, which we call USELs
(Ultra-Strong Emission Liners). We then describe (§3)
the spectroscopic followups for 161 of these galaxies us-
ing multi-object masks with the DEIMOS spectrograph
(Faber et al. 2003) on the 10-m Keck II telescope. Sam-
ple spectra for each class of object are shown. Flux
calibration and equivalent width distributions are pre-
sented in §4, and the resulting measured line ratios are
discussed. In §5 luminosity functions are constructed and
star formation rates are estimated for the sample. These
galaxies are estimated to contribute roughly 10% to the
measured star-formation rate (without extinction correc-
tions) at this epoch. Analysis of the metallicities is given
in §6. Their morphologies and dynamical masses are dis-
cussed in §7 and a final summary discussion is given in
§8. We use a standard H0 = 70 km s
−1 Mpc−1, Ωb =
0.3, ΩΛ = 0.7 cosmology throughout.
2. THE NARROW BAND SELECTION
The emission line sample was chosen from a set of
narrow band images obtained with the SuprimeCam
camera on the Subaru 8.2-m telescope. The data were
obtained in a number of runs between 2001 and 2005
under photometric or near photometric conditions. We
used two ∼120 Å (FWHM) filters centered at nominal
5000 6000 7000 8000 9000 10000
WAVELENGTH (A)
Fig. 1.— Schematic illustration of the selection process and a
typical spectrum of the galaxies we find. The objects are chosen
based on their excess light in one of two narrow band filters at 8160
Å and 9140Å. The present case corresponds to an Hα emission
line object found in the 9140Å filter (shown with the narrow solid
curve). Also illustrated are the broad band V (dash-dot), R (solid),
I (dashed), and z′ (dotted) filters use to measure the continuum.
The spectrum shown corresponds to object 205 in Table 3 and is an
Hα emitter at z = 0.3983. The easily visible lines are the Balmer
series and the [O III] lines at λλ5007, 4959, and 4363Å.
wavelengths of 8150 Å and 9140 Å in regions of low
sky background between the OH bands. The nominal
specifications for the Subaru filters may be found at
http://www.naoj.org/Observing/Instruments/SCam/sensitivity.html
and are also described in Ajiki et al. (2003). We shall
refer to these filters as NB816 and NB912.
About 5 hour exposures were obtained with NB816 and
∼10 hour exposures with NB912 yielding 5 sigma limits
fainter than 26 mags in both bands. Deep exposures in
B, V , R, I and z′ were also taken for the fields. The data
were taken as a sequence of dithered background-limited
exposures and successive mosaic sequences were rotated
by 90 degrees. Only the central uniformly covered areas
of the images were used. Corresponding continuum expo-
sures were always obtained in the same observing run as
the narrowband exposures to avoid false identifications
of transients such as high-z supernovae, or Kuiper belt
objects, as emission-line candidates. A detailed descrip-
tion of the full reduction procedure for images is given in
Capak et al. (2004). All magnitudes are given in the AB
system (Oke 1990). These were measured in 3′′ diameter
apertures, and had average aperture corrections applied
to give total magnitudes.
The primary purpose of the program was to study Lyα
emitters at redshifts of z ∼ 5.7 and z ∼ 6.6 (Hu et al.
2004, 2007) but the narrow band imaging is also ideal
for selecting lower redshift emission line galaxies and it
is for this purpose that we use these data in the present
paper. The fields which we use and the area covered (ap-
proximately a half square degree in each bandpass) are
summarized in Table 1. These are distributed over the
sky to deal with cosmic variance. We selected galaxies in
the narrowband NB816 filter using the Cousins I band
filter as a reference continuum bandpass and including
all galaxies with NB816 < 25 and I − NB816 greater
than 0.8. We selected galaxies in the NB912 filter with
the z filter as the reference continuum bandpass and in-
http://www.naoj.org/Observing/Instruments/ SCam/sensitivity.html
THE YOUNGEST GALAXIES 3
TABLE 1
Narrowband Survey Area Coverage
Field RA (J2000) Dec (J2000) (lII, bII) EB−V
a NB816 NB912
(arcmin2) (arcmin2)
SSA22 22:17:57.00 +00:14:54.5 ( 63.1,−44.1) 0.07 674 591
SSA22 new 22:18:24.67 +00:36:53.4 ( 63.6,−43.9) 0.06 278 278
A370 new 02:41:16.27 −01:34:25.1 (173.4,−53.3) 0.03 278 278
HDF-N 12:36:49.57 +62:12:54.0 (125.9,+54.8) 0.01 710 528
Total 1940 1675
Note. — An adjacent field to A370 new is a site of a gravitational lensing cluster at
z ∼ 0.375, and was omitted from the suvey.
a estimated using http://irsa.ipac.caltech.edu/applications/DUST/ based on
Schlegel et al. (1998)
20 21 22 23 24
N(AB)
20 21 22 23 24
N(AB)
WITH SPECTRA
Fig. 2.— Continuum – Narrow band magnitude versus narrow
band magitude for all objects with narrow band magnitude brighter
than 24. The diamonds show the narrowband excess emission mag-
nitude of the NB912 sample and the squares the NB816 sample.
Galaxies which would be included in an R < 24 continuum selected
sample are shown with solid symbols. The upper panel shows the
complete sample while the lower panel shows the subsample which
has been spectroscopically identified.
cluded all galaxies with NB912 < 25 and z − NB912
greater than 1. The selection process is illustrated for
a galaxy found in the NB912 filter in Figure 1. Both
selections correspond roughly to choosing objects with
emission lines with rest frame equivalent widths greater
than 100 Å. The exact equivalent width limit depends
on the precise position of the emission line in the filter
and the redshift of the galaxy which in turn depends on
which emission line is producing the excess light in the
narrow band.
The final USEL sample consists of 542 galaxies (267
in the NB816 filter and 275 in the NB912 filter). Tabu-
lated coordinates, multi-color magnitudes, and redshifts
(where measured) for these objects are summarized in
Table 2. Very few of these objects would be included
in continuum-selected spectroscopic samples. Figure 2
shows the narrow band excess as a function of narrow-
band (NAB) magnitude for objects with narrow band
magnitudes brighter than 24. The open symbols show
the present sample while the solid symbols show ob-
jects which would be included in an R < 24 continuum-
selected sample.
3. SPECTRA
Spectroscopic observations were obtained for 161
USELs from the sample using the Deep Extragalac-
tic Imaging Multi-Object Spectrograph (DEIMOS;
Faber et al. 2003) on Keck II in a series of runs between
2003 and 2006. The emission line objects were included
in masks designed to observe high-z Lyα candidates and,
as can be seen in the lower panel of Figure 2, constitute
an essentially random sample of the emission line galax-
The observations were primarily made with the G830
ℓ/mm grating blazed at 8640 Å and used 1′′ wide slitlets.
In this configuration, the resolution is 3.3 Å, which is
sufficient to distinguish the [O II]λ3727 doublet struc-
ture. This allows us to easily identify [O II]λ3727 emit-
ters where often the [O II]λ3727 doublet is the only emis-
sion feature. The spectra cover a wavelength range of
approximately 4000 Å and were centered at an average
wavelength of 7800 Å, though the exact wavelength range
for each spectrum depends on the slit position with re-
spect to the center of the mask along the dispersion di-
rection. The G830 grating used with the OG550 blocker
gives a throughput greater than 20% for most of this
range, and ∼ 28% at 8150 Å. The observations were
not generally taken at the parallactic angle, since this
was determined by the mask orientation, so considerable
care must be taken in measuring line fluxes as we dis-
cuss below. Each ∼ 1 hr exposure was broken into three
subsets, with the objects stepped along the slit by 1.5′′
in each direction. Some USELs were observed multiple
times, resulting in total exposure times for these galaxies
of 2−3 hours. The two-dimensional spectra were reduced
following the procedure described in Cowie et al. (1996)
and the final one-dimensional spectra were extracted us-
ing a profile weighting based on the strongest emission
http://irsa.ipac.caltech.edu/applications/DUST/
4 Kakazu et al.
4000 4500 5000 5500 6000 6500 7000
REST WAVELENGTH (A)
Fig. 3.— Spectrum of an Hα emission galaxy selected in the NB912 filter. In the upper plot we have decreased the scale of the vertical
axis by a factor of 10 to show the continuum and the weaker lines. The more important emission line features are labelled and marked
with the dotted lines.
3500 4000 4500 5000 5500
REST WAVELENGTH (A)
Fig. 4.— Spectrum of an [O III] galaxy in the NB816 selected sample. The lower plot shows the relative strengths of the very strong
emission lines in the spectrum. In the upper plot we have decreased the scale of the vertical axis by a factor of 10 to show the continuum
and the weaker lines. The more important emission line features are labelled and marked with the dotted lines.
THE YOUNGEST GALAXIES 5
3500 4000 4500 5000 5500
REST WAVELENGTH (A)
Fig. 5.— Spectrum of an [O III] galaxy selected in the NB912 filter. The lower plot shows the relative strengths of the very strong
emission lines in the spectrum. In the upper plot we have decreased the scale of the vertical axis by a factor of 10 to show the continuum
and the weaker lines. The more important emission line features are labelled and marked with the dotted lines.
3000 3200 3400 3600 3800 4000 4200 4400
REST WAVELENGTH (A)
Fig. 6.— Spectrum of an [O II] galaxy selected in the NB816 filter. The plot shows the relative strengths of the very strong emission
lines in the spectrum. The more important emission line features are labelled and marked with the dotted lines.
6 Kakazu et al.
0.0 0.5 1.0 1.5
REDSHIFT
[OIII]
Hα [OII]
0.0 0.5 1.0 1.5
REDSHIFT
10-18
10-17
10-16
10-15
Fig. 7.— (a) Distribution of redshifts for the spectroscopically
identified sources. [O III] λ5007 emitters are the most common.
Since the focus of this paper is on intermediate-redshift (z . 1)
strong emission line galaxies, we did not plot high redshift Lyα
galaxies (z >> 5) in our NEO sample. High-z Lyα emitters are
discussed in Hu et al. (2004, 2007). (b) Flux versus redhift for the
spectroscopically identified sample. Squares are Hα, diamonds are
[O III] λ5007, and triangles are [O II] λ3727. The solid line shows
the flux limit corresponding to the narrow band magnitude limit of
N(AB)=25 for an emitter with very large equivalent width. Some
objects with lower equivalent widths fall below this limit.
line in the spectrum. A small number of the spectra
were obtained with the ZD600 ℓ/mm grating giving a
correspondingly lower resolution but a wider wavelength
coverage. These observations were centered at 7200 Å.
Essentially all of the emission line candidates which
were observed were identified, though two of the objects
in the NB816 sample are stars where the absorption line
structure mimics emission in the band. Sample spec-
tra are shown in Figures 3, 4, 5, and 6. The measured
redshifts are given in Tables 2 and 3. The narrow band
emission line selection produces a mixture of objects cor-
responding to Hα, [O III]λ5007, and [O II]λ3727 and, at
the faintest magnitudes (> 24), high redshift Lyα emit-
ters. The number of objects seen in each line and the
redshifts where they are found are shown in Figure 7.
The spectroscopically identified sample from both bands
contains 13 Hα, 92 [O III]λ5007, and 23 [O II]λ3727 emit-
ters. In the remainder of the paper we shall focus primar-
ily on the Hα and [O III]λ5007 selected galaxies which lie
between redshifts zero and one.
Since only 30% of the USELs are spectroscopically
identified we must apply a substantial incompleteness
correction in computing the line luminosity function
and the universal star formation histories. Because
the type mix may vary as a function of magnitude we
have adopted a magnitude dependent weighting for each
galaxy equal to the total number of galaxies at this mag-
nitude divided by the number of spectroscopically iden-
tified galaxies. However, the analysis is not particularly
sensitive to the adopted scheme since the fraction of iden-
tified galaxies is relatively constant with magnitude.
4. FLUX CALIBRATIONS
Generally our spectra were not obtained at the par-
allactic angle since this is determined by the DEIMOS
mask orientation. Therefore flux calibration using stan-
dard stars is problematic due to atmospheric refraction
effects, and special care must be taken for the flux cali-
bration. We thus employed three independent methods
for the flux calibration. In §4.1 we introduce “primary
fluxes” which are absolute fluxes of the emission lines
used to select the galaxies. Primary fluxes are computed
from the SuprimeCam broadband and narrowband mag-
nitudes. We use these fluxes to derive luminosity func-
tions of Hα and [O III]λ5007 emitters at z < 1 (§5.1).
In §4.2 we measure line fluxes from the spectra. Rela-
tive line fluxes can be measured from the spectra without
flux calibration as long as we restrict the line measur-
ments to over short wavelength range where the DEIMOS
response is essentially constant. For example, one can as-
sume the response of neighboring lines (e.g. [O III]λ4949
and [O III]λ5007) are the same and therefore one can
measure the flux ratio without calibration. For bright
galaxies, we can absolutely calibrate the fluxes by inte-
grating spectra and equating them to Subaru broadband
fluxes. These line fluxes derived from the spectra are
used as a check of the primary fluxes. We show that
the ratio of [O III]λ5007/[O III]λ4959 is indeed close to
1/3, and that the fluxes computed from the spectra are
highly consistent with the primary fluxes measured in
§4.1. In §4.3, we show Balmer flux ratios f(Hβ)/f(Hα) of
bright Hα emitters are close to the Case B conditions,
suggesting very little reddening.
Metallicity measurements by the direct method re-
quire four emission lines that are widely displaced
over the spectral wavelength range ([O III]λλ4959, 5007,
[O III]λ4363, and [O II]λ3727). To calibrate these lines,
we used neighboring Balmer lines with the assumption
of Case B conditions, and this is described in §4.4.
4.1. Narrow Band Fluxes − Primary Fluxes
For the emission lines used to select each galaxy we
may compute the equivalent widths and absolute fluxes
directly from the narrow band magnitudes (N) and
the corresponding continuum magnitudes (C) from our
SuprimeCam imaging data. For example, in the case of
the NB816 selected emission-line galaxies, N corresponds
to the NB816 magnitude and C is the I band magnitude.
We shall refer to the values calculated in this way as the
primary fluxes and use this quantity to compute the lu-
minosity functions for each emitter in §5.1.
Defining the quantity
R = 10−0.4∗(N−C)
THE YOUNGEST GALAXIES 7
1.5 2.0 2.5 3.0
LOG REST FRAME EW ([OIII]5007)
1.5 2.0 2.5 3.0
LOG REST FRAME EW (Hα)
Fig. 8.— (a) Distribution of the rest frame equivalent widths
determined from the narrow band magnitudes for the spectroscop-
ically identified [O III] λ5007 sources. (b) Distribution of the rest
frame equivalent widths for the Hα selected sample.
the observed frame equivalent width becomes
where φ is the narrow band filter response normalized
such that the integral over wavelength is unity and ∆λ is
the effective width of the continuum filter. The narrow
band filter is often assumed to be rectangular in which
case φ becomes 1/δλ where δλ is the width of the narrow
band but as can be seen from Figure 1 this is not a very
good approximation in the present case. For very high
equivalent width objects the denominator in this equa-
tion becomes uncertain unless the broad band data are
very deep, and this can result in a large scatter in the
very highest equivalent widths where the continuum is
poorly determined.
In the case of the [O III]λ5007 line we must include the
second member of the doublet which also lies within the
narrow band filter. We have computed these cases as-
suming the flux of the [O III]λ4959 line is 0.34 times that
of the [O III]λ5007 line. Then φ = φ1+0.34×φ2 where φ1
is the filter response at the redshifted 5007 Å wavelength
and φ2 is the filter response at redshifted 4959 Å.
The distribution of the rest frame equivalent widths for
the Hα and [O III]λ5007 samples is shown in Figure 8.
The [O III]λ5007 sample selects objects with rest frame
equivalent widths above about 100Å while the lower red-
shift Hα sample selects objects with rest frame equivalent
widths above about 150Å. Since the [O III]λ5007 lines are
also generally stronger than the Hα lines the [O III] selec-
tion chooses less extreme objects than the Hα selection
and will include a larger fraction of galaxies at the given
10-18 10-17 10-16 10-15
f(OIII 5007) erg cm-2 s-1
Fig. 9.— The ratio of the [O III] λ4959 line to [O III] λ5007. The
errors are plus and minus 1 sigma. The median ratio is 0.338 and
the scatter around this value is consistent with that expected from
the statistical errors.
redshift.
The high observed frame equivalent widths make the
line fluxes insensitive to the continuum determination
and these may simply be found from
f = A
10−0.4N − 10−0.4C
where A is the AB zeropoint at the narrow band wave-
length in units of erg cm−2 s−1 Å−1. The flux depends
on the filter response at the emission line wavelength and
correspondingly is most uncertain at the edges of the fil-
ters where this quantity changes rapidly. Primary fluxes
defined here are measured by using narrowband (N) and
broadband (C) magnitudes from Subaru imaging data
with the object redshift information from Keck spectra
for the filter response at the exact location of emission
line wavelength (φ). We use these primary fluxes to con-
struct the luminosity functions of Hα and [O III]λ5007
selected emitters as we discuss in §5.1.
4.2. Line Fluxes from the Spectra
For the short wavelength range where DEIMOS re-
sponse is essentially constant, we may also compute the
relative line fluxes from the spectra without calibration.
For each spectrum we fitted a standard set of lines. For
the stronger lines we used a full Gaussian fit together
with a linear fit to the continuum baseline. For weaker
lines we held the full width constant using the value mea-
sured in the stronger lines and set the central wavelength
to the nominal redshifted value. We also measured the
noise as a function of wavelength by fitting to random
positions in the spectrum and computing the dispersion
in the results.
These fits should provide accurate relative fluxes over
short wavelength intervals where the DEIMOS response
is similar, but may be expected to be poorer over longer
wavelength intervals where the true calibration can vary
from the adopted value. We tested the relative flux cali-
bration for neighboring lines and the noise measurement
by measuring the ratio of the [O III]λ4959/ [O III]λ5007
lines. This is expected to have a value of approximately
0.34. The ratio is shown as a function of the [O III]λ5007
8 Kakazu et al.
10-17 10-16 10-15
FLUX (erg cm-2 s-1)
Fig. 10.— Ratio of fluxes computed from the spectra and the
broad band magnitudes versus those from the narrow band mag-
nitudes. Hα lines are shown as solid boxes, [O III] λ5007 lines as
diamonds and [O II] λ3727 lines as crosses.
flux in Figure 9. The measured values scatter around the
expected value and the dipsersion is consistent with the
noise determination. This result supports our assump-
tion of [O III]λ4959/[O III]λ5007 = 0.34 in the primary
fluxes measurements described in §4.1.
The brighter objects may be absolutely calibrated us-
ing the continuum magnitudes obtained from our Sub-
aru data. We integrated the spectrum convolved with
each SuprimeCam filter response and set this equal to
the broad band flux to normalize the spectrum in each
of the filters. We then used the Gaussian fits to obtain
the spectral line fluxes for lines lying within that broad
band. This procedure only works for sources with well
determined continuum magnitudes (C < 24.5) where the
sky subtraction can be well determined in the spectra.
For these objects the spectrally determined fluxes are
shown versus the primary fluxes in Figure 10 where we
plot the ratio of the spectral to the primary flux versus
the primary line flux. The values agree extremely well
though the measured spectral line fluxes are on average
about 80 − 90% of the primary flux values. This may
reflect a selection bias in the choice of the objects or the
narrow band filters could be slightly narrower than the
nominal profiles.
4.3. Balmer Ratios
We now measured the Balmer ratios for the sample of
objects selected in Hα and where the continuum magni-
tudes were bright enough to absolutely flux calibrate the
spectra. The ratio of f(Hβ)/f(Hα) is shown in Figure 11.
The values average closely to the Case B ratio which is
shown as the solid line and at brighter fluxes the indi-
vidual values also closely match to this value suggesting
that the galaxies have very little reddening. However,
at fainter fluxes the scatter about the average value is
considerably higher than the statistical errors. At the
faintest fluxes it appears that the systematic uncertainty
can be as high as a multiplicative factor of two.
4.4. Final Flux Calibration for Metallicity Analysis
For the metallicity analysis we adopted the pro-
cedure of normalizing the longer wavelength lines
([O III]λλ4959, 5007, [O III]λ4363) to their nearest
10-18 10-17 10-16 10-15
f(Hα) erg cm-2 s-1
Fig. 11.— The ratio of the Hβ/Hα fluxes versus Hα flux. The
values average to the unreddened Balmer decrement shown by the
solid line but at the lower fluxes the scatter is larger than expected
from the statistical errors reflecting the calibration uncertainties for
the fainter sources. The figure shows the ten objects detected in the
Hα line with continuum magnitudes above 24.5 in the bandpasses
corresponding to the lines.
Balmer line to determine the unreddened fluxes. For ex-
ample, in the case of the Hα emission selected galaxies,
we can measure Hα absolute fluxes by the primary fluxes
method described in §4.1. We can then derive Hβ and Hγ
fluxes from Hα fluxes by assuming Case B recombination
[e.g., f(Hα) = 2.85 × f(Hβ), f(Hγ) = 0.469 × f(Hβ) at T
= 104 K and Ne ∼ 10
2−104cm−3; Osterbrock 1989]. As
Hβ and [O III]λλ4959, 5007 have very similar DEIMOS
response, the relative fluxes should remain the same with
or without the flux calibration and this can be expressed
by a simple equation:
f0(Hβ)
f0([O III]λ4959, λ5007)
f(Hβ)
f([O III]λ4959, λ5007)
where f0(Hβ) and f0([O III]λ4959, λ5007) are the flux
counts in the un-calibrated, reddened DEIMOS spec-
tra, while f(Hβ) and f([O III]λ4959, λ5007) are absolute,
unreddened fluxes. Since we know f(Hβ) from f(Hα)
with the Case B assumption and f0(Hβ)/f0([O III]λ4959,
λ5007) from the DEIMOS spectra, we can derive
f([O III]λ 4959, λ5007) using this simple formula. Sim-
ilary, we can absolutely calibrate [O III]λ4363 lines by
using its neighboring Balmer line, Hγ:
f0(Hγ)
f0([O III]λ4363)
f(Hγ)
f([O III]λ4363)
where f0(Hγ) and f0([O III]λ4363) are again the counts
in flux uncalibrated, reddened DEIMOS spectra, and
f(Hγ) and f([O III]λ4363) are absolute fluxes.
In the case of the [O III] selected emitters, we first de-
rive [O III]λλ4959, 5007 absolute fluxes using the primary
fluxes method (§4.1), and then use the above formula to
get absolute fluxes of Hβ, then Hγ (by the Case B ratio),
and finally [O III]λ4363.
This flux calibration technique using neighboring
Balmer line should work well for the [O III]λλ4959, 5007,
λ4363 lines and the [N II] lines which all lie close to
Balmer lines but may be slightly more approximate for
the [S II] lines. The higher order Balmer lines are too un-
certain to apply this procedure due to inadequate S/N of
THE YOUNGEST GALAXIES 9
40 41 42 43
L(Hα) (erg s-1)
1 z=0.24
40 41 42 43
L(Hα) (erg s-1)
1 z=0.39
Fig. 12.— The luminosity function of Hα at z = 0.24 (top panel)
and at z = 0.39 (bottom panel). In each case the open boxes
show the luminosity functions determined from the spectroscopic
sample alone while the solid boxes show the function corrected for
the incompleteness in the spectroscopic identification. The errors
are plus and minus 1 sigma and at the highest luminosity we show
the 1 sigma upper limit.
the lines, and we have used the continuum flux calibrated
values with no reddening for the [O II]λ3727 and [Ne III]
lines. These values will have correspondingly higher flux
uncertainties. Fortunately the [O II]λ3727 line is very
weak in most of the objects and the uncertainty has lit-
tle effect on the metallicity determinations. However,
ionization analyses based on the [Ne III] line should be
undertaken with caution.
5. STAR FORMATION HISTORY
5.1. Hα and [O III]λ5007 Luminosity Functions
Because of the high observed frame equivalent widths
the primary fluxes are insensitive to the continuum de-
termination. However, they do depend on the filter re-
sponse at the emission line wavelength so we first restrict
ourselves to redshifts where the nominal filter response
is greater than 70% of the peak value. This also has
the advantage of providing a uniform selection and we
assume the window function is flat over the defined red-
shift range. Now the volume is simply defined by the se-
lected redshift range for all objects above the minimum
luminosity which we take as corresponding to a flux of
1.5×10−17 erg cm−2 s−1 (Figure 7). The luminosity func-
tion is now obtained by dividing the number of objects
40 41 42 43
L(OIII) (erg s-1)
1 z=0.63
40 41 42 43
L(OIII) (erg s-1)
1 z=0.83
Fig. 13.— The luminosity function of [O III] λ5007 emitters at
z = 0.63 (top panel) and at z = 0.83 (bottom panel). In each case
the open boxes show the luminosity functions determined from the
spectroscopic sample alone while the solid boxes show the function
corrected for the incompleteness in the spectroscopic identification.
The errors are plus and minus 1 sigma and at the highest luminosity
we show the 1 sigma upper limit.
in each luminosity bin by the volume. The incomplete-
ness corrected luminosity function is obtained from the
sum of the weights in each luminosity bin divided by the
volume. The 1 sigma errors shown are calculated from
the Poissonian errors based on the number of objects in
the bin. The calculated Hα luminosity function is shown
for the two redshift ranges corresponding to the NB816
and NB912 selections in Figure 12 and the corresponding
[O III]λ5007 luminosity functions in Figure 13.
5.2. Star Formation Rates
The individual objects have Hα luminosities stretching
up to about 3×1041 erg s−1 where, at the higher redshifts,
we use the Hβ luminosity to infer the Hα value assuming
there is no reddening. For a steady formation this would
require a star formation rate of a few solar masses per
year if we adopt the Kennicutt (1998) conversion rate for
his Salpeter mass function.
Since the objects are more probably caused by star-
bursts the true star formation rate will depend on the
evolutionary history. However, the Hα luminosity den-
sity should give a reasonable estimate of the universal
star formation density of the objects provided only that
most of the ionizing photons are absorbed in the galax-
10 Kakazu et al.
Fig. 14.— The star formation history inferred from the Hα or
Hβ luminosity density as a function of redshift. The data from our
sample are shown in red. The open squares show the total rate
from the entire sample while the solid squares show the values for
objects with Hα rest frame equivalent widths in excess of 200Å or
Hβ rest frame equivalent widths in excess of 70Å. The diamonds
show the UV star formation rates (uncorrected for extinction) from
the ground based work of Wilson et al. (2002) and the triangles
the Galex results of Wyder et al. (2005) and Schiminovich et al.
(2005). Hα measurements from the literature as summarized in
Ly et al. (2007) are shown with filled circles. In all cases the errors
are ±1σ.
ies. We first formed the total Hα or Hβ luminosity den-
sity of the galaxies by summing over the incompleteness
weighted luminosities in each redshift interval. We only
included detected objects and did not attempt to extrap-
olate to lower luminosities but the result are not partic-
ularly sensitive to this because the luminosity functions
are relatively flat at the lower luminosities (Figures 13
and 14). We then converted these to star formation rates
with the Kennicutt (1998) conversion.
The results are shown in Figure 14. We first plot
the rate for the total samples at each redshifts shown
by the open squares. We have shown star formation
rates for UV continuum samples for comparison and the
present sample of strong emitters gives star formation
rates which are about 10% of the UV values over the
redshift interval. For comparison, we also show the star
formation rates from Hα selected samples reported in the
literature and summarized in Ly et al. (2007). In order
to better understand the evolution we have also restricted
our own sample to objects with rest frame equivalent
widths of Hα in excess of 200Å at low redshifts and Hβ
equivalent widths in excess of 70Å at the higher redshifts.
The star formation rates for this sample are shown with
the solid squares. This provides a more uniform selection
with redshift and gives a slower increase than the total
inferred star formation rate. For this sample the SFR is
evolving roughly as (1+z)3 broadly similar to other UV
and optically determined formation rates in this redshift
interval. These more restricted objects comprise about
5% relative to the UV star formation rates at the higher
redshifts.
6. GALAXY METALLICITIES
6.1. [O III] emitters
The spectra are of variable quality and, in order to
measure the metallicities, we need very high signal to
noise observations. It is also important that Balmer lines
are well detected since our flux calibrations are done us-
ing the neighboring Balmer lines (§4.4). We therefore
restricted ourselves to [O III] emitters whose Hβ fluxes
are detected above 15 sigma. Among 92 [O III] emitters
in our total spectroscopic sample, 8 such [O III] emitters
were chosen in the NB912 sample, and 10 in the NB816.
These emitters have Hγ detected above 4 sigma. Tables 4
and 5 give the oxygen line fluxes normalized by their Hβ
fluxes for the NB816 and NB912 selected emitters, re-
spectively. 1σ upper limits are listed when the measured
flux is below 1σ.
The most direct method to estimate the gas-phase
oxygen abundance is to use the electron temperature of
the HII region. Higher gas metallicity increases nebular
cooling, leading to lower electron temperatures. There-
fore electron temperature is a good indicator of the
gas metallicity. The electron temperature can be de-
rived from the ratio of the [O III]λ4363 auroral line to
[O III]λλ5007,4959. This procedure is often referred to
as the ‘direct’ method or Te method (e.g., Seaton 1975;
Pagel et al. 1992; Pilyugin & Thuan. 2005; Izotov et al.
2006c). One well-known problem with the direct method,
however, is that [O III]λ4363 is generally very weak even
in the low-metallicity galaxies. For higher metallicity
systems, the far-IR lines become the dominant coolant
and therefore the optical auroral line is essentially not
detectable. However, the majority of our sample exhibit
[O III]λ4363, already suggesting that these are metal-
deficient systems. To derive Te[O III] and the oxygen
abundances, we used the Pagel et al. (1992) calibrations
with the Te[O II]−Te[O III] relations derived by Garnett
(1992). The results are shown in Table 4 (for NB816
selected [O III] emitters) and Table 5 (for NB912 se-
lected [O III] emitters). The Izotov et al. (2006c) for-
mula, which was developed with the latest atomic data
and photoionization models, gives consistent abundances
within 0.1 dex. The [S II]λλ6717, 6731 lines that are usu-
ally used for the determination of the electron number
density, are beyond the Keck/DEIMOS wavelength cov-
erage for our [O III] emitters. Therefore we assumed ne
= 100 cm−3. However the choice of electron density has
little effect as electron temperature is insensitive to the
electron density; we get the same results even when we
use ne = 1000 cm
The 1σ upper and lower limits on Te[O III] and the
oxygen abundances are also shown in the tables. Because
the [O III]λ4363 flux is weak (< 4σ), the range of our
metallicity measurements are quite wide. However, of
18 [O III] emitters, even the upper metallicity limits on 6
emitters satisfy the definition of XMPGs [12 + log(O/H)
< 7.65; Kunth & Östlin 2000]. All our emitters, except
the ones that only have lower limits on metallicties such
as ID31 in Table 4, have very low metallicities – even the
upper metallicity limits are about 0.02 - 0.2 Z⊙. A few
emitters may even have metallicities that are comparable
to the currently known most metal-poor galaxies [I Zw
18 and SBS0335−052W; 12 + log(O/H) ∼ 7.1 − 7.2].
However our current metallicity errors are too large to
measure the baseline metallicity accurately and higher
S/N spectra will be necessary for this purpose.
Our discovery rate of XMPGs appers to be signif-
icantly higher than those of other surveys. Only 14
new XMPGs have been discovered from the analysis of
THE YOUNGEST GALAXIES 11
Fig. 15.— [O III]λ4959+λ5007/[O III]λ4363 versus
[O II]λ3727/[O III]λ5007 for the [O III] and Hα selected emitters
in Table 4 and 5. The electron temperature of the HII region is
also shown.
∼530,000 galaxy spectra in the SDSS and they are all
located nearby (z < 0.005) (SDSS DR3: Kniazev et al.
2003, SDSS DR4: Izotov et al. 2006a). At higher red-
shift, 17 metal-poor (7.8 < 12 + log(O/H)< 8.3) galax-
ies have been found at z ∼ 0.7 in the initial phase of the
DEEP2 survey of 3,900 galaxies and the Team Keck Red-
shift Survey of 1,536 galaxies (Hoyos et al. 2005). But
none of these galaxies satisfies the condition of XMPGs.
The present sample may be the first XMPGs at inter-
mediate redshift (z ∼ 1) whose oxygen abundances are
securely measured by the direct method. The narrow-
band method is very powerful for finding not only high-
redshift (z >> 5) galaxies, but also strong emission-line,
extremely metal-deficient galaxies at intermediate red-
shifts (z < 1).
Figure 15 shows the electron temperature sensi-
tive line ratio, [O III](λ4959+λ5007)/[OIII]λ4363 ver-
sus [O II]λ3727/[O III]λ5007. If we have an estimate of
the metallicity, as in the present case, we can use the
[O II]λ3727/[O III]λ5007 ratio to estimate the ionization
parameter q. The ionization parameter q is defined as
the number of hydrogen ionizing photons passing through
a unit area per second per unit hydrogen number den-
sity (Kewley & Dopita 2002). For the low metallicity
systems with strong [O III]λ4363 lines, we can see from
Figure 15 that [O II]λ3727 is very weak compared to
[O III]λ5007 with values ranging downwards from 0.3.
Assuming the metallicity is less than 0.2 Z⊙ this would
place a lower bound of q = 108 cm s−1 on the ioniza-
tion parameter based on the Kewley & Dopita (2002)
model. The higher metallicity objects have stronger
[O II]λ3727/[O III]λ5007 which, while in part due to the
metallicity, also requires these objects to have lower ion-
ization parameters suggesting we are seeing an evolution-
ary sequence.
6.2. Hα emitters
Among 13 Hα emitters in our spectroscopic sample,
only 3 NB912 selected emitters have Hβ fluxes ade-
quate (> 15σ) for the purpose of metallicity measure-
ments. Their Te[O III] and oxygen abundances were mea-
sured using the direct method described above, and are
shown in the Table 5 together with the data for the
[O III] emitters. The [O II]λ3727 line of ID266 is out-
side the Keck/DEIMOS wavelength coverage. In order
to derive an upper-limit on the metallicity, we assumed
[O II]λ3727/[O III]λ5007 = 1.0, which is the maximum
value in our sample (see, Fig. 15). All our Hα emitters
are metal poor (Zupper < 0.45 Z⊙), but none of them are
XMPGs.
6.3. Composite Spectrum
As can be seen in Figure 15 the objects with low
[O II]λ3727/[O III]λ5007 have relatively uniform values
of ([O III]λ5007+λ4959)/[OIII]λ4363 and similar metal-
licities. Given the relatively low signal to noise of the
individual spectra it therefore seems of interest to form a
composite spectrum. Such a spectrum will have weight-
ings on the lines which depend on the individual ion-
ization parameters and metallicity but will give a rough
estimate of the average metallicity and temperature of
the sample.
In Figure 16 we show the composite spectrum of
the 8 objects with [O II]λ3727/[OIII]λ5007 less than
0.1. The [O III]λ4363 is now strongly detected with a
value of 16.7 ± 2.1 or eight sigma. The mean temper-
ature is 19, 500 ± 1, 500 K and the average abundance
12+log(O/H)=7.5± 0.1 and the mean rest frame equiv-
alent width of Hβ is 57Å. The results are similar to the
values obtained by averaging the individual analyses of
the eight objects.
7. MORPHOLOGIES
The morphology of the USELs may give us a clue to the
mechanism of their high star formation activity (SFR ∼
a few M⊙/yr) and star formation history; what has trig-
gered the star formation −mergers, gas infall, or galactic
winds? High resolution HST/ACS images are available
for GOODS-North (GOODS-N) region (Giavalisco et al.
2004) which is one of our survey fields. There are 17
NB816 selected USELs in the GOODS-N, and 16 in the
NB912 sample. Figures 17 and 18 show thumbnails of
the NB816 and NB912 selected USELs in the GOODS-N
field (respectively) with each thumbnail 12.′′5 on a side.
The white dashes point to the largest galaxy. We used
continuum broadband images to show underlying stellar
populations: HST/ACS B, V, z′-band images were used
for NB816 emitters, and B, V, I-band for NB912 emit-
ters. High-redshift Lyα emitters (z >> 5) are very red
because of the continuum absorption below Lyα emission
caused by neutral hydrogen in the intergalactic medium.
We do not have spectra for most of the USELs in the
GOODS-N field yet, and none of the USELs in GOODS-
N have metallicity measurements either due to lack of
spectra or low spectral S/N. But we can qualitatively
argue that the USELs at intermediate redshift (z < 1)
exhibit widespread morphologies from relatively compact
high surface brightness objects to very diffuse low surface
brightness ones.
8. DISCUSSION
The present emitters differ from the local dwarf HII
galaxies in a large number of ways though they appear
much more similar to the XMPGs found in the SDSS
samples. They are much more luminous, have large
[O III]/[O II] ratios, and they are a relatively high frac-
tion (about 10% by number from Figure 13) of other
12 Kakazu et al.
3600 3800 4000 4200 4400 4600 4800 5000 5200
REST WAVELENGTH (A)
Fig. 16.— Composite spectrum of the 8 emitters with [O II]λ3727/[O III]λ5007 less than 0.1. The eight spectra have simply been summed
without weighting. The lower panel shows the stronger lines and the upper the continuum and the weaker lines. The stronger emission
lines are labelled and marked with the vertical dotted lines.
Fig. 19.— The oxygen abundance versus the absolute rest frame B magnitude for the [O III] selected samples (red squares). One sigma
errors are shown for the oxygen abundances and one sigma lower limits are shown with upward pointing arrows. The solid line shows the
(Skillman et al. 1989) relation for the nearby dwarf irregulars. As with the local XMPGs (filled circles, Kniazev et al. 2003; Kewley et al.
2007 and GRB hosts (open squares, Stanek et al. 2006; Kewley et al. 2007, the present galaxies are much more luminous at a given
metallicity than the local irregulars. Metal-poor luminous galaxies (but not XMPGs) at z ∼ 1 from Hoyos et al. 2005. are shown as
triangles. A few of our emitters may have oxygen abundances comparable to the most metal-deficient galaxis, I Zw 18 (12 + log O/H =
7.17± 0.01, Thuan & Izotov 2005) and SBS 0335-052W (12 + log O/H = 7.12 ± 0.03, Izotov et al. 2005).
THE YOUNGEST GALAXIES 13
galaxy populations at these redshifts. Taken together
this suggests that we are seeing much more massive
galaxies in the early stages of formation and, since we
need these to have relatively long lifetimes in order to
understand their frequency, that we are seeing objects
undergoing continuous star formation rather than single
starbursts. For the case of constant star formation with a
standard Salpeter IMF a forming galaxy can have equiv-
alent widths above 30Å for 109 yr (Leitherer et al. 1999)
which would allows us to understand the observed num-
ber density of strong emitters relative to the total galaxy
population.
In this type of model we would expect the metal-
licity to grow with time and that higher metallicity
galaxies would have higher continuum magnitudes and
lower equivalent widths in Hβ. We plot the absolute
rest frame B magnitudes versus the Oxygen abundance
in Figure 19. As with the case for the local XMPGs
found in the SDSS (filled circles, Kniazev et al. 2003;
Kewley et al. 2007) and the metal-poor galaxies (7.8 <
12 + log O/H < 8.3) at z ∼ 1 (triangles, Hoyos et al.
2005), the present emitters (red squares) are much more
luminous at a given metallicity than is found for the
local dIrrs (solid line, Skillman et al. 1989). Further-
more there does indeed seem to be a trend to higher
continuum luminosities at higher metallicity consistent
with ongoing star formation raising the luminosity. Re-
cently Kewley et al. (2007) reported the similarity be-
tween XMPGs and long duration GRB hosts ; they share
similar SFRs, extinction levels, and both lie in a similar
region of the luminosity-metallicity diagram. Our sam-
ple metal-deficient galaxies, which also lie in the region of
GRB hosts, may be additional support of the connection
between XMPGs and GRB hosts.
Of the six galaxies with continuum magnitudes
brighter than -18 all but one have metallicities or
lower limits which would place them near or above
12+log(O/H)=8 while the lower luminosity galaxies pri-
marily have 12+log(O/H) in the range 7.1 − 7.8. If we
assumed that the metallicity were a simple linear func-
tion of the age then the more luminous galaxies would be
several times older than the less luminous ones which is
not quite enough to account for the luminosity increase
(see e.g. Leitherer et al. 1999) suggesting that the en-
richment process may be more complex. However, the
accuracy of our current metallicity measurements may
be inadequate for measuring the lowest metallicities in
the sample and we could be underestimating the amount
of metallicity evolution.
The relation between the metallicity and the Hβ equiv-
alent width is shown in Figure 20. There clearly is a large
scatter in metallicity at all equivalent widths suggesting
that the star formation may be episodic with periods in
which bursts of star formation enhances the Hβ equiva-
lent widths in objects where previous star formation has
raised the metallicity.
With better spectra and more accurate metal estimates
we should be able to refine these tests and also determine
whether the number density of objects versus metallic-
ity is consistent with that expected in a simple growth
model. Perhaps even more importantly as larger spectro-
scopic samples are obtained we should be able to deter-
mine if there is a floor on the metallicity corresponding
10 100 1000
EW (H Beta)
Fig. 20.— The oxygen abundance versus the rest frame Hβ equiv-
alent width for the [O III] selected samples. One sigma errors are
shown for the oxygen abundances and one sigma lower limits are
shown with upward pointing arrows. The dotted line shows the
metallicity of 1Zw-18.
to the enrichment in the intergalactic gas. Within the er-
rors we have yet to find an object with lower metallicity
than the lowest metallicity local galaxies but this could
easily change as the observations are improved.
9. SUMMARY
We have described the results of spectrscopic observa-
tions of a narrowband selected sample of extreme emis-
sion line objects. The results show that such objects are
common in the z = 0 − 1 redshift interval and produce
about 5-10% of the star formation seen in ultraviolet or
emission line measurements at these redshifts. A very
large fraction of the strong emitters are detected in the
[O III]λ4363 line and oxygen abundances can be mea-
sured using the direct method. The abundance primar-
ily lie in the 12+log(O/H) range of 7−8 characteristic of
XMPGs.
The results suggest that we are seeing early chemical
enrichment of startup galaxies at these redshifts which
are forming in relatively chemically regions. As we de-
velop larger samples of these objects and improve the
accuracy of their abundance estimates we should be able
to test this model and to determine if there is a floor on
the metallicity of the galaxies.
We are indebted to the staff of the Subaru and Keck
observatories for their excellent assistance with the obser-
vations. We acknowledge Subaru/SuprimeCam support
astronomer, Hisanori Furusawa, for his help over several
years with the observations. Y. Kakazu acknowledges
invaluable advice from Lisa J. Kewley and Roberto Ter-
levich on metallicity measurements. This work was sup-
ported in part by NSF grants AST04-07374 (LLC) and
AST06-87850 (EMH), and Spitzer grant JPL 1289080
(EMH).
14 Kakazu et al.
REFERENCES
Ajiki, M. et al. 2003. AJ, 126, 2091, astro-ph/0307325
Ajiki, M., et al. 2006, PASJ, 58, 113, astro-ph/0511471
Capak, P. et al. 2004, AJ, 127, 180, astro-ph/0312635
Corbin, M. R., Vacca, W. D., Cid Fernandes, R., Hibbard, J. E.,
Somerville, R. S., & Windhorst, R. A. 2006, ApJ, 651, 861,
astro-ph/0607280
Cowie, L. L., Songaila, A., Hu, E. M., & Cohen, J. G. 1996, AJ,
112, 839, astro-ph/9606079
Faber, S. M. et al. 2003, Proc. SPIE, 4841, 1657
Garnett, D. R. 1992, AJ, 103, 1330
Giavalisco, M. et al. 2004, ApJ, 600, L93, astro-ph/0309105
Hoyos, C., Koo, D. C., Phillips, A. C., Willmer, C. N. A., &
Guhathakurta, P. 2005, ApJ, 635, L21, astro-ph/0511066
Hu, E. M., Cowie, L. L., & McMahon, R. G. 1998, ApJ, 502, L99,
astro-ph/9801003
Hu, E. M., Cowie, L. L., Capak, P., Hayashino, T., & Komiyama,
Y. 2004, AJ, 127, 563, astro-ph/0311528
Hu, E. M., Cowie, L. L., Kakazu, Y., & Capak. P. 2007, in
preparation
Izotov, Y. I., Thuan, T. X., & Guseva, N. G., 2005, ApJ, 415, 87,
astro-ph/0506498
Izotov, Y. I., Papaderos, P., Guseva, N. G., Fricke, K. J., &
Thuan, T. X. 2006a, A&A, 454, 137, astro-ph/0604234
Izotov, Y. I. 2006b, in ASP Conf. Ser. 353, Stellar Evolution at
Low Metallicity: Mass Loss, Explosions, Cosmology, ed.
H. J. G. L. M. Lamers, N. Langer, T. Nugis, & K. Annuk (San
Francisco: ASP), 349
Izotov, Y. I., Stasińska, G., Meynet, G., Guseva, N. G., & Thuan,
T. X. 2006c, A&A, 448, 955, astro-ph/0511644
Kennicutt, R. C., Jr. 1998, ARA&A, 36, 189, astro-ph/9807187
Kewley, L. J., & Dopita, M. A. 2002, ApJS, 142, 35,
astro-ph/0206495
Kewley, L. J., Brown W. R., Geller, M. J., Kenyon, S. J., &
Kurtz, M. J. 2007, ApJ, in press, astro-ph/0609246
Kniazev, A. Y., Grebel, E. K., Hao, L., Strauss, M. A.,
Brinkmann, J., & Fukugita, M. 2003, ApJ, 593, L73,
astro-ph/0307401
Kobulnicky, H. A., & Kewley, L. J. 2004, ApJ, 617, 240,
astro-ph/0408128
Kobulnicky, H. A., et al. 2003, ApJ, 599, 1006, astro-ph/0310346
Kunth, D., & Östlin, G. 2000, A&A Rev., 10, 1, astro-ph/9911094
Leitherer, C. et al. 1999, ApJS, 123, 3, astro-ph/9902334
Lilly, S. J., Carollo, C. M., & Stockton, A. N. 2003, ApJ, 597,
730, astro-ph/0307300
Ly, C. et al. 2007, ApJ, 657, 738, astro-ph/0610846
Maier, C., Lilly, S. J., Carollo, C. M., Meisenheimer, K.,
Hippelein, H., & Stockton, A. 2006, ApJ, 639, 858,
astro-ph/0511255
Melbourne, J., & Salzer, J. J. 2002, AJ, 123, 2302,
astro-ph/0202301
Miyazaki, S., et al. 2002, PASJ, 54, 833, astro-ph/0211006
Oey, M. S. 2006, in ASP Conf. Ser. 353, Stellar Evolution at Low
Metallicity: Mass Loss, Explosions, Cosmology, ed.
H. J. G. L. M. Lamers, N. Langer, T. Nugis, & K. Annuk (San
Francisco: ASP), 253
Oke, J. B. 1990, AJ, 99, 1621
Osterbrock, D. E. 1989, Astrophysics of gaseous nebulae and
active galactic nuclei (Mill Valley, CA: University Science
Books)
Pagel, B.E.J., Simonson, E. A., Terlevich, R. J., & Edmunds, M.
G. 1992, MNRAS, 255, 325
Papovich, C. et al. 2004, ApJ, 600, L111, astro-ph/0310888
Pilyugin, L. S. & Thuan, T. X. 2005, ApJ, 631, 231
Popescu, C. C., Hopp, U., Hagen, H. J., & Elsaesser, H. 1996,
A&AS, 116, 43, astro-ph/9510127
Pustilnik, S. A., Engels, D., Kniazev, A. Y., Pramskij, A. G.,
Ugryumov, A. V., & Hagen, H.-J. 2006, Astronomy Letters, 32,
228, astro-ph/0508255
Salzer, J. J., et al. 2000, AJ, 120, 80, astro-ph/0004074
Sargent, W. L. W., & Searle, L. 1970, ApJ, 162, L155
Schiminovich, D., et al. 2005, ApJ, 619, L47, astro-ph/0411424
Schlegel, D. J., Finkbeiner, D. P., & Davis, M. 1998, ApJ, 500,
525, astro-ph/9710327, used to estimate dust extinctions from
source location by the NASA/IPAC Infrared Science Archive
Dust Reddening and Extinction Tool
(http://irsa.ipac.caltech.edu/applications/DUST/)
Seaton M. J. 1975, MNRAS, 170, 475
Skillman, E. D., Kennicutt, R. C., & Hodge, P. W. 1989, ApJ,
347, 875
Stanek, K. Z. et al. 2006, Acta Astron., 56, 333, astro-ph/0604113
Stasińska, G., Schaerer, D., & Leitherer, C. 2002, Ap&SS, 281,
335 McMahon, R. G. 2003, MNRAS, 342, 439,
astro-ph/0302212
Stern, D., Bunker, A., Spinrad, H., & Dey, A. 2000, ApJ, 537, 73,
astro-ph/0002239
Stockton, A., & Ridgway, S. E. 1998, AJ, 115, 1340,
astro-ph/9801056
Thuan, T. X., Izotov, Y. I. 2005, ApJS, 161, 240,
astro-ph/0507209
Tran, K.-V. H., Lilly, S. J., Crampton, D., & Brodwin, M. 2004,
ApJ, 612, L89, astro-ph/0407648
Turnshek, D. A., Bohlin, R. C., Williamson II, R. L., Lupie,
O. L., Koornneef, J., & Morgan, D.H. 1990, AJ, 99, 1243
Ugryumov, A. V., et al. 1999, A&AS, 135, 511
Wilson, G., Cowie, L. L., Barger, A. J., & Burke, D. J. 2002, AJ,
124, 1258, astro-ph/0203168
Wyder, T. K., et al. 2005, ApJ, 619, L15, astro-ph/0411364
http://arxiv.org/abs/astro-ph/0307325
http://arxiv.org/abs/astro-ph/0511471
http://arxiv.org/abs/astro-ph/0312635
http://arxiv.org/abs/astro-ph/0607280
http://arxiv.org/abs/astro-ph/9606079
http://arxiv.org/abs/astro-ph/0309105
http://arxiv.org/abs/astro-ph/0511066
http://arxiv.org/abs/astro-ph/9801003
http://arxiv.org/abs/astro-ph/0311528
http://arxiv.org/abs/astro-ph/0506498
http://arxiv.org/abs/astro-ph/0604234
http://arxiv.org/abs/astro-ph/0511644
http://arxiv.org/abs/astro-ph/9807187
http://arxiv.org/abs/astro-ph/0206495
http://arxiv.org/abs/astro-ph/0609246
http://arxiv.org/abs/astro-ph/0307401
http://arxiv.org/abs/astro-ph/0408128
http://arxiv.org/abs/astro-ph/0310346
http://arxiv.org/abs/astro-ph/9911094
http://arxiv.org/abs/astro-ph/9902334
http://arxiv.org/abs/astro-ph/0307300
http://arxiv.org/abs/astro-ph/0610846
http://arxiv.org/abs/astro-ph/0511255
http://arxiv.org/abs/astro-ph/0202301
http://arxiv.org/abs/astro-ph/0211006
http://arxiv.org/abs/astro-ph/0310888
http://arxiv.org/abs/astro-ph/9510127
http://arxiv.org/abs/astro-ph/0508255
http://arxiv.org/abs/astro-ph/0004074
http://arxiv.org/abs/astro-ph/0411424
http://arxiv.org/abs/astro-ph/9710327
http://irsa.ipac.caltech.edu/applications/DUST/
http://arxiv.org/abs/astro-ph/0604113
http://arxiv.org/abs/astro-ph/0302212
http://arxiv.org/abs/astro-ph/0002239
http://arxiv.org/abs/astro-ph/9801056
http://arxiv.org/abs/astro-ph/0507209
http://arxiv.org/abs/astro-ph/0407648
http://arxiv.org/abs/astro-ph/0203168
http://arxiv.org/abs/astro-ph/0411364
THE YOUNGEST GALAXIES 15
TABLE 2
NB816 Emission-Line Sample
No. RA(2000) Dec(2000) N(AB) Z(AB) I R V B z
1 40.115555 −1.694722 23.92 25.28 24.84 25.44 −99.00 −99.00 −1.0000
2 40.116665 −1.617361 24.26 25.88 25.52 25.74 −99.00 −99.00 −1.0000
3 40.138332 −1.405639 23.49 24.94 24.70 25.22 −99.00 −99.00 0.6343
4 40.174721 −1.704750 24.27 25.46 25.35 25.82 −99.00 −99.00 0.6355
5 40.183056 −1.495417 24.58 25.28 25.85 26.86 −99.00 −99.00 5.6886
6 40.216946 −1.494805 24.80 26.08 26.31 26.14 −99.00 −99.00 0.2416
7 40.250832 −1.744639 23.51 24.35 24.33 24.54 −99.00 −99.00 −1.0000
8 40.276112 −1.518139 24.72 24.78 25.53 25.75 −99.00 −99.00 −1.0000
9 40.276390 −1.623250 24.55 25.36 25.36 25.76 −99.00 −99.00 −1.0000
10 40.284168 −1.453583 21.82 23.20 23.02 22.56 −99.00 −99.00 0.2480
11 40.298889 −1.447389 24.32 25.95 25.42 25.71 −99.00 −99.00 −1.0000
12 40.304165 −1.391694 20.88 22.35 22.06 22.32 −99.00 −99.00 −1.0000
13 40.306946 −1.638500 24.91 25.67 25.72 26.02 −99.00 −99.00 −1.0000
14 40.311111 −1.535111 24.08 25.93 25.49 25.71 −99.00 −99.00 0.6240
15 40.318333 −1.548222 24.60 25.03 25.50 25.96 −99.00 −99.00 −1.0000
16 40.318890 −1.430889 23.37 23.77 24.32 24.42 −99.00 −99.00 1.1804
17 40.319168 −1.446333 23.60 23.85 24.43 24.56 −99.00 −99.00 1.1873
18 40.320835 −1.778028 20.70 21.57 21.60 21.86 −99.00 −99.00 −1.0000
19 40.324165 −1.409972 24.29 24.38 25.18 25.45 −99.00 −99.00 −1.0000
20 40.326111 −1.709805 23.24 24.75 24.62 24.96 −99.00 −99.00 −1.0000
21 40.336388 −1.570194 24.26 24.73 25.06 25.33 −99.00 −99.00 −1.0000
22 40.337223 −1.388194 24.81 27.37 26.71 26.77 −99.00 −99.00 −1.0000
23 40.337502 −1.658306 24.89 25.20 25.76 26.05 −99.00 −99.00 −1.0000
24 40.340279 −1.689472 24.55 26.37 26.07 26.51 −99.00 −99.00 −1.0000
25 40.340557 −1.551889 24.99 25.70 25.93 26.35 −99.00 −99.00 −1.0000
26 40.340832 −1.371222 22.42 23.43 23.29 23.16 −99.00 −99.00 −1.0000
27 40.340832 −1.516250 24.89 25.06 25.69 25.78 −99.00 −99.00 −1.0000
28 40.341110 −1.493500 23.37 24.62 24.55 24.23 −99.00 −99.00 −1.0000
29 40.341389 −1.484139 24.48 25.46 25.49 25.73 −99.00 −99.00 −1.0000
30 40.342777 −1.599528 24.65 25.31 25.50 25.85 −99.00 −99.00 −1.0000
31 40.343056 −1.438833 23.17 24.60 24.54 24.54 −99.00 −99.00 0.6226
32 40.347500 −1.403833 24.86 27.45 26.16 26.35 −99.00 −99.00 0.6324
33 40.349724 −1.598472 23.11 23.87 24.11 24.46 −99.00 −99.00 1.1956
34 40.356388 −1.515722 24.94 26.19 26.18 26.42 −99.00 −99.00 −5.0000
35 40.372223 −1.390361 23.85 24.68 24.72 24.62 −99.00 −99.00 −1.0000
36 40.373611 −1.722528 24.93 25.85 25.74 25.88 −99.00 −99.00 −1.0000
37 40.377777 −1.701889 23.45 23.82 24.26 24.70 −99.00 −99.00 −1.0000
38 40.388054 −1.697361 23.96 24.32 24.79 24.90 −99.00 −99.00 −1.0000
39 40.388889 −1.573361 24.79 25.13 25.65 26.07 −99.00 −99.00 −1.0000
40 40.394444 −1.521389 22.73 23.69 23.84 24.06 −99.00 −99.00 0.6292
Note. — Magnitudes are measured in 3′′ diameter apertures. An entry of ‘−99’ indicates that
no excess flux was measured. −1.0000 in the redshift column means no spectroscopic data were
obtained for the object. This is a sample table showing the first entries of the electronic version
of the table that will accompany the published paper.
16 Kakazu et al.
TABLE 3
NB912 Emission-Line Sample
No. R.A. (J2000.0) Decl. (J2000.0) N(AB) z′(AB) I R V B zspec
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10)
1 40.131668 −1.408361 23.86 25.08 25.89 25.97 −99.00 −99.00 0.8371
2 40.133888 −1.575222 24.76 25.92 26.13 25.84 −99.00 −99.00 1.4498
3 40.148056 −1.593555 23.21 24.62 25.53 26.11 −99.00 −99.00 0.8207
4 40.148335 −1.725417 24.76 25.97 26.98 26.39 −99.00 −99.00 0.8111
5 40.150833 −1.737556 23.31 24.35 24.88 25.32 −99.00 −99.00 0.8269
6 40.153332 −1.536833 23.64 25.00 26.18 26.87 −99.00 −99.00 0.8301
7 40.156387 −1.765833 21.93 23.05 23.51 23.74 −99.00 −99.00 −1.0000
8 40.165833 −1.580056 24.94 26.20 25.92 25.84 −99.00 −99.00 1.4482
9 40.183334 −1.389583 21.77 22.79 23.33 23.52 −99.00 −99.00 0.8325
10 40.184444 −1.596444 23.09 24.50 25.30 25.55 −99.00 −99.00 0.8293
11 40.193611 −1.690083 24.32 25.40 25.83 25.69 −99.00 −99.00 0.8266
12 40.194168 −1.373722 23.93 24.93 25.21 25.35 −99.00 −99.00 −1.0000
13 40.194721 −1.373917 24.87 25.90 26.30 26.47 −99.00 −99.00 −1.0000
14 40.196667 −1.378333 24.07 25.42 26.05 26.18 −99.00 −99.00 −1.0000
15 40.202221 −1.584472 24.22 25.30 25.81 26.19 −99.00 −99.00 0.8289
16 40.203335 −1.471861 24.77 26.20 25.78 25.54 −99.00 −99.00 0.3965
17 40.214722 −1.519917 23.14 24.40 25.38 25.94 −99.00 −99.00 0.8288
18 40.220276 −1.753778 24.34 25.41 26.40 26.29 −99.00 −99.00 −1.0000
19 40.220833 −1.388556 23.24 24.36 25.13 24.99 −99.00 −99.00 −1.0000
20 40.226944 −1.542111 23.02 24.47 25.77 25.48 −99.00 −99.00 0.8208
21 40.229168 −1.720889 24.99 27.85 27.11 · · · −99.00 −99.00 6.4800
22 40.229443 −1.376472 23.75 24.98 25.72 24.66 −99.00 −99.00 −1.0000
23 40.245834 −1.578972 24.61 25.82 27.22 26.91 −99.00 −99.00 0.8285
24 40.280556 −1.421056 24.82 25.86 26.13 25.91 −99.00 −99.00 −1.0000
25 40.290833 −1.746361 23.89 25.16 25.81 25.03 −99.00 −99.00 0.0000
26 40.323891 −1.697667 24.73 25.80 26.47 25.71 −99.00 −99.00 −1.0000
27 40.330833 −1.612389 23.03 24.68 26.06 25.47 −99.00 −99.00 0.0000
28 40.339722 −1.395361 23.87 25.54 26.98 25.45 −99.00 −99.00 0.3889
29 40.346668 −1.448305 23.93 24.98 25.37 25.48 −99.00 −99.00 0.8274
30 40.382500 −1.554056 23.52 24.87 25.60 25.08 −99.00 −99.00 0.3930
31 40.393055 −1.713694 24.94 26.83 26.70 26.50 −99.00 −99.00 −1.0000
32 40.398888 −1.466417 23.87 24.91 24.83 24.64 −99.00 −99.00 1.4590
33 40.403889 −1.530167 24.88 26.08 27.44 27.30 −99.00 −99.00 0.8223
34 40.409443 −1.369222 21.52 23.44 24.47 24.04 −99.00 −99.00 −1.0000
35 40.411667 −1.691417 24.41 25.50 25.68 25.55 −99.00 −99.00 −1.0000
36 40.424999 −1.454028 21.91 22.93 23.42 23.41 −99.00 −99.00 −1.0000
37 40.430000 −1.501111 23.13 24.21 24.76 24.92 −99.00 −99.00 0.8267
38 40.446110 −1.676139 24.91 26.09 26.37 26.04 −99.00 −99.00 −1.0000
39 40.478889 −1.534278 24.87 26.06 25.79 25.56 −99.00 −99.00 0.3861
40 40.506111 −1.755111 22.73 24.27 25.24 25.19 −99.00 −99.00 −1.0000
41 40.511665 −1.596944 24.73 25.96 25.55 25.67 −99.00 −99.00 −1.0000
42 40.518055 −1.666139 22.47 23.73 24.37 24.45 −99.00 −99.00 −1.0000
Note. — Magnitudes are measured in 3′′ diameter apertures. An entry of ‘−99’ indicates that no
excess flux was measured. −1.0000 in the redshift column means no spectroscopic data were obtained for
the object. This is a sample table showing the first entries of the electronic version of the table that will
accompany the published paper.
TABLE 4
Line fluxes and Oxygen Abundance for L816 selected emitters
Object f([OIII]5007) f([OIII]4959) f([OIII]4363) f([OII]3727) Te[OIII] 12+log(O/H)
[OIII] emitters
31 513.6 ± 24.0 222.3 ± 11.4 < 6.60 54.4 ± 4.62 < 1.19 > 8.09
40 577.9 ± 21.6 191.3 ± 8.05 9.40 ± 3.40 140.9 ± 6.22 1.14 < 1.34 < 1.53 7.86 < 8.03 < 8.25
51 401.5 ± 12.6 146.9 ± 5.52 9.40 ± 4.50 < 2.39 1.19 < 1.55 < 1.90 7.51 < 7.62 < 7.94
76 464.4 ± 10.5 191.3 ± 4.86 < 2.90 344.5 ± 8.07 < 0.95 > 8.55
118 492.6 ± 29.7 193.9 ± 12.9 34.4 ± 12.8 11.3 ± 2.61 2.16 < 3.08 < 4.32 6.93 < 7.16 < 7.44
195 335.0 ± 21.4 129.5 ± 10.2 24.0 ± 10.9 97.1 ± 9.71 2.02 < 3.17 < 4.86 6.78 < 7.06 < 7.44
206 597.1 ± 19.5 204.1 ± 7.41 21.6 ± 9.20 < 1.72 1.48 < 1.97 < 2.48 7.42 < 7.55 < 7.84
208 658.0 ± 30.9 249.8 ± 12.3 15.6 ± 8.00 < 22.1 1.16 < 1.56 < 1.93 7.67 < 7.85 < 8.22
223 242.9 ± 15.3 83.3 ± 7.53 23.7 ± 18.9 < 10.1 1.45 < 4.64 < 19.62 6.14 < 6.61 < 7.53
252 466.8 ± 9.32 157.9 ± 3.57 9.20 ± 4.00 139.0 ± 3.71 1.16 < 1.45 < 1.72 7.68 < 7.87 < 8.14
Note. — Only emitters with > 15σ Hβ fluxes are listed. All fluxes are normalized by their f(Hβ) and multiplied by
100. 1σ upper limits are listed for [OII]3727 flux below 3σ and [OIII]4363 below 1σ. The units of Te[OIII] are 10
−4 [K].
THE YOUNGEST GALAXIES 17
Fig. 17.— HST/ACS (B, V, z’) composite images of NB816 emitters in the GOODS-N field with overlaid object IDs from Table 2 and
redshifts, where known. Fields are 12.′′5 on a side.
TABLE 5
Line Fluxes and Oxygen Abundances for L912 selected emitters
Object f([OIII]5007) f([OIII]4959) f([OIII]4363) f([OII]3727) Te[OIII] 12+log(O/H)
[OIII] emitters
3 550.9 ± 12.9 187.9 ± 4.91 23.9 ± 7.90 8.6 ± 2.5 1.74 < 2.20 < 2.71 7.26 < 7.43 < 7.65
6 588.1 ± 35.1 216.0 ± 14.2 18.4 ± 11.0 52.0 ± 8.6 1.20 < 1.79 < 2.39 7.40 < 7.68 < 8.14
9 442.3 ± 15.3 154.7 ± 6.42 < 12.3 157.2 ± 6.88 < 1.70 > 7.70
10 490.0 ± 11.9 178.7 ± 5.14 13.7 ± 4.40 < 2.95 1.42 < 1.69 < 1.97 7.55 < 7.61 < 7.82
17 342.5 ± 20.0 129.7 ± 9.29 15.9 ± 9.40 < 7.72 1.41 < 2.26 < 3.28 7.03 < 7.22 < 7.70
20 418.7 ± 17.4 135.1 ± 6.96 16.8 ± 5.50 24.5 ± 2.45 1.69 < 2.11 < 2.57 7.18 < 7.36 < 7.58
239 202.4 ± 10.2 75.6 ± 6.40 < 8.20 190.6 ± 10.2 < 2.08 > 7.34
270 351.7 ± 15.1 149.7 ± 7.73 12.4 ± 3.40 30.7 ± 2.72 1.59 < 1.87 < 2.16 7.28 < 7.43 < 7.61
Hα emitters
52 589.1 ± 10.0 179.2 ± 3.42 18.3 ± 1.59 < 1.56 1.75 < 1.83 < 1.92 7.62 < 7.67 < 7.72
60 619.1 ± 33.5 206.7 ± 12.5 10.7 ± 7.77 48.4 ± 8.7 0.90 < 1.37 < 1.77 7.67 < 7.96 < 8.57
266 682.8 ± 10.3 217.7 ± 3.57 14.7 ± 2.42 ... 1.40 < 1.52 < 1.63 < 8.3
Note. — Same as Table 4 but for NB912 emitters. The [OII]λ3727 of ID266 is beyond the DEIMOS wavelength
coverage and thus was not being measured.
18 Kakazu et al.
Fig. 18.— HST/ACS (B, V, I) composite images of NB912 emitters in the GOODS-N field with overlaid object IDs from Table 3 and
redshifts, where known. Fields are 12.′′5 on a side.
