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UNIVERSITY OF HAWAI'I LIBRARY
A WIDE-FIELD SURVEY FOR HIGH-REDSHIFT QUASARS
A DISSERTATTON SUBMITTED TO THE GRADUATE DIVISION OF THE NTVERSITY OF HAWAI'I IN PARTIAL FULFILLMENT OF T HE
REQ UIREIVIENTS FOR THE DEGREE OF
DOCTOR OF PHILOSOPHY
IN
ASTRONOMY
MAY 2008
By Yuko K. M. Kakazu
Dissertation Commi ttee:
Esther M. Hll , Cha irperson Michael C. Liu
Antoinette S. Cowie Alan N. Stockton
Richard J . Wainscoat John G. Learned
We certify that we have read this dissertation and that, in our opinion, it is
satisfactory in scope and quality as a dissertation for the degree of Doctor
of Philosophy in Astronomy.
DISSERTATION COMMITTEE
Chairpe n
ii
© Copyright 2008 by
Yuko K. M. Kakazu All Rights Reserved
iii
Dedicated to my mother and father
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Acknowledgements
First and foremost, I would like to thank my supervisor, Esther M. Hu, for her guidance
and enthusiOBtic encouragements throughout my graduate school life in Hawaii. I am
particularly indebted to her patience and painstaking corrections on my English. Special
thanks go to Len Cowie. who WOB, in many respects, my second supervisor. His deep insight
and sound advice have made OBtrophysics fun for me. I am deeply grateful to Michael Liu
for his valuable advice and assistance in my studies of brown dwarfs. I also thank him for
his careful reading of the manuscript. I wish to thank Toni Cowie for her supervision on
my first 699 project, and Amy Barger for her continuous support and advice particularly
with my job-hunting.
It is my pleOBure to thank Alan Stockton, Richard Wainscoat, and John Learned for
their participation in the committee of this thesis. Special thanks to Richard for obtaining
near-IR data of the brown dwarfs candidates.
It is impossible to overstate my appreciation to Lisa Kewley, who hOB been my role model
of a female researcher. Her continuous encouragements and advice have been enormous in
the iDBt two years. Thank you, Lisa.
I am sincerely grateful to my former office mates, Peter Capak and Wei-Hao Wang
for their help and guidance on computer programming, observations, and data reduction.
Without their help, I would have 'drowned' in the massive data flow from Subaru, or, I
would have destroyed my machine after hours of unsuccessful compilation of programs.
The observations presented in this thesis were obtained with various telescope facilities
on Mauna Kea. I owe my most sincere gratitude to the observatory stuff, particularly
v
two support astronomers of Subaru/Suprime-Camj Komiyama-san and Furusawa-san. This
work would not have been possible without their devotion to the observatory.
On a more personal note. I warmly thank all my friends who have gotten me through
this long journey. Special thanks to my office mates, CJ and Ben for enduring with my
music, Yun-Han for being a great housemate, and my lovely "Hula sisters" - Miko, Yumi,
Nanase, and Keiko. I also wish to thank auntie Lei'ohu, Maydeen, Kumu Hula Eko and
Debbie for opening a door to Hawaiian culture for me.
It would be a crime if I did not thank the 'team Jewitt' - Dave Jewitt, Jing Li, Suu-Suu,
Pedro Lacerda, Nuno Peixinbo, and Jan Kleyna. They have been my enourmous supporters,
particularly in the final stage of my Ph.D., which was perhaps the most intense period in
my entire life. I sincerely thank Dave and Jing for feeding me and making sure that I was
alive and happy. Spending time with Suu-Suu has been wonderful refreshment especially
after hours of debugging and thesis writing. It was my great pleasure to see her growth
from a toddler to a girl who is now bilingual and even plays violin. I am truly grateful
to Pedro, Nuno, and Jan, whom I could always count on when I needed quirky humors
or commiserations. I especially thank my "Savior", Pedro, who fixed a bug in my thesis
presentation file which had happened 18 hours before the defense.
Finally, but not least, my deep gratitude goes to my whole family for their continuous
and unconditional support. Thanks to my mother, Miyoko Kakazu, for raising me and
supporting my decision to go abroad. Her optimistic ways of thinking have always been my
resort. I am truly grateful to my sister, Mariko, and my two brothers, Koji and Masato,
for their support and belief in me. The memory of my deceased father, Kosaburo Kakazu,
has been a persistent motivation, inspiration, and encouragement for my journey of life. I
hope that my father would have been proud to see my accomplishment now.
vi
Abstract
The present thesis reports the results from the Hawaii QUIlB8.l' and T dwarf survey (HQT
survey), which is a wide-field optical imaging survey conducted with Subaru/Suprime-Cam.
The HQT survey was designed to search for low-luminosity qUllB8.l'S (Ml450 < -22.5) at high
redshift (z > 5.7) as well as T dwarfs, both of which are selected by their very red optical
1- z' colors. We developed a new color selection technique using a narrowband N B816 filter
in order to break a well-known color degeneracy between qullB8.l'S and foreground M and
L dwarfs. The follow-up Keck/DEIMOS spectroscopy and near-ffi inIaging with various
instruments on Mauna Kea have demonstrated the effectiveness of our technique, and have
successfully revealed six faint T dwarfs (J < 20). These dwarfs are among the most distant
spectroscopically known (60 - 170 pc) and they provide an indirect support for the high
binary fraction at LIT transition. The non-detection of z > 5.7 qullB8.l'S in our survey is
consistent with the present picture of the cosmic reionization in which qullB8.l'S are negligible
contributor to the cosmic reionization. With our survey area coverage (9.3 deg2) and depths
(ZAB < 23.3), we were able to set strong constraints on the faint-end slope of the qUllB8.l'
luminosity function.
Majority of our candidate qullB8.l'S turned out to be strong emission line galaxies at
Z < 1, whose large equivalent widths and low metal contents suggest they are very young
systems which have just undergone starbursts within a few Myrs. In order to systematically
search for these Ultra-Strong Emission Line galaxies (USELs), we used narrowband selected
samples from Hu's ultra-deep multiwavelength data. The followup Keck/DEIMOS spectra
have revealed their high star formation density (5-10% of UV measurements at z = 0-1),
vii
which is a significant contribution at a epoch when cosmic star formation is in its peak.
Many of the USELs show [OIIIjA4363 auroral lines and about a dozen satisfy the criteria
for eXtremely Metal Poor Ga1axies (XMPGs). Our XMPGs are the most distant known
today. Our high yield rate of XMPGs suggests that narrowband method is powerful in
finding such populations. Strikingly, a few of our XMPGs have metallities close to the
most metal-poor galaxy. Our discovery indicates that ga1axies are still forming in relatively
chemically pristine sites at z ~ 1.
viii
Table of Contents
Acknowledgements . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .. v
Abstract . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .. vii
List of Tables . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .. xi
List of Figures ....................................... xiii
Chapter 1: Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .. 1
1.1 Motivation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .. 1
1.2 Current High-z Quasar Surveys. . . . . . . . . . . . . . . . . . . . . . .. 3
1.3 Hawaii Quasar and T dwarf survey (HQT survey). . . . . . . . . . . . . .. 9
1.4 Dissertation Outline . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .. 13
Chapter 2: Hawaii Quasars and T-dwarf (HQT) Survey I.
Method and Discovery of faint field ultra-cool dwarfs . . . . . . . . . . . . . . .. 16
2.1 Introduction.................................... 17
2.2 Color Selection Technique for High-z QSOs (z > > 5) and T dwarfs . . . .. 20
2.3 Observations ................................... 24
2.3.1 Subaru/Suprime-Cam Imaging Survey. . . . . . . . . . . . . . . .. 24
2.3.2 QSO and T dwarf Candidate Selection . . . . . . . . . . . . . . . .. 27
2.3.3 Follow-up Near-IR Imaging . . . • • • • • • . . . . . • • • • • • • .. 29
2.4 Analysis...................................... 32
2.4.1 Classification with Colors . . . . . . . . . . . . . . . . . . . . . . .. 32
2.4.2 Classification of illtracool dwarfs with near-ffi spectra . . . . . . .. 36
ix
2.4.3
2.4.4
Distance Measurements
Proper Motion . . . . .
2.4.5 Number Counts and Galactic Scale Height. .
2.5 Conclusion ...................... .
Chapter 3: Hawaii Quasars and T-dwarf (HQT) Survey II.
Constraints on the z = 6 Quasar Luminosity Function .
3.1 Introduction ..................... .
3.2 QSO Candidate Selection in HQT Supcime-Cam field.
3.3 Follow-up Observations ..... .
3.3.1 Near-IR Imaging Follow-up
3.3.2 Keck/DEIMOS Spectroscopic Follow-up
3.4 Results.......................
3.4.1 Strong Emission Line Galaxies at z < 1
3.4.2 Limits on the Quasar Luminosity Function
3.5 Discussion and Conclusion . . . . . . . . . . . . .
Chapter 4: Extremely Low-Metallicity Galaxies at z < 1
4.1 Introduction ........ .
4.2 The Narrow Band Selection
4.3 Spectra ...... .
4.4 Flux Calibrations . .
4.4.1 Narrow Band Fluxes - Primary Fluxes ...
4.4.2 Line Fluxes from the Spectra . . . . .
4.4.3 Balmer Ratios ............ .
4.4.4 Final Flux Calibration for Metallicity Analysis
4.5 Star Formation History ................ .
4.5.1 Ha and [OIII]A5007 Luminosity Functions .. .
4.5.2 Star Formation Rates
4.6 Galaxy Metallicities .....
x
37
39
41
44
60
61
63
65
65
68
69
69
76
81
87
88
90
93
101
102
104
108
108
109
109
112
114
4.6.1 [OlII] emitters
4.6.2 Ha. emitters . .
4.6.3 Composite Spectrum .
4.7 Morphologies .
4.8 Discussion.
4.9 Summary .
Chapter 5: Summary and Future Work
5.1 Summary ..
5.2 Future Work
A new color criteria 5.2.1
5.2.2 Quasar contribution to cosmic reionization
5.2.3 Multiwavelength Studies of USELs at z < 1
xi
114
118
118
120
120
127
147
147
149
149
149
150
List of Tables
1.1 Summary of High-Redshift (z > 6) Quasar Surveys
1.2 Known High-Redshift (z > 5.7) Quasars
1.2 Known High-Redshift (z > 5.7) Quasars
2.1 Survey Area and Ultra-cool Dwarf Number Counts
2.2 Coefficients of Polynomical Fit to Colors of Ultra-Cool Dwarfs ...... .
2.3 Photometric Properties of Candidates .. . . . . . . . . . . . . . . . . . . .
2.4 Log of Near-ffi Imaging Observation . . . . . . . . . . . . . . . . . . . . . .
2.5 Classification and Distance Measurements . . . . . . . . . . . . . . . . . . .
2.6 Proper Motion Measurements
6
7
8
51
52
53
54
54
55
3.1 Candidate High-redshift (z > 5.7) Quasars ...•.............. 65
3.2 Log of Near-ffi Imaging Observation . . . . . . . . . . . . . . . . . . . . .. 68
3.3 KeckII/DEIMOS Spectroscopy Observations of HQT Quasar Candidates . 69
3.4 Faint-end Slope and Break Magnitude for the z = 6 Quasar LF at Several
Confidence Levels. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .. 78
4.1 Narrowband Survey Area Coverage . . . . . . . . . . . . . . . . . . . . . .. 92
4.2 Narrowband Survey Area Coverage . . . . . . . . . . . . . . . . . . . . . .. 127
4.3 NB816 selected USEL Sample . . . . . . . . . . . . . • . • . . . . . • • . .. 128
4.3 NB816 selected USEL Sample . . . . . . . • . . . . . . . . . . . . . . . . .. 129
4.3 NB816 selected USEL Sample . . . • . . . . • • . . . • • . . . . . . . . . .. 130
xii
4.3 NB816 selected USEL Sample . . . . . . . . . . . . . . . . . . . . . . . . .. 131
4.3 NB816 selected USEL Sample . . . . . . . . . . . . . . . . . . . . . . . . .. 132
4.3 NB816 selected USEL Sample . . . . . . . . . . . . . . . . . . . . . . . • .• 133
4.4 NB912 selected USEL Sample . . . . . . . . . . . . . . . . . . . . . . . . .. 134
4.4 NB912 selected USEL Sample . . . . . . . . . . . . . . . . . . . . . . . . .• 135
4.4 NB912 selected USEL Sample . . . . . . . . . . . . . . . . . . . . . . . . .• 136
4.4 NB912 selected USEL Sample . . . . . . . . . . . . . . . . . . . . . . . . .. 137
4.4 NB912 selected USEL Sample . . . . . . . . . . . . . . . . . . . . . . . . .. 138
4.4 NB912 selected USEL Sample . . . . . . . . . . . . • . . . . . . . . . . • .. 139
4.4 NB912 selected USEL Sample . . . . . . . . . . . . . . . . . . . . . . . . .. 140
4.5 Line fluxes and Oxygen Abundance for NB816 selected USELs .. . . . .. 141
4.6 Line Fluxes and Oxygen Abundances for NB912 selected USELs .. . . .. 142
xiii
List of Figures
1.1 Histogram of M141lO distribution of known z > 5.7 quasars . . . . • . . . . .
l.2 Limiting magnitude and area coverage of major surveys . . . . . . . . . . .
l.3 Predicted quasar number at survey redshift range 5.8 :5 z :5 6.4 . . . . . . .
2.1 A model quasar spectrum at z = 6 and T dwarf spectrum ......... .
2.2 N B816 - z' VB. [- z' color-color diagram . . . . . . . • . . . . . . . . . . .
2.3 HQT Survey Area as a function of limiting magnitude . . . . . . . . . . . .
2.4 Optical Suprime-Cam and near-m MOmCS images of HQT-9 ...... .
2.5 Subaru/CISCO JH spectrum of HQT-2 ................... .
2.6 Subaru/CISCO J H spectrum of HQT-3 ................... .
2.7 Keck/NIRSPEC J-band spectrum of HQT-11 ..............•..
2.8 I - z' VB. z' - J diagram for candidates . . . . . . . . . . . . . . . . . . . .
2.9 z' - J VB. J - H diagram . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.10 z' - J VB. J - K diagram
2.11 Expected number of L dwarfs and T dwarfs as a function of Galactic scale
height ...................................... .
2.12 Expected number of L dwarfs and T dwarfs as a function of field Galactic
5
11
12
21
25
28
29
33
34
35
46
47
48
49
latitude .. . . . . • . . . . . . . . . . . . . . . . . . . . . . . . . • . . . .. 50
3.1 N B816 - z' VB I - z' diagram for candidate quasars . • . . . . . • . . . .. 66
3.2 Images of the HQT candidate quasars . . . • • . . . . . . . . . . . . . . .. 67
xiv
3.3 I - z' versus. z' - J diagram for C8Jldidate quasars . . . . . . . . . . • . .. 70
3.4 Keck/DEIMOS spectra of HQT Jl64554.25+65301O.93 . . . . . . . . . . .. 72
3.5 Keck/DEIMOS spectra of HQT J221432.35+oo3119.05 . . . . . . . . . . .• 73
3.6 Keck/DEIMOS spectra of HQT Jl63552.45+660706.71.
3.7 Keck/DEIMOS spectra of HQT J222037.35+014558.05
3.8 Constraints on the faint-end slope and break magnitude
3.9 Quasar luminosity function at z = 6 . . . . . . . . . • . . . . . . . . . . . .
3.10 032 VB. R23 emission line diagnostics .................... .
3.11 OIII/OIl VB. NelII/HB .........................•....
4.1 Sample Ha emitter with SuprimeCam filter profile ....•.........
4.2 Continuum mag versus. Narrowband magnitude .............. .
4.3 Sample Spectrum of Ha emitter in NB912 sample. . • . . . . . . . . . . . .
4.4 Sample spectrum of OIII emitter in NB816 sample
4.5 Sample spectrum of OIII emitter in NB912 sample
74
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94
95
96
97
4.6 Sample spectrum of 011 emitter in NB816 sample. . . . . . . . . . . . . .. 98
4.7 Redshift Distribution of all narrowband selected emitters .......... 100
4.8 Distribution of Rest frame equivalent width of Ha and 0111 lines . . • . .. 103
4.9 OIII.M959/ A5OO7 flux ratio as a function of f(OIIIAS007) .......... 105
4.10 Ratio of fluxes measured from images and spectra .............. 106
4.11 Balmer Ratios versus Ha flux . . . . . . . . . . . . . . . . . . . . . . . . .. 107
4.12 Luminosity function of Ha emitters. . . . . . . . . . . . . . . . . . . • . •. 110
4.13 Luminosity function of olIr emitters . . . . . . . . . . . . . . . . . . . . .. 111
4.14 Star Formation Rate History .......................... 113
4.15 Metallicity versus 0IlA3727/0IIIA5OO7 . . . . . . . . . . . . . . . . . . . •. 117
4.16 Composite Spectrum . . . . . . . . . . . . . . . . . • . . . . . . . . . . . .. 119
4.17 HST / ACS image of NB816 emitters in GOODS field . . . . . . . . . . . .. 121
4.18 HST / ACS image of NB912 emitters in GOODS field . . . • . . . . . . . .. 122
xv
4.19 Oxygen abundance versus M(B). . . . . . . . . . . . . . . . . . . . . . . .. 123
4.20 The Oxygen abandance versus M(B) . . . . . . . . . . . . . . . . . . . . .. 124
4.21 Oxygen abundance versus Rest frame H,8 equivalent width. . . . . . . . .. 126
xvi
1.1 Motivation
Chapter 1
Introduction
One of the fundamental questions in current cosmology is the era and energy sources of
·cosmic reionization", which is a landmark event in the cosmic history after which the
"When WIlB the universe reionized, ending the 'dark ages'?" "What are the sources
responsible for reionization - a first generation of galaxies or AGNs?" .
Identifying the objects at high-redshift (z » 5) is a direct way of probing the source of
cosmic reionization and the structure formation in the early universe. Ultra-deep surveys
using narrowband filters centered on the Lya emission line of high-z galaxies have proved
to be highly successful, and have revealed more than 100 star-forming galaxies at z ~ 5.7,
z ~ 6.6 (see review by Hu et al. 2006), and now even at z ~ 7.0 (lye et al. 2006; Ota et
al. 2007). Additional high-z galaxies have also been found using the Lyman break method,
which selects UV-bright galaxies using broad-band colors (Bunker et al. 2006; Bouwens
& llIingworth 2006 and references therein). Very faint galaxies at 7 < z < 10 have been
identified using strong gravitational lensing magnification by clusters, though their faintness
makes it difficult to spectroscopically confirm that they are truly z > 7 galaxies (Ellis et al.
2001; Stark et al. 2007 and references therein). These discoveries of z > > 5 galaxies and
the filamentary spatial distribution of Lya emitting galaxies at z ~ 5.7 over angular scales
of ~ 60 Mpc (Hu et al. 2004) suggest that early structure formation hBB already talren
1
place by a redshift of z ~ 6, implying significant field-to-field variations exist even at these
high redshifts.
However, the studies of quasars at high-redshift have been less well advanced than the
faint galaxy studies until very recently. This is because of the extreme rarity of high-z
quasars, and the overwhelmingly large number of cool dwarfs (mainly type M and L), that
resemble high-z quasars in the optical colors. Discriminating between these two different
populations requires either deep near-ffi imaging or direct spectroscopy, both of which are
time intensive to perform for these large survey fields.
The use of the reddest optical band, z', in combination with wide J-band images from
2MASS (Two Micron All Sky Survey; Skrutskie et al. 1997) have resulted in the discovery
of more than 20 z ~ 5.8 luminous quasars in the Sloan Digital Sky Survey (SDSS) which
covers ~ 8,000 deg2 of sky (e.g., Fan et al. 2oo6a). Their absorption spectra have shown
that the z ~ 6 is the tail-end of the reionization epoch (e.g., Becker et al. 2001; Fan et al.
2006b).
However, the relatively shallow survey depth means that at z ~ 6 the SDSS only probes
to M 1450 = -27, about three magnitudes brighter than the average luminosity of AGNs
at lower redshifts (Pei 1995). Therefore, the bright surveys such as SDSS cannot place
constraint on the faint-end slope of the luminosity function, where majority of UV ionizing
flux arises. In order to have a complete picture about quasar evolution and its relation to
the reionization history, a deeper survey than the SDSS is needed. Hence as my Ph.D. thesis
project, we have conducted a wide-field (~ 14 deg2) survey for high-z quasars (z ~ 5.8) at
significantly fainter limits than the SDSS (z' AB < 24). Our goal is to place constraints on
the faint-end slope of the quasar luminosity function at high-z, which will lead to the better
understandings of the quasar contribution to the early UV background radiation, and the
formation of supermassive black holes.
2
1.2 Current High-z Quasar Surveys
At the time when my Ph.D. thesis proposal was approved (2004 May), there were only
seven very luminous quasars (Ml450 < -27) known at z > 5.8 and they were all found by
the SDSS survey (Fan et aI. 2000, 2001, 2003). Now the number of bright SDSS quasars
has increased to about 20 (Fan et aI. 2004, 2006; Goto et aI. 2006). Many attempts of
finding lower luminosity quasars have been carried out as well, however, majority of them
have resulted in negative results. Willott et aI. (2005) inlaged a 3.8 deg2 area down to
z'AB = 23.35 in the Canada-France High-redshift Quasar Survey (CFHQS), but did not find
any quasars. In the 2 Ms X-ray exposure of the Chandra Deep Field-North, Barger et aI.
(2003) found no quasars at z > > 5 with z~B < 25.1. Similarly, Cristiani et aI. (2004)
reports the non-detection of high-z AGNs based on deep X-ray imaging of the Hubble Deep
Field-North and Chadra Deep Field-South. The lack of z » 5 quasars in these deep
surveys have already demonstrated that quasars are not likely the major source of cosmic
reionization (e.g., Shankar & Mathur 2007).
The recent development in wide-field near-IR cameras, combined with the deeper optical
imaging data, have finally began to shed light on fainter, hence more "normal", quasars at
high-z. Just within the past year, about a dozen z > 5.7 quasars whichare~ 1-2 magnitude
fainter than the SDSS luminous quasars, have been found. SDSS deep survey, which covers
an area of 260 deg2 to z'AB.lim = 21. have identified five quasars with Ml450 "" -26 (Jiang et
aI. 2007). One of their quasars at z = 5.86 was independently discovered by Venemans et aI.
(2007) in the UKIDSS survey using Y filter. Cool et aI. (2006) found a quasar at z = 5.85
from an area of 7.7 deg2 in the NOAO Deep Wide-Field Survey Bootes Field (NDWFS;
Jannuzi & Dey). Their quasar was selected from the Spitzer mid-infrared data, and is about
half a magnitude fainter than the bright SDSS quasars. More recent CFHQS survey that
is shallower (z'AB = 22.5) but covers significantly wider area (~ 4OOdeg2) than their initial
search, Willott et aI. (2007) discovered four quasars at z > 6. Table 1.1 summarizes the
3
previous and current surveys for hlgh-z qUI!88.I'S, and Table 1.2 shows a list of currently
known z > 5.7 qul!88.I'S.
Ongoing surveys mentioned above (e.g., SDSS deep; CFHQS; UKIDSS) are expected to
find more qul!88.I'S at fainter limits than the SDSS main survey. However, qul!88.I'S found by
these surveys are still more than two magnitudes brighter than the characteristic magnitude
of qul!88.I'S at z = 2, which is Mi450 = -25 (Croom et aI. 2004). The bright SDSS survey
has shown that the quasar number density declines from z = 2 to z = 6, suggesting a
fainter MO at z = 6 (Fan et aI. 2004). Figure 1.1 shows the Ml450,AB distribution of
currently known z > 5.7 qUI!88.I'S, and it is clearly shown that majority of known qUl!88.I'S
are extremely bright. Therefore current surveys are unlikely to probe the faint-end of the
quasar luminosity function at which the bulk of the quasar population reside. It is therefore
necessary to conduct a new survey for hlgh-z quasars which goes down to fainter limits, yet
covers sufficiently large area in order to better sample quasars and to overcome the cosmic
variance that are already known to be present at z > > 5 from the galaxy studies (e.g., Hu
et aI. 2004).
4
Ul 10 ... c Ul C ::l 8 I-0
" 1li A 6 l-N
i
I , C ~ c 4 c I ~ -c ... 2 I-CD ~
E ::l Z 0
-28 -27 -26 -25 -24 M'450, AS
Figure 1.1 Histogram of Ml450 distribution of known z > 5.7 qllllSarS. Break magnitude of quasar luminosity function at z = 2 is indicated as a red dotted line (Mi450 AB = -25). It can be seen that all known z > 5.7 quasars are brighter than this magnitud~.
5
CD
Table 1.1. Summary of High-Redshift (z > 6) QUaBar Surveys
Survey Name Published Planned Color selection Limiting mag M1460 RedshIft Expected Identllled Reference Area Area range Number QSOs
[deg2[ [deg') [AB) [AB) [per deg')
HQTSurvey 9.3 20.0 I z' NB816 Zl = 23.3 < -22.4 5.8-6.4 1.8 0 1
PrevIous Surveys
SDSS main 8000 8000 i' Zl J Z' = 20.2 < -26.7 5.8 - 6.4 0.006 19 2 CFHQS deep 3.83 3.83 i' z' J z' = 23.4 < -22.4 5.8-6.4 1.8 0 3
Ongoing Surveys
SDSS deep 260 300 i' Zl J z' = 21.0 < -25.9 5.8-6.4 0.02 5 4 UKIDSS LAS 106 4,100 it z' Y J y= 20.8 < -25.8 5.8-7.2 0.01 1 5 CFHQS 400 900 " Zl J Z· = 22.5 < -24.4 5.8-6.4 0.40 4 6 Palomar QUEST 15,000 i' z' J z· = 19.5 < -27.4 5.8- 6.4 0.001 0 7
Note. - Our survey (top row), previous surveys (middle row), and ongoing surveys (bottom row) for z > 6 quasars are Usted. Color selection teclmique used In current z > 6 q1UlSW" surveys require both optical (i', z') and nesr-IR (J) data, while our method only needs optical data (I, N 8616, z') to separate out quasars from their contaminants. The expected z I'V 6 quasar number per square degree is computed from the most recent SDSS luminosity function (Jiang et aI. 2007). If a break In the quasar luminosity function occurs wltbln the survey magnitude range, which is the ceae for low-z quaaars, the SDSS luminosity function will overpredlct tbe number of quasars expected In tbe surveys deeper than tbe SDSS main and deep surveys.
References. - (I) Hawaii Quasar and T dwarf survey [HQT survey) - this work (Kakszu et aI. 2007 In preparation); (2) Sloan Digital Sky Survey northern Galactic cap survey [SDSS main) - Fan et aI. (2006; Paper IV); (3) Canada-France High-redshit Quasar Survey [CFHQS deep) - Wlllott et aI. (2005); (4) Sloan Digital Sky Survey southern deep survey [SDSS deep) - Jiang et aI. (2007); (5) UKIRI' Infrared Deep Sky Survey Large Area Survey [UKIDSS LAS) - Venemanns et aI. (2007); (6) Canada-France High-redshit Quasar Survey [CFHQS] - WlIIott et aI. (2007); (7) Palomer Quest survey - Djorgovskl et aI. (2008).
Table 1.2. Known High-Redshift (z > 5.7) Quasars
Survey Object redshift Ml450 i Z J H K referenre
CFHQS J232908-030158 6.43 -25.23 > 25.08 21.76± 0.05 21.56± 0.25 W07 SDSS main J114816.64+525150.3 6.42 -27.82 23.30± 0.30 20.0l± 0.09 18.25± 0.05 F03 SDSS main J103027.1O+052455.0 6.28 -26.89 23.23± 0.43 20.05± 0.10 18.87± 0.10 FOI SDSS main JlO4845.05+463718.3 6.23 -27.55 22.38± 0.19 19.86± 0.09 18.4O± 0.05 F03 SDSS main JI62331.81+3112oo.5 6.22 -26.67 24.52± 0.62 20.09± 0.10 19.15± 0.10 F04 CFHQS JOO3311-012524 6.13 -25.03 24.82± 0.15 22.41± 0.08 21.58± 0.20 W07 SDSS main JI25051.93+313021.9 6.13 -26.87 22.14± 0.18 19.53± 0.08 > 16.50 F06 CFHQS JI50941-174926 6.12 -26.98 23.11± 0.05 20.26± 0.02 19.68± 0.10 W07 SDSS deep J231546.57-002358.1 6.12 -25.38 24.90± 0.28 20.88± 0.08 19.94± 0.08 J07 SDSS main JI60254.18+422822.9 6.07 -26.82 22.78± 0.38 19.89± 0.10 18.46± 0.05 F04 .... SDSS deep J030331.40-001912.9 6.07 -25.43 23.92± 0.23 20.85± 0.07 19.46± 0.10 J07 SDSS deep J205406.49-000514.8 6.06 -26.11 23.30± 0.22 20.72± 0.09 19.18± 0.06 J07 SDSS main JI63033.90+401209.6 6.05 -26.11 23.38± 0.35 20.42± 0.12 19.38± 0.10 F03 SDSS deep J035349.72+010404.4 6.05 -26.49 24.03± 0.30 20.54± 0.08 18.55± 0.06 J07 CFHQS JI64121 +375520 6.04 -25.48 23.69± 0.20 21.3l± 0.04 21.24± 0.14 W07 SDSSmain J113717.73+354956.9 6.01 -26.85 22.57± 0.30 19.54± 0.07 18.4l± 0.05 F06 SDSS main J081827.4O+172251.8 6.00 -27.14 22.19± 0.22 19.60± 0.08 18.54± 0.05 F06 SDSS main JI30608.26+035626.3 5.99 -27.19 22.58± 0.05 19.47± 0.05 18.77± 0.10 FOI SDSS main J084119.52+290504.48 5.96 -26.86 22.54 19.84 20.8 GOO SDSS main JI33550.81+353315.8 5.93 -26.58 22.67± 0.99 20.1O± 0.11 18.97± 0.05 F06 SDSS main JI41111.29+121737.4 5.93 -26.75 23.43± 0.37 19.63± 0.07 18.95± 0.05 F04 UKIDSSLAS J020332.38+001229.2b 5.86 -26.20 23.70± 0.20 20.90± 0.10 19.08± 0.10 99.00±99.00 19.2l± 0.08 V07 SDSS deep J020332.39+001229.3b 5.85 -25.72 23.72± 0.22 20.87± 0.10 19.05± 0.08 J07 SDSS deep J000552.34-000655.8c 5.85 -25.82 23.40± 0.34 20.54± 0.10 19.87± 0.10 J07 NDWFS JI42516.3+325409 5.85 -26.03 21.57± 0.06 20.68± 0.06 COO SDSS main J084035.09+562419.9 5.85 -26.40 22.43± 0.34 19.76± 0.10 19.00± 0.10 F06 SDSS main J000552.34-000655.8c 5.85 -26.46 23.4O± 0.34 20.54± 0.10 19.87± 0.10 F04 SDSS main JI43611.74+500706.9 5.83 -26.28 22.76± 0.28 20.00± 0.12 19.04± 0.10 F06
00
Table 1.2-Continued
Survey Object redsbift Ml450 z J H K reference
SDSS main J083643.85+oo5453.3 5.82 -27.88 21.04± 0.08 18.74± 0.05 17.89± 0.05 FOl SDSS main JOOO239.39+ 255034.8 5.BO -27.66 21.47± 0.11 IB.99± 0.05 > 16.50 F04 SDSS main JI04433.04-012502.2 5.BO -27.50 21.Bl± 0.19 19.23± 0.07 17.02 ± 0.04 FOO SDSS main J092721.B2+200123.7 5.79 -26.55 22.12± 0.17 19.BB± O.OB 19.0l± 0.10 F06
"This QSO was discovered by Gato et aI. (2007) using the SDSS fourth public data release. Photometry errors are not available.
bThese two are the same QSO tbat was first reported by the UKIDSS survey. and was later independently found by the SDSS deep survey.
<These two are tile same QSO. This QSO was dlscovered by Fan et aI. (2004: Paper IV). and was recovered by Jiang et aI. (2007).
Note. - Known z > 5.7 quasars are lIsted in the order of redsbift as of October 2007. Optical magnitudes in (i' and z') are given in AB system, whlIe near-infrared magnitudes (J, H. K) are iu Vega-based system. Note tbat I-band of NDWFS JI42516.3+325409 (Cool et aI. 2oo7) is a Cousins I filter.
References. -
• SDSS main survey - FOO (Fan et aI. 2000). FOI (Fan et aI. 2001: Paper I). F03 (Fan et aI. 2003: Paper II), F04 (Fan et aI. 2004: Paper ill). F06 (Fan et aI. 2006; Paper IV).
• SDSS deep survey - J07 (Jiang et aI. 2oo7).
• NOAO Deep and Wide Field Survey [NDWFS] - C06 (Cool et aI. 2006).
• UKIDSS Large Area Survey [ULAS] - V07 (Venemans et aI. 2007).
• Canada-France High-z QSO survey [CFHQS] - W07 (Willott et aI. 2007).
1.3 Hawaii Quasar and T dwarf survey (HQT survey)
In order to move beyond current uncertainties in the faint-end of the q1l8B8.I' luminosity
function (LF) and the field-to-field variations due to cosmic variance, deeper surveys over
sufficiently wide areas are needed. Hence as my Ph.D. thesis project, we have conducted a
wide-field (~ 10 dlllf) survey for high-z quasars (z ~ 5.8) at significantly fainter limits than
these other surveys (~ 4 magnitudes deeper than the SDSS survey), using Suprime-Cam
on 8.2m Subaru telescope. High-z quasar surveys are known to be effective in identifying
ultra.cool dwarfs since they exhibit similarly red I - z' colors to high-z q1l8B8.I'S due to
strong KI absorption bands (e.g., Fan et al. 2006a). These ultra.cool dwarfs have ma9'JeB
between stars and planets, and therefore their studies lead to the better understanding of
the formation and evolution of stars and planets. Since the current surveys for ultracool
dwarfs are very shallow, they can probe only very nearby ones « 30 pc in case for the
SDSS). With the depth of our survey, we can probe T dwarfs out to 100 pc, and may
be able to constrain their Galactic distribution for the first time. Although the primary
purpose of our survey is to identify quasars at high-redshift, we therefore also examine T
dwarfs. Hereafter we refer to our survey as the Hawaii Quasar and T dwarf (HQT) survey.
Since quasars have larger characteristic luminosity, but smaller number density than
galaxies, quasar surveys are generally shallower but wider than the galaxy surveys. We
determine the survey limiting magnitude to be faint enough to probe quasars at significantly
fainter limits than other major quasar surveys, yet covering sufficiently large area to better
sample quasars. Figure 1.2 shows the depths VB. area coverage of the major high-z gala.xy
and quasar surveys. The depth of the HQT survey is similar to the CFHQS deep survey
(Willott et al. 2005), but covers four times larger area.
Expected Quasar numbers
Since the current z > 5 quasars are all very luminous, faint end of the quasar LF is poorly
determined. We therefore calculate the expected number of quasars in the HQT survey by
9
extrapolating the SDSS quasar LF to lower luminosity. In this analysis, we used a composite
quasar spectrum from Vanden Berk et aI. (2001) and applied Lya forest law of Songaila
& Cowie (2002) to compute K correction for quasars at the redshift range of our survey
(5.7 < z < 6.4). Figure 1.3 shows the resulting number of quasars per square degree as a
function of survey limiting magnitude in z'. Three LF models were used to estimate the
number of quasars. Solid line indicates the quasar LF of Fan et aI. (2001) derived from 39
luminous quasars at 3.6 < z < 5.0. Blue solid line represents the LF models constructed
based on nine z ~ 6 quasars found in SDSS main survey (Fan et aI. 2004). This shows the
steepening of the bright-end slope at z = 5 -> 6. Recently, Jiang et aI. (2007) measured
a z = 6 quasar LF using a total of 20 quasars found in the SDSS main and deep surveys.
Their model is shown as a red solid line, and is consistent with the LF of Fan et aI. (2004).
The z ~ 6 quasar space densities from SDSS main survey (Fan et aI. 2oo6a) and SDSS
deep survey (Jiang et aI. 2007) are plotted as squares. The discovery of one luminous quasar
in the UKIDSS survey appears to be consistent with the SDSS predictions (Venemans
et aI. 2007). The vertical lines indicate limiting magnitudes for the SDSS main survey
(zAB = 20.2), the SDSS deep survey (rAB = 21), the CFHQS survey (zAB = 22.5), and the
HQT survey (zAB = 23.5). We roughly estimate about two quasars at the depth of HQT
survey according to the quasar LF of Jiang et aI. (2007). The expected number for other
quasar surveys are summarized in Table 1.1. IT the break magnitude M* occurs within
the magnitude range of a survey, the number of quasars becomes less than the prediction.
However, we address that at the magnitude limit and field coverage of the HQT survey, even
non-detections provides a significant constraint on the faint-end of the luminosity function
at high-z quasars.
10
~ .ll
N
26
24 I-
HDF • COSMOS •
CFHQS ..Qeep
• ~QT CFHQS •
- -22
- 24
22 1-......................................... ...... .. ......... .... ...... ....... ... ............... ... ...... .. .
:'i a '" "
SDSS Deep • ~
- - 26
20 I-SDSS Main •
QUEST • -28
1 8 ~~~~~~~~~~~~~ .. ~, ~~ 10- 1 100 10 1 102 103 104 105
Survey Area [deg 2]
Figure 1.2 Limiting magnitude and area coverage of major surveys for z = 6 objects. Quasar surveys plotted here are summar ized in Table 1.1. The HQT survey is shown as a fed circle. The Q EST survey is a next generation all-sky survey for extremely luminous quasars (Djorkov ·ki et a1. 2008). Both the QUEST and CFHQS surveys are indicated by blue circles, since the area plotted here for these surveys are proposed values. Two galaxy su rvey fields, the COSMOS (Scovi lle et a1. 2007) and Hawaii Hubble Field North (Capak et a1. 2004) are also plotted for comparison to show that galaxy surveys generally probe to much fainter depths t han quasar surveys, however cover smaller area. Horizontal dotted line indicates a break magnitude of quasar LF at z = 2 (Mi450 = -25.0).
11
M -32 -30 -28 1 45~26 -24 -22
104 5.8 < z < 6.4
N _HQT Survey
C1' Q)
1:l 102 _._CFHQS
"- ...... SOSS deep '-Q) __ SOSS main
.D E 100 :J
Z
1:l luminosity Functions Q)
10- 2 _ Jiang . 1 01. (2007) ...., U Q) _ 'on .1 01. (2004) 0-X _'on .1 01. (2001)
W
10- 4
14 16 18 20 22 24 26 z' [AS]
Figure 1.3 Predicted quasar number at survey redshift range 5.8 :0; z :0; 6.4. The vertical lines indicate limiting magni tudes for the major quasar surveys described in Table 1.1; SDSS main survey (dashed line; z' < 20.2), SDSS deep survey (dotted line; z' < 21), CFHQS survey (dash-dot line; z' < 22), and HQT survey (pink solid line; z' < 23.3). Three quasar LF model are used to calculate the expected number of quasars per square degree; Jiang et a!. (2007) [red line], Fan et a!. (2004) [blue line], and Fan et a!. (2001) [black line]. In the HQT survey, we expect about 2 z ~ 6 quasar per deg2 , though the number is overestimated if the break magnitude is within the survey magnitude range. The predicted numbers for other major quasar surveys are listed in Table 1.1.
12
1.4 Dissertation Outline
We present the results from the Hawaii QU8B8.l' and T dwarf survey (HQT survey) which was
designed and conducted by Y. Kakazu as a dissertation project to search for low luminosity
quasars at high-redshift (z > > 5) as well as distant T dwarfs.
In Chapter 2, we describe the survey imaging data taken with the Suprime-Cam
mosaic CCD camera on the 8.2-m Subaru telescope. We introduce a unique color-selection
technique to separate out quasars from their major contaminants (M and L dwarfs) without
relying on direct spectroscopy or near-infrared imaging. Using this technique, we select
extremely red objects (I - Zl > 2.0) found in an area of 9.3 deg; in the HQT survey, with
the purpose of finding T dwarfs and quasars at the highest redshift range probed by the
reddest optical filter z'. This Chapter will be soon submitted to Astrophysical JournoJ. with
authors Kakazu, Y., Hu, E. M., Liu C. M., Wang, W.-H., Wainscoat, J. R., and Capak, L.
P.
In Chapter 3, we focus on moderately red objects (0.7 < 1-Zl > 2.0) to look for quasars
at z > 5.7. Keck/DEIMOS followup spectra of candidate quasars, however, showed that
none of them is z > > 5 quasar and majority of our candidates are strong emission line
galaxies at z < 1. Using the non-detection of high-z quasars, we constrain the faint-end
slope a and the break magnitude M* of the qU8B8.l' LF at z = 6. With the area coverage
and depths of our survey, our constraints are significantly tighter than the previous studies.
The strong emission line galaxies found by the HQT survey as high-z quasar candidates,
turned out to be an interesting class of galaxies which have just undergone major starbursts
within a few Myr. Motivated by our findings, we have conducted more systematic studies
of these galaxies using Hu's ultradeep narrowband data. In Chapter 4. we investigate the
physical properties of these galaxies. This Chapter has been published in the Astrophysical
Journal (vol 668, p853) with authors, Kakazu, Y., Cowie, L. L., & Hu, E. M. Finally in
Chapter 5, we briefly summarize our results and discuss future work.
13
References
[lJ Barger, A., et al. 2003, ApJ, 5841, 61
[2J Becker, R., et al. 2001, AJ, 122, 2850
[3J Bunker, A., et al. 2006, New Astronomy Reviews, 50, 94
[4J Bouwens, R. J., & Illingworth, G. D. 2006, Nature, 443, 189
[5J Cool, R. J., et al. 2006, AJ, 132, 823
[6J Cristiani, S., et al. 2004, ApJ, 6ooL, 119
[7J Croom, S. M., Smith, R. J., Boyle, B. J., Shanks, T., Miller, L., Outram, P. J., &
Loaring, N. S. 2004, MNRAS, 349, 1397
[8J Djorgovski, S. G., et al. 2008, Astronomische Nachrichten, 329, 263
[9J Ellis, R., et al. 2001, ApJ, 560, Ll19
[lOJ Fan, X., et al. 2000, AJ, 120, 1167
[l1J Fan, X., et al. 2001, AJ, 122, 2833 (Paper I)
[12J Fan, X., et al. 2003, AJ, 125, 1649 (Paper II)
[13J Fan, X., et al. 2004, AJ, 128, 515 (Paper III)
[14J Fan, X., et al. 2006, AJ, 131, 1203 (Paper IV)
14
[15] Fan, X., et aI. 2006, AJ, 132, 117
[16] Goto, T. 2006, MNRAS, 371, 769
[17] Hu, E. M., et aI. 2004, AJ, 127, 563
[18] Hu, E. M. & Cowie, L. L. 2006, Nature, 440, 7088, 1145
[19] lye, M., et aI. 2006, Nature, 443, 7108, 186
[20] Jannuzi, B. T., & Dey, A. 1999, in ASP Conf. Ser. 191, Photometric Redshits and
High-Redshift GaIaxies, ed. R. J. Weymann, L. J., Storrie-Lombardi, M. Sawicki, & R.
J. Brunner (San Francisco: ASP), 111
[21] Ota, K., et aI. 2007, submitted to ApJ, astro-ph/0707.1561
[22] Pei, Y. 1995, ApJ, 438, 623
[23] Shankar, F. & Mathur, S. 2007, ApJ, 660, 1051
[24] Skrutskie, M. F. et aI. 1997, in The Impact of Large Scale Near-IR Sky Surveys, ed.
F. GarzOn et aI. (Dordrecht: Kluwer), 25
[25] Songaila, A. & Cowie, L. L. 2002, AJ, 123, 2183
[26] Jiang, L. et aI. 2007, submitted to AJ, astro-ph/0708.2578
[27] Vanden Berk, et aI. 2001, AJ, 122, 549
[28] Venemans, B. P. et aI. 2007, MNRAS, 376, L76
[29] Willott, C. J., Delfosse, X., Forveille, T., Delorme, P., & Gwyn, S. D. J. 2005, ApJ,
633,630
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15
Chapter 2
Hawaii Quasars and T-dwarf (HQT) Survey I.
Method and Discovery of faint field ultra-cool
dwarfs
Abstract
The Hawaii Quasar and T dwarf survey (HQT Survey) is a wide-field, red optical survey
carried out with the Suprime-Cam mosaic CCD camera on the 8.2-m Subaru telescope.
The HQT survey is designed to search for low-luminosity (MABl450 < -23) quasars at
high-redshift (z > 5.7) as well as T dwarfs, both of which are selected by their very
red I - z' colors. We use an optical narrowband filter N B816 to break a well-known
I - z' color degeneracy between high-z quasars and foreground M and L dwarfs, which
are more numerous than quasars. In this Chapter, we report on the discovery of six faint
(19 :s: J :s: 20) ultracool dwarfs found over a ~ 9.3 deg2 area with limiting magnitude
of zAB :s: 23.3. These dwarfs were confirmed by near-IR imaging and/or spectroscopy
conducted at various facilities on Mauna Kea. With estimated distances of 60 - 170
pc, these are among the most distant spectroscopically confirmed brown dwarfs to date.
Limits on the proper motion of these ultracool dwarfs suggest that they are old members
of Galactic disk, though future follow-up observations are necessary to minimize the errors.
Our finding rate of ultracool dwarfs is within the model predictions. However, the large
16
brightening amplitude (~1 mag) previously reported for the LIT transition objects appear
to overpredict the numbers. Our results agree well with the recent relations of Liu et
al. (2007), and indirectly suggest a high binary rate for these transition objects. We also
examine how the survey field latitude affect the survey sensitivity to the vertical scale height
of ultracool dwarfs.
2.1 Introduction
Brown dwarfs are substellar objects with masses too low to sustain hydrogen-burning nuclear
fusion in their cores. They fill the mass and temperature gap between stars and planets.
The studies of these ultra-cool, low-mass objects therefore provide us with insight into the
formation of starlike objects as well as planets.
The last several years have seen tremendous progress in identifications of these ultracool
dwarfs thanks to the advent of wide-field surveys carried out on 1 - 2m-class telescopes:
the Sloan Digital Sky Survey (SDSS; York et al. 2000), the Two Micron All Sky Survey
(2MASS; Skrutskie et al. 1997), and the Deep Near-Infrared Survey of the Southern Sky
(DENIS; Epchtein et al. 1999). To date, over 100 brown dwarfs with spectral types T,
which are the coolest known substellar objects ( with Tefl ;§ 1400 K; Golimowski et al.
2004; Vrba et al. 2004), have been successfully discovered by these surveys (Chiu et al.
2006 and references therein). However, the shallow depths of these surveys mean that they
can ouly probe brown dwarfs in solar neighborhood. For example, the shallow SDSS survey
can ouly probe T dwarfs within 40 pc.
It is crucial to sample more distant brown dwarfs to measure their number density and
the distribution in the Galaxy. The recent advent of large-format near-ffi cameras on 4-m
class telescopes has made it possible to identify such population: Two T dwarfs, one located
at ~8Opc, were found in the United Kingdom Infrared Telescope (UKIRr) Infrared Deep
Sky Survey (UKIDSS; Lawrence et al. 2006) early data release (Kendall et al. 2007). Eight
17
new mid to late type T dwarfs (Lodieu et al. 2007) 8Jld one very late T dwarf (T8.5; Warren
et al. 2007) are further found in the UKIDSS Large Area Survey first data release.
These ultra.cool dwarfs emits ~ 90% of their energy at wavelength longward of Ittm with
a peak around 1.2ttm, 8Jld thus near-ill photometry 8Jld spectra are advantageous to identify
8Jld study them. However, a substeIlar object survey conducted ouly at near-ill wavelengths
is insensitive to late type L 8Jld early type T dwarfs, as the strong C~ absorption in 1.6ttm
(H-b8Jld) 8Jld 2.2ttm (K-b8Jld) makes these objects look similar to main sequence stars
8Jld asteroids in near-lR colors. On the other h8Jld, optical surveys with far-red b8Jldpasses
C8Jl search for T dwarfs over wide r8Jlge of temperatures. Moreover, the large field-of-view
of optical mosaic CCD cameras with good response at far-red wavelengths enable efficient
searches for these ultracool dwarfs, as the discovery of a number of T dwarfs in optical
surveys demonstrate (e.g., SDSS survey; Chiu et al. 2006, !fA deep survey; Liu et al.
2002).
Intriguingly, quasars at high-redshift (z ~ 5.8) exhibit similarly red 1-:1 colors to L 8Jld
T dwarfs due to strong absorption caused by the intervening intergalactic medium (lGM).
Indeed, many low-mass stars and brown dwarfs have been found by high-z quasar surveys
where these dwarfs were considered as contaminants (eg, Fan et al. 2001; Sharp et al. 2004;
Willott et al. 2005). Cool dwarfs, particularly M and L types, are much more common than
high-z quasars at the typical depths probed by these surveys. For example, in the shallow
SDSS survey, the number density of M/L dwarfs is more than 30 times larger than that of
z > 5.7 quasars (Fan et al. 2003). T dwarfs are rarer than M/L dwarfs, being 5 times more
numerous than quasars. This demonstrates the rarity of quasars and the large numbers of
cool dwarfs, particularly M and L type dwarfs that are identified by optical color selection.
In order to break the optical color degeneracy between quasars and ultra.cool dwarfs,
near-ill follow-up photometry and/or direct spectroscopy has been used. Despite the recent
advent of large near-ill mosaic cameras, however, it is still not so efficient to conduct wide
field near-ill imaging to the required depths. Hence in this Chapter, we introduce an
18
alternative method which uses a narrowband filter NB816 mounted on Subaru/Suprime
Cam, to separate quasars from ultra.cool dwarfs without relying on near-ffi data.
Equipped with this color selection technique, we have conducted a wide-field, deep
optical survey for high-z quasars (z > 5.7) and T dwarfs using Suprime-Cam on 8.2-m
Subaru telescope. Hereafter we refer to this as the HQT (Hawaii Quasar and T dwarf)
survey. The HQT survey covers an area of about 10 deg2 to a limiting magnitude of
z' ~ 24. Our survey is about four maguitudes deeper than the SDSS northern Galactic cap
survey, which now covers 7,500 d~ area (Fan et aI. 2006). Our goals are: (1) to identify
low-luminosity quasars (Ml450 < -22.5) at high-z (z > 5.7), which are considered to form
the bulk of quasar population at high-redshift, and (2) to identify faint, field T dwarfs and
understand their number density and distribution in the Galaxy.
In this Chapter, we describe a survey color selection procedure and report the discovery
of several faint (J > 19) ultracool dwarfs. Other results (e.g., constraints on the quasar
luminosity function, discovery of strong emission-line galaxies with extremely high ionization
parameters at z < 1) will be presented in the following Chapters. The outline of this Chapter
is as follows: In §2, we show our survey color selection technique. Survey observations with
Suprime-Cam and follow-up observasions using various ffi instruments on Mauna Kea are
described in §3. In §4.1, we present classification of our candidates based on their optical
and near-IR colors. Near-ffi spectra were obtained for three candidates. In §4.2, we show
that these are indeed ultra.cool dwarfs, and estinIate their spectral types from the spectra.
Distances and proper motions of the newly identified T dwarfs are measured and described
in §4.3 and §4.4. In §4.5, we calculate the expected number of dwarfs using a sinIple analytic
model and examine their Galactic vertical scale height. We also investigate the effect of
field latitude on the their scale height for future surveys. Finally we summarize our results
in §5. Optical maguitudes in this dissertation are given in the AB system (Oke et aI. 1990),
while near-ffi magnitudes are in the Mauna Kea Observatory near-ffi systenJ (MKO-Nffi
system; Tokunaga & Vacca 2005; Leggett et aI. 2006) which is a Vegarbased photometry
system.
19
2.2 Color Selection Technique for High-z QSOs (z» 5) and
Tdwarfs
At z > 5.7, the Lya emission line of quasars begin to move out of the I band and the
strong IGM H I absorption below Lya emission suppresses quasar I band fluxes, making
them very red in 1- z'. Thus z > 5.7 quasars are selected as I-dropouts. Ultracool dwarfs
of spectral types M, L, and T are more numerous than high-z quasars. These cool dwarfs
possess similarly red I - z' colors due to their cool atmospheres, mimicking high-z quasars
in optical color selection. Figure 2.1 illustrates the spectral features of a high-z quasar and
an ultracool dwarf.
In order to separate quasars from these foreground cool dwarfs, z' - J colors have
traditionally been used in optical quasar surveys (e.g., SDSS; Fan et al. 2001), since quasars
are bluer due to declining power-law continuum above Lya, whereas cool dwarfs are red.
However it requires a lot of telescope time to carry out follow-up J-band imaging for all
candidates selected by the I -dropout technique. For example, our survey with a limiting
depth of z' ~ 24 requires J-band imaging data as deep as J ~ 22 in order to detect M and
L dwarfs, and remove them from the quasar candidates. This means very intensive near-ffi
observations (> 100 hours even with the the Wide Field Camera (WFCAM) on UKIRI',
one of the largest ffi cameras) to image the whole survey area (~ 10 deg2).
We therefore use an optical narrow band filter N B816 as a way of breaking the color
degeneracy between quasars and stars without relying on near-ffi data. This filter was
originally designed to capture Lya emission from high-z galaxies (e.g. Hu et al. 2004;
Ajiki et al. 2(03). However as we show here, it is also powerful for the clean selection
of high-z quasars and T dwarfs in I-band drop-out surveys. The N B816 imaging with
Suprime-Cam on Subaru can reach the required depth in ten times less observation time
when compared with the UKIRI' /WFCAM. Another advantage of the N B816 filter is that
it is a1most centrally located in the Cousins I filter, thus we can calibrate the NB816 data
using Cousins I data. The N B816 filter has a width of ~ 120 A (FWHM), and is centered
20
C ::J
>. '-o
30
20
'-...., 10 .D 'o ..........
-- z = 6 Model QSO --T2 dwarf
NB816 z' J
-10L-____ -L ____ ~~~ __ ~~ ____ ~ ____ ~
0.6 0.8 1 .0 1.2 1.4 Wavelength [JLm]
Figure 2.1 A model quasar spectrum at z = 6 (solid line) and an IRTF / Spex and Keck/ LRIS spectrum ofT2 dwarfSDSS1254-0122 from Burgasser et al. (2002) and Cusbing et aI. (2005) (dotted line). The transmission profi les of Subaru / Suprime-Cam optical fi lters (I , N 8816, and z') and UKIRT/ UFTI J-band fi lter are plotted below the spectra as reference. The strong Lya emission line of quasars begins to move out of the I -band at th is reclshift. Ultracool dwarfs show a strong infrared excess in comparison to a z = 6 quasars and therefore the z - J color has been used in current higb-z quasar su rveys to separate quasars from cool dwarfs.
21
at a nominal wavelengths of 8150 A in a region of low sky background between OH bands.
The nominal specifications for the Subaru filters may be found at the instrument website 1
and are also described in Ajiki et al. (2003).
Figure 3.1 presents the NB816 - z' versus 1- z' color-color diagram for known stars
and brown dwarfs with the predicted quasar track as a function of redshift overplotted.
The quasar color (blue solid line) was computed from the SDSS composite quasar spectrum
of Vanden Berk et al. (2001) at redshifts 5.7 - 6.4. We first extrapolated the spectrum
with a Iv ~ v-O•46 power law continuum below Lya, and then applied the Songaila &
Cowie (2002) IGM transmission law for a given redshift of quasar spectrum. Finally,
we convolved results through the Suprime-Cam filter response profiles, which include the
ceo quantum efficiency, atmospheric absorption, and optics effects. To compute stellar
colors, we convolved the stellar spectral flux library of Pickles (1998) with the Suprime
Cam filter response profile. For ultracool stars, we used Sandy Leggett's online database 2
and IRTF jSpeX spectra from Cushing et al. (2005). An eight-order polynomial fit to the
synthesized colors was computed as a function of spectral type of ultra-cool dwarfs and is
shown as the green track in Figure 3.1. Table 2.2 lists the coefficients of the polynomial fits
to optical and near-IR colors of dwarfs of M - T types. The optical-near-IR colors (z' - J)
and pure near-IR colors of dwarfs were computed in the same manner as optical colors, but
for near-IR bandpasses we used Mauna Kea consortium filter set installed in UKIRT jUFTI.
The fits to pure near-IR colors (J - Hand J - K) are made from both synthesized and
observational data from Knapp et al. (2004).
As can be seen in this diagram, z > 5.7 quasars are selected as N B816-dropouts (i.e.,
have red N B816 - z, colors). On the other hand, M and L dwarfs, which are the dominant
population of I-dropouts at the typical depths probed by quasar surveys (e.g., Fan et al.
2001), have bluer NB816 - z' colors. Hence we can separate quasars from M and L dwarfs
by using their NB816 - z' color. The reason why high-z quasars are selected as NB816-
dropouts is easily understood. At z ~ 5.7, Lya emission from the quasar comes into the
·http://www.naoj.org/ObservingfInstruments/SCam/sensitivity.html 'http://www.jach.hawaii.edu/pub/ukirt/skl
22
NB816 filter and at z ~ 5.8 it moves out of the NB816 filter. This quasar redshift transition
z = 5.7 -> 5.8 is marked by a huge increase in the NE816 - z color term. At z > 5.8, the
strong Lyo forest absorption depresses the flux in the N E816 band, leaving very little flux
from the quasar in this band. This effect shows up as nearly constant N E816 - z' colors.
At z ~ 6.1, Lyo flux comes into the z' filter and the forest absorption now starts to affect
the z' -band flux as well. This makes N E816 - z' slightly bluer at the highest redshift range
(z > 6.1).
In the extremely red regime (e.g., I- z' > 2.0), however, N E816 - z' colors of quasars
and T-dwarfs become degenerate. This is because both quasars and T-dwarfs have very
little flux in the N E816 and I band filters; the former due to strong IGM absorption and
the latter due to strong K I absorption. Thus near-IR follow-ups or direct spectroscopy are
required to disentangle these two populations.
However, since the majority of I-dropouts in the high-z quasar survey are Galactic M
and L dwarfs (i.e., T dwarfs and quasars are much rarer), our technique is useful to remove
such populations from the candidates. In our survey, we are interested in both high-z
quasars as well as T dwarfs - both of which are of great scientific interest on their own
right. The color criteria for z > 5.7 quasars and T dwarfs is shown as a region surrounded
by dashed lines in Figure 3.1, and it is expressed as:
I- z' > 0.7
and
NE816 - z' > 1.7
where magnitudes are in AB system.
In this Chapter, we discuss extremely red objects with I - z' > 2.0 colors found in
the HQT survey, with the aim of finding very high-z quasars (z > 6) and T dwarfs. This
selection is shown as the shaded area in Figure 3.1. The high-z (5.7 < z < 6.0) quasar
23
candidates with moderately red 1- z' colors (0.7 < 1- z' < 2.0) found by the HQT survey
will be discussed in Chapter 3.
m the HQT survey, instead of using the Sloan i' filter, we chose to use the Cousins I
filter since its effective wavelength matches better with the N B816. Sloan i' filter has a
shorter effective wavelength than the Cousins I, making i' - z' of a quasar at the same
redshift greater than I - z'. Nonetheless, in either I or i', one can use N B816 - z' color
criteria to eliminate M and L dwarfs from the candidates.
2.3 Observations
2.3.1 Subaru/Suprime-Cam Imaging Survey
The optical imaging data were obtained in a number of runs between 2002 September and
2006 September with Suprime-Cam (Miyazaki et al. 2002) on the 8.2m Subaru telescope.
Suprime-Cam consists of ten 2048 x 4096 ceo chips, yielding a field of view of 34' x 27' with
a pixel scale of fi!2. Its wide area coverage and good response at red optical wavelengths
make this instrument ideal for surveys for high-z objects and ultra-cool dwarfs, both of
which are very red at optical wavelengths.
The HQT survey area consists of three contiguous fields surrounding Abell 2218 (A2218-
wide; RAceD "" 16 : 30, DECcen "" +66 : 00), SSA22 (SSA22-wide; RAceD = 22 : 15,
DECceD = +01 : 00), and Abell 370 (A370-wide; RAceD = 02 : 41, DECcen = -01 : 44).
Two survey fields. A2218-wide and A370-wide include massive galaxy clusters. However the
area affected by these clusters is very small compared to the total survey area. Therefore
little effect from cluster properties such as a lensing magnification is expected. The fields
were chosen to be at low-extinction and overlapped with known, well-studied fields for
which we already have ultra-deep Suprime-Cam imaging data for the central region (~0.2
deg2) from narrowband high-z galaxy studies (Hu et al. 2002; Hu et al. 2004; Hu et al.
2005). The main advantage to choosing such fields is that it permits us to check the wide
24
3
o
-1
-z > 5.7 OSOS --Ultrocool Dwarfs
• Stars (0 - II) + L dwarf. lIE T dwarfs
o
L _ _ _____ _ _____ _ __ _ _ _ __ ____ _
2 3 4 NB816 - Z' [AB)
Figure 2.2 N 8 816 - z' vs. J - z' color·color diagram. Foreground objects are shown as asterisks (T dwarfs) , crosses (L dwarfs), alld circles (main sequence stars) . The solid blue li lle is the t rack of simulated quasar colors as a function of redshift . The sky hlue line is a track of synthesized colors of ultra-cool dwarfs (M5 - T8). Shaded area represents our color selection used in th is work to search for very high- z (z > 6.1) quasars and T dwarfs. Area surrounded by dotted lines ind icates z > 5.8 quasar and T-dwarf selection. Our candidate quasars and T dwarfs a re shown as red symbols where star symbols indicate confirmed late type L and T dwarfs and squares are other type of objects. Objects that were detected in only z'-band are shown as double squares. Annotated numbers correspond to the labeling in Table 2.3.
25
shallow photometry and also provides us with color calibration information in the event of
non-photometric conditions.
Observations with Suprime-Cam were taken as a series of three point dithers, with
typical exposure times on each Suprime-Cam field of 180 sec in I, and 540 sec in NB816
and z1. Data that were taken under poor weather conditions - that is, extinction greater
than 0.5 magnitudes and/or seeing sizes greater than I" - were re-taken under better
weather conditions. In order to achieve uniform survey depths, we only use the central
25' x 25' area from each Suprime-Cam pointing. Most of our I, NB816, and z1-band images
were obtained nearly simnltaneously to avoid transients such as supernovae and Kuiper
Belt objects. The data were reduced in a similar manner to that described in Capak et aI.
(2004) which is bSBed on Imeat software developed by Nick Kaiser 3. Calibrations were done
using spectrophotometric standard stars when available. Data that were obtained under
non-photometric weather conditions were cross-calibrated using stars in the overlapping
area for each Suprime-Cam field. For the SSA22-wide and A370-wide fields, existing ultra
deep Suprime-Cam data of the central field were used to cross-check the photometric zero
points. The I - N B816 and I - zI colors of field stars were also examined to ensure
the accuracy of photometric zero points. Photometric measurements were made using
SExtractor version 2.2.2 (Bertin & Amouts 1996) with "double image" mode to me8Bure I
and NB816 magnitudes for all objects detected in z'. Since each Suprime-Cam field were
observed under various seeing conditions (but always with better than l"seeing), we used
an aperture diameter of 2.5 x Seeing(FWHM) which is between 1'!6 - '1!4.
In order to determine survey area versus depths relations, we first formed a sensitivity
map (i.e. RMS map) by taking the inverse square root of the exposure map. The sensitivity
map was then scaled to have the flux unit using magnitude zero points and aperture diameter
used for each Suprime-Cam field. Next, we made the histogram of the scaled sensitivity
map to derive the relation of depths and number of pixels which have the same depths.
Survey area was then computed by integrating pixels. Fig 2.3 shows the resulting survey
3Imcat is available at http://www.lfa.hawaU.edu/-kaiser/lmcat
26
area versus depth relation in three filter bandpasses (I, NB816, and z'). This shows our
survey reaches limiting magnitudes of 25 [AB] in all three bandpasses, covering"" 14 deg2
to 50' limiting depths of lAB "" 23.8, N B816AB "" 23.4, and z'AB "" 23.5. However. since
z > 6.1 QSO and T dwarfs are estimated to be more than two magnitudes fainter in I and
NB816 than in z' (see Fig 3.1), we need deep I and NB816 images. Our candidate QSOs
and T dwarfs were selected to have magnitudes brighter than z'AB = 23.3. This means that
the faintest candidate from the survey is z'AB = 23.3, with I and NB816 magnitudes of
lAB ~ 25.3 and NB816AB ~ 25.0. Hence, in our candidate selection, we only used areas
where I and NB816 2u limiting magnitudes are lAB ~ 25.3 and NB816AB ~ 25.0. This
restriction limited our total survey area to 9.3 deg2, of which 1.9 deg2 is in A2218-wide, 3.0
deg2 in A370-wide, and 4.4 deg2 in SSA22-wide. The location of our survey field and area
coverage are summarized in Table 2.1.
2.3.2 QSO and T dwarf Candidate Selection
With the purpose of finding T dwarfs and quasars at very high-z (z > 6.1), we selected
extremely red objects with I - z' > 2.0 and Z~B < 23.3. Quasars at the lower redshift
range (5.7 < Z < 6.1) with moderately red colors (1.0 < 1- z' < 2.0) will be discussed in
Chapter 3. When applying 1- z' color criteria, we also included objects whose lower limit
on I - z satisfy the criteria of I - z' > 2.0. We examined all selected objects with eye to
remove transients and spurious detections due to spikes from bright stars. cosmic rays, and
artifacts.
We then applied N B816 - z' color criteria, that is N B816 - z' > 1.7, to eliminate
potential contaminants from M and L type stars. When applying N B816 - z' criteria, we
allowed objects whose upper limit of N B816 - z' satisfy the criteria. This was done because
photometric errors in N B816 are relatively large due to shallowness of the image in this
bandpass. In total, 10 objects (5 in A2218-wide, 3 in A370-wide, and 2 in SSA22-wide)
were chosen as final high-z quasar (z > 6) and T-dwarf candidates. Their coordinates and
photometric properties are summarized in Table 2.3.
27
20
15
~ N
'" " ~ 10 0
~ 5
0 22.0 22.5 23.0 23.5 24.0 24.5 25.0
Umiting Mag [5u, AS]
20
L816 15
~ N
'" " ~ 10 0
~ 5
0 22.0 22.5 23.0 23.5 24.0 24.5 25.0
Umiting Mag [Su, AS]
20
Z' 15
~ N
'" " ~ 10 0
~ 5
0 22.0 22.5 23.0 23.5 24.0 24.5 25.0
Umlting Mag [5u, AS]
Figure 2.3 Survey Area as a function of limiting magnitude in I, N B816, and z-band 28
NB816 z· J
b
••
Figure 2.4 Optical Suprime-Cam and near-IR MOIRCS images of HQT-9 are shown with a circle of 1.2/fradius. Each image is a 20/fsidc. This object was selected because of its extrernely red J - z' and N B 816 - z' colors, suggesting it is an eit her high-z quasru· or T dwarf. Follow-up near-IR imaging shows clear detection F\lrt her spectroscopy confirmed this object is a T6 dwar f.
Figure 3.2 presents opt ical ruld near-IR images of one of Our candidate, HQT-ll. as an
example. Our cand idates exhibi t extremely red 1 - z' and N B 816 - z' colors (i.e., luminous
in z' but very faint in I and N B 816), suggestive of high-z quasars or T dwarfs.
2.3 .3 Follow-up Nea r-IR Imaging
We carried out follow-up near-IR observations usiug various facili t ies on (vlauna Kea.
Table 2.4 summarizes our near- lfl. observations. Photometric measurements were made
again with an aperture diameter of 2.5 x Sccing(F WHM). In the following subsections, we
describe our follow-up observations.
UKIRT / WFCAM Imaging
T he J -band imaging data of two candidates (HQT-6 and HQT-7) were obtained with a near
IR Wide Field Camera (WFCAM: Casali et al. 2007) installed on the 3.8-m Uni ted Kingdom
Infrared Telescope (UKIRT). WFCAlVI is made of four 2048 x 2048 arrays, achieving a large
field of view of 0.2 deg2 with a relatively large pixel scale of 0.4". Total integration time
fo r each quasar/ T-dwarf candidate is 4 minutes. T he data were pipeline processed by the
29
Cambridge Astronomical Survey Unit (CASU), and field 2MASS stars were used to calibrate
the data.
Our data were obtained during the initial phase of WFCAM commis.<rioning when there
was an issue with image quality and cosmetics (see, Dye et al. 2006). As a result, the
stellar point spread function (PSF) was as large as 2/1 and therefore we used a 1arge aperture
diameter (~ 5") for the photometric measurements of candidates.
Subaru/CISCO & MOmCS Imaging
All our candidates except two (HQT-6 and HQT-7) were observed with the near-m
instruments on Subaru; CISCO (the Cool Infrared Spectrograph and Camera for OHS;
Motohara et al. 2(02) and MOmCS (Multi-Object Infrared Camera and Spectrograph;
Ichikawa et al. 2(06). CISCO is a 1024 x 1024 HAWAII array, covering 1.8'x 1.8'area
• with a pixel scale of 0.~11. MOmCS is a successor to CISCO and has capabilities for both
wide-field imaging (~4'x 7', 0'!117/pixel) and multi-object spectroscopy. MOIRCS employs
two HAWAII-2 2048 x 2048 arrays.
All the data except for HQT -8 were obtained under photometric conditions with better
than 0.5" seeing, and were calibrated by using UKIRT faint standard stars FS23 and FS139
in the Mauna Kea near-m photometric system (Leggett et al. 2(06). The HQT-lO data
were obtained under cirrus and thus were calibrated by 2MASS stars in the field. Data were
processed by the IDL based pipeline SIMPLE (SIMPLE Imaging and Mosaicking Pipeline;
W.-H Wang in preparation).
UKIRT /UIST Imaging
The H-band imaging data of HQT-9 and HQT-lO were taken with the UIST (UKIRT
Imager-Spectrometer; Ramsay Howat et al. 2004) on the 3.8-m UKIRT telescope under
photometric condition. UIST is made of a 1024 x 1024 InSb array with two plate scales
30
of O.l2"and 0.06", yielding a FOV of 2' x 2' or l' x 1'. Data were pipeline-processed by
ORAC-DR recipes provided by the UKIRT 4.
Subaru/CISCO spectroscopy of HQT-2 and HQT-3
The JH grism spectra of HQT-2 and HQT-3 were obtained with Subaru/CISCO on 2005
July 20 UT under conditions of cirrus with a typical seeing FWHM of 1". A set of twelve
and fourteen exposures of 300s were obtained for HQT-2 and HQT-3 respectively, both in an
ABBA dither pattern with a 10" offset. We used the 1" slit, yielding spectral coverage from
1.1-1.8ILm with resolving power R = >'1 a>. 0:; 200. The AOV star SAO 030187 was observed
immediately after the target exposures at a similar airmass. Data were reduced using
custom-made IDL scripts and IRAF I apall for the spectral extraction. First images were
pair-wise subtracted and divided by a median-combined fiat-field image. Images obtained
at the same nod position were then median-combined. Extractions of one-dimensionaI
specta were done with IRAF I apaIl. The transform between pixel and wavelength is fixed
for the CISCO and thus wavelength calibration was done by using the relation provided
at the CISCO website 5. Wavelength was independently calibrated using the night sky
lines, and this method gave the consistent results. Hydrogen absorption features intrinsic
to the atmosphere of a AO V star were removed by linear interpolation. In order to correct
for telluric absorption and instrument response, we divided the extracted spectra by the
spectrum of an AO V calibration star and mUltiplied by a 9,480K black body curve. Finally
the spectra obtained at different nod positions (A and B) were combined. Due to the
faintness of the targets (J ~ 20) and presence of cirrus, the SIN of the resulting spectra
was quite low. We then smoothed the spectra with a boxcar of 50 pixels. Figures 2.5 and
2.6 show final spectra for HQT-2 and HQT-3 with a spectral resolution R = 40.
'http://www.oracdr.org "http://www.naoj.org/Observlng/Instruments/OHS/spec/skyspec.html
31
Keck/NIRSPEC spectroscopy of HQT-9
The J-band spectrum of HQT-9 was obtained on 2004 December 5 UT with the
KeckII/NIRSPEC (Near-Infrared Echelle Spectrograph; McLean, I. S. et al. 1998) under
non-photometric, very windy (> 30 knots) weather conditions. NIRSPEC employs a
1024 x 1024 InSb array and has a high (Echelle) resolution mode (R <:::: 25,000), as well
as a low resolution mode (R <:::: 2,200). A total of 10 minutes of exposure was obtained in
low-resolution mode using the NIRSPEC-3 blocking filter which approximates J-band. A
0.76"slit width was used. The observation was taken at two different positions along the
slit, at a position angle of 111.480 east of north to avoid possible photon contamination from
a neighboring bright object '}!' away. An AO V calibration star (HD 215143) was observed
immediately after the exposures of HQT-9. The data were reduced using the REDSPEC
package developed by the NIRSPEC instrument group at the UCLA ill lab 6. The resulting
spectrum had poor SIN and was subsequently smoothed with a boxcar of 50 pixels, yielding
a spectral resolution of R <:::: 100 for final spectrum. Figure 2.7 shows the final spectrum of
HQT-9.
2.4 Analysis
2.4.1 Classification with Colors
The z - J color information is essential to separate candidate quasars from candidate T
dwarfs. Furthermore, in the case of L and T dwarfs, z - J provides a better constraint on
spectral types than I - Zl. This is because the photometric errors in optical bandpasses are
generally larger than those in near-ill bandpasses for these ultra-cool dwarfs which emit
majority of their light in the near-ill. We thus obtained follow-up J-band imaging data for
all our candidates. Figure 3.3 presents their I - z' VB. z' - J color-color diagram. The
symbols used follow those of Figure 3.1. The shaded area denotes the expected color range
for T dwarfs, and it is now well separated from the model high-z quasar track. H-band
6REDSPEC is available at http://www2.keck.hawaii.edu/inst/nirspec/DataRed.htmI
32
---.. -u CI) VI
o
H,o
20 !--,-CH.-,-_
/: ....
15
E 10 o c
5
o
1.2
:" '.
1.3
HQT-2
CH.
..... .... ;' .......•...
... \ ...... .
... \ ......................... ...... .
.......................... T2 T3 4 T5
1.4 1.5 1.6 1.7 wavelength (J.Lm)
Figure 2.5 Subaru/ CISCO J H spectrum of HQT-2 is shown as a black solid line. Also shown for comparison are the standard spectra for T dwarfs from Burgasser et aI. 2006. Major H20 and CH4 absorption bands, which are indicative of T dwarfs, are clearly detected . This object is classified as T2 - T4 based on the scheme of Burgasser et al. (2006).
33
15
o 10 E ..... o c
5
H,o
CH. ......
1.2
HQT-3
...... " ... - L9 TO T1 T2
1.3 1.4 1.5 1.6 1.7 wavelength (J.Lm)
Figure 2.6 Subaru/ CISCO J H spectrum of HQT-3 is shown as a black solid line. The comparison with the overplotted standard spectra of Burgasser et al. 2006 suggests this object is late-type L or early-ty pe T.
34
15
eli. HQT-9
"'" H,o ........ 101-~~-0 (l) Ul
o E '-o 5 c:
'-" -<
LL
o .... ,,'.
1.15
T7 T6 TS T4
1.20 1.25 1.30 1.35 wavelength (JLm)
Figure 2.7 Keck/ NIRSPEC J-band spectrum of HQT-9 is shown as a black solid line. This object appear to be T4 - T6.
35
imaging data exist for six candidates, and they are shown as red symbols in the r - J
VB. J - H diagram (Figure 2.9). Similarly, the r - J VB. J - K color plot is shown in
Figure 2.10.
Using these color-color diagrams, we classified objects into three categories: QSOs, 1-,
and T-type dwarfs. In case of the 1- and T-type dwarfs, we further estimated their spectral
types using a color-spectral type relation obtained from known dwarfs. This relation is
shown as the green dotted line in Figures 3.1 - 2.10 and was derived by fitting colors of
known dwarfs as a function of their spectral types with eighth order polynomials. Optical
colors of known dwarfs were computed from the spectra as our filter sets (1, N B816, r)
are unique, and the procedure is described in §2. Table 2.2 lists the coefficients of the
polynomial fits to colors. Near-ill colors come from photometric measurements of known L
and T dwarfs from Knapp et al (2004).
The candidate classification based on their color information is shown in the "Class"
section of Table 2.5. Six candidates out of a total of 10 are probable T dwarfs, though
one candidate (HQT-3) could be a late-type 1. These probable T dwarfs are shown with
star symbols in Figure 3.1 - 2.10, and are marked with asterisks in Tables 2.3 and 2.5.
The remAining three candidates are tentatively classified as high-z quasars, though optical
spectra are necessary to confirm if they are truly quasars. One quasar candidate, HQT-6,
could be a transient such as Kuiper Belt object or supernova. This is because its z'-band
images were not taken during the same observing run as its 1 and N B816-band images,
and HQT-6 was only detected in r. The Keck/DEIMOS spectra were obtained for two
candidate quasars (HQT-1 and HQT-7), and they will be discussed in a forthcoming paper
(Kaks:zu et al. in preparation). The present Chapter focuses on the ultracool dwarfs found
by the HQT survey.
2.4.2 Classification of Ultracool dwarfs with near-ffi spectra
In order to confirm that our T dwarf candidates are truly ultracool objects, we have
obtained near-ill spectra for three probable T dwarfs, HQT-2, HQT-3, and HQT-9. HQT-2
36
and HQT-3 were observed with Subaru/CISCO with the JH grism and their spectra are
presented in Figures 2.5 and 2.6. Both objects clearly exhibit ~ and H20 absorption,
which are the hallmark features of T dwarfs. Since the SIN of our spectra is low, we classify
the object by directly comparing its overall spectral shape with those of standard template
T dwarf spectra from Burgasser et al (2006), instead of using spectral diagnostic indices.
The closest fit to HQT-2 appears to be T3, though there is a ±1 subclass uncertainty in
the spectral type due to poor SIN in the spectra. Similarly, we assign a spectral type of L9
- TO to HQT-3.
Figure 2.7 shows the Keck/NIRSPEC J-band spectrum of HQT-9 with standard T
dwarf spectra from Burgasser et al. (2006) overlaid. The spectra were obtained under non
photometric, strong wind conditions, which resulted in poor SIN. However, the detection of
C~ and H20 absorption bands ensures that this is a T dwarf. The closest template spectra
is T6, though below 1.26JLm, there is an absorption feature that makes it deviates from T6.
We therefore assign a broad spectral type of T4 - T6 to HQT-9. All three T dwarfs that
are spectroscopically confirmed have spectral types consistent with those derived from their
optical/near-ffi colors. Table 2.5 lists spectral classifications for LIT candidates derived
from colors and spectra.
2.4.3 Distance Measurements
Generally, the distance estimates of field brown dwarfs with no parallax measurements are
made using absolute magnitude versus spectral type relations obtained from nearby brown
dwarfs. It has been observationally known that the dwarfs between late-type L and early
type T show an unusual strong brightening (- 1 mag) in thair J-band absolute magnitudes.
The similar brightening appears in other wavelengths (Z, H, and K), though the effect is
the greatest in J band. The identifications of LIT transition binaries by Liu et al. (2006)
and Burgasser et al. (2007) suggest that this brightening may be enhanced by a high
incidence of unresolved binaries at LIT transition and that the true brightening amplitude
is more modest. A growing number of resolved binaries supports this claim (Liu et al.
37
2007, in preparation). Here, to estimate distances of HQT candidates, we use the absolute
magnitude versus spectral type relations from Liu et al. (2006) and Liu et al. (2007).
Liu et al. (2006) presents two sets of polynomial fits - one dervied after removing known
binaries in Knapp et al. (2004) sample (L06-A), and another derived after removing known
and possible binaries (L06-B). The latter fits show almost no brightening "hump" at LIT
transitions, and thus the distances derived from this fits are smaller than the ones derived
from the former fits. The new relations of Liu et al. (2007, in preparation [L07]) are closer
to L06-B relations, however have slightly larger brightening amplitude.
The brightening effect is more modest in H-band than in J-band, and thus the
uncertainty caused by unknown binary fraction is smaller. Hence we used apparent H
band magnitudes to estimate distances for the HQT candidates with H-band photometry.
For objects without H-band data (HQT-5 and HQT-8), J-band photometry were used.
Table 2.5 shows the resulting distance measurements based on three relations described
above (L06-A, L06-B, and L07). The errors in distance measurements are calculated as
a combination of J-band photometric errors, errors due to spectral type uncertainties
in classifications, and MH (or MJ) scatter in the absolute magnitude VB. spectral type
relations. Three spectroscopically confirmed T dwarfs (HQT-2, HQT-3, HQT-9) - though
HQT-3 could be a late-type L - have distances of 60 - 170 pc. There are ouly a
handful of spectroscopically confirmed T dwarfs known today at distances beyond 60pc;
ULASJl452+0655 (T4.5) at 83 pc (Kendall et al. 2007), SOri 70 (T6) at 400 pc (Zapatero
Osorio et al. 2002) - though this was later disputed by Burgasser et al. (2004) to be at
75-lOOpc -, NTTDF1205-0744 (T6) at 90 pc, and two newly found T dwarfs from UKIDDS
at around 60 pc (ULAS J100759.90-010031.1 [T5.5] and ULAS J223955.76+003252.6
[T4.5j; Lodieu et al. 2007). Therefore, the three spectroscopically confirmed T dwarfs
found by our survey may be among the most distant, spectroscopically confirmed T dwarfs.
The discovery of distant T dwarfs from the HQT survey demonstrates the effectiveness of
our optical color selection.
38
2.4.4 Proper Motion
Since there is a greater than one year epoch difference between Suprime optical imaging
survey and J-band follow-up observations, it is possible to place limits on proper motions
and hence tangential velocities for our candidates. Space velocities of stars and substellar
objects are expected to increase as they age through interaction with the Galactic disk
stars. Therefore, proper motion measurements of field ultracool dwarfs can provide an
insight into their age and distribution in the Galaxy. For instance, at a given spectral type,
a lower velocity suggests a younger, lower-mass object, while higher velocity suggests an
older, higher-mass object. The detection of a high proper motion ultracool dwarf might
imply that it is an old member of Galactic disk population (Burgasser et al. 2007).
An extremely high proper motion might indicate membership in a low-metallicity, halo
population. Furthermore, proper motion measurements can be used as a way of separating
out ultra-cool dwarfs from high-z galaxies/quasars that have the similar optical colors as
dwarfs, since extragalactic objects have zero proper motion.
In order to measure proper motion, we need to register a follow-up J-band image of each
candidate onto its Suprime-Cam z' -band image. The registration can be done using the field
stars which appear in both J and z' -band images. We chose such astrometric reference stars
in the following way. We first selected objects that lie within 2' of each candidate. This area
constraint was set to minimize the effect of non-linear distortion. We then choose objects
that are detected in both optical z' and J-band at > 100', but which are not saturated.
We finally selected objects whose SExtractor star parameters are greater than 0.8 and with
flags parameters of zero. These criteria yielded typically 15 - 20 astrometric reference stars
for each candidate.
Using these reference stars, we derived the third order transformation coefIicients
between J-band image and z'-band image. The astrometric registration tool in the SIMPLE
software (W.-H Wang 2007, in preparation) was used in this process, which finds a solution
in an iterative manner by rejecting stars with registration errors greater than three sigma.
Proper motion was then measured by directly comparing absolute coordinates of each
39
C8Jldidate in Suprime z' and in the registered J image. The errors in proper motion were
computed by taking a median of position differences of reference stars between z' and J
images. We also measured proper motion independently by measuring the change in the
median of distances between a C8Jldidate and its astrometric reference stars. This alternative
measurement was carried out in order to examine a possible systematic offset between J and
z' -band images arising from the registration. However this method yielded the consistent
results with the previous method within one sigma errors. In the following discussion, we
adopt the values derived by the direct comparison of C8Jldidate absolute coordinates in z'
and J band images.
Table 2.6 summarizes the results of our proper motion measurements in units of milli
arcsecond [mas] per year. Note that none of the candidate quasars were examined for
motions since their J-band magnitudes are too faint (::;; 5 sigma) to unambiguously measure
their positions in the follow-up J band images. The proper motion of HQT-3 was not
measured since its J band image lacks a sufficient number of reference stars to register onto
the z'-band image. Land T dwarfs found in our survey are among the most distant known
ultracool dwarfs, and hence a long time baseline tlT is required to accurately measure their
proper motion. As such, the short baseline tlT of our observations have resulted in large
errors in proper motions, and consequently in tangential velocities, Vtan = 4.74Jtd km s-l.
Future follow-up observations are necessary to minimize the errors and to draw reliable
conclusions. However, the spectroscopically confirmed T dwarf with the longest baseline
observations (HQT-9: T4-T6: tlT = 3.6 year) in our sample, has a reasonably good proper
motion measurement (/1 = 170 ± 19 mas year-l). Furthermore, this object was observed
at three different epochs (Suprime-Cam; UT 2003.10.26 & UT 2005.09.05 and MOmCS;
UT 2007.06.12), allowing three independent proper motion measurements, which all agree
with each other within the errors. Here, we adopt the proper motion derived from the first
epoch Suprime-Cam z'-band image and follow-up MOmCS J-band image, since this is the
longest time baseline and yielded smallest error.
40
The tangential velocity of HQT -9 is Vtan = 65 ± 14 kin S-1 , which is higher than those of
nearby T4-T6 dwarfs reported in Vrba et al. (2004). The highest tangential velocity in the
Vrba et al. sample of T4 - T6 dwarfs is Vtan = 51.3± 2.6 kin s-1 (2MASS J235654-1553;
T6) and the mean velocity for T4 - T6 dwarfs is Vtan = 29.25 ± 1.3 kin S-I. Among all
18 T dwarfs of Vrba et al. (2004), there are only three T dwarfs whose Vtan exceeds 60 kin
s-l; 2MASS J041519-0935 (T8/T9) with Vtan = 61.4 ± 1.0, SDSS J015141+1244 (T1±1)
with Vtan = 75.3 ± 5.5 kin 6-1, and 2MASS J104753+2124 (T6.5) with Vtan = 86.5 ± 3.5
kin s-l. The large tangential velocity of HQT-9 may suggest that it is an older and more
massive T dwarf than those in Vrba et al. (2004). The remaining candidate L and T dwarfs
appear to have the velocities consistent with Vrba et al. (2004) (Vtan = 20 - 60 kin s-1
for T dwarfs), though given the large errors, future follow-up observations are necessary.
The large velocity of the newly found T dwarf, HQT-9, and the lack of T dwarfs with small
velocities may imply that T dwarfs are old members of the Galactic disk.
2.4.5 Number Counts and Galactic Scale Height
Few studies have been done so far on the vertical scale height of L and T dwarfs, as this
requires probing distant dwarfs over wide areas at different Galactic latitudes (e.g., Liu et
al. 2002). Ryan et al. (2007) estinIated the scale height of L/T dwarfs to be 350 ± 50 pc
based on 28 faint dwarf candidates (21 < zAB < 25) found in the deep HST/ACS parallel
fields. Their estimate is consistent with the known trend of increasing scale height with
decreasing stellar mass. However, given the small area coverage (~ 0.04 deg2), ideally one
would like to have dwarfs from a wider area.
We have identified 6 T dwarfs from the HQT survey, of which three were
spectroscopically confirmed including one possibly late L dwarf, HQT-3. Although the
number of detected ultracool dwarfs by the HQT survey is still small, our survey area
coverage of ~ 10 deg2 is significantly larger than the previous work on the vertical scale
height [~ 0.5 deg2 of !fA deep survey in Liu et al. (2002); ~ 10 arcmin2 of HUDF in Pir2kal
et al. (2005); ~ 0.04 deg2 in Ryan et al. (2007)]. Thus our data may provide a better
41
constraint on the scale height. Using a simple analytic model of Galactic disk, we estimate
the expected number of dwarfs and examine if the HQT survey can set a useful limit on
the Galactic scale height of these ultracool dwarfs.
Following the steps described in Liu et al. (2002) , we calculate the expected number
of L and T dwarfs from the HQT survey as a function of Galactic scale height. We first
assume that the distribution of T dwarfs follows the exponential Galactic disk model of
Wainscoat et al. (1992) with a canonical radial scale length of 3.5 kpc. We can estimate
the maximum distances of dwarfs that can be probed by the HQT survey from the survey
magnitude criterion of z'AB = 23.3 and the absolute magnitude VB. spectral type relations.
The effect of brightening at LIT transitions are more modest in K band than that in J
band. AB such, the uncertainty in the brightening amplitude due to unrecognized binaries
is expected to be smaller in K band. Hence we use MK VB. Spectral type relations for our
analysis and transform MK to Mz' using the z' - K synthesized colors of dwarfs presented
in Table 2.2.
With L07 relations, we can probe L dwarfs out to 700 pc and T dwarfs out to 190 pc. If
we use L06-A relations, the probable distances become slightly larger (T dwarfs out to 200
pc) due to a larger brightening amplitude in the early T dwarfs. Consequently the survey
volume become larger for the L06-A relations. AB a result, the number of dwarfs found by
the survey is expected to be larger with L06-A relations. We carry out our number counts
analysis based on three relations (L06-A, L06-B, and L07).
We adopt a T dwarf space density of 0.006 pc-3 and assume an equal uumber of objects
per spectral subclass from TO to T8. This space density was derived observationally from
the local2MASS T dwarfs (Burgasser 2001) and agrees well with the IMF model of Chabrier
et al (2005). For L dwarfs, we adopt 0.01 pc-3 which is consistent with the observational
data (Reid et al. 1999) and theoretical predictions (Chabrier 2002). Similarly as T dwarfs,
we assign an equal number of objects per spectral subclass from LO to L8.
Given these quantities, we shall now calculate the expected number of T dwarfs for each
spectral subclass by multiplying the exponential disk density of Wainscoat et al. (1992) by
42
survey surface area, and integrating it over the distance to which our survey is sensitive.
Finally we summed the expected number of each subclass of T dwarf to derive the number
of the entire T dwarf population. Figure 2.11 presents expected number of L dwarfs (top)
and T dwarfs (bottom) in entire HQT survey area as a function of Galactic scale height of
these dwarfs. Model predictions are presented as solid lines [L06-A (black), L07-B (redj, L07
(blue)]. One sigma errors derived from the uncertainty in the absolute magnitude relations
are also shown as dotted lines. The effect of the use of different relations is almost negligible
in the number of L dwarfs, but is clearly evident in that of T dwarfs. Our results suggests
that the number of L dwarfs discovered by the HQT survey is sensitive to their Galactic
scale height. However, this is not the case with T dwarfs where numbers become essentially
constant at the scale height greater than 200 pc. This is because the HQT survey is only
sensitive to T dwarfs within 200 pc.
Our findings of 6 T dwarfs in the HQT survey are indicated as a horizontal dotted line
with a shaded area for a 10" Poisson statistical uncertainty. Although the large uncertainty
is seen in the expected number of T dwarfs, our findings are roughly consistent with the
model predictions. Furthermore, L06-B and L07 relations appear to match better with our
observational results than the L06-A relations - that were derived based on Knapp et aI.
(2004) data. This may be another evidence for the large fraction of unresolved binaries at
the LIT transitions.
It is known that the number ratio of LIT dwarfs at different latitudes is a function of
their Galactic scale height (Liu et aI. 2002; Ryan et aI. 2005). Thus by comparing the
number of identified dwarfs in the fields at different latitudes. one may be able to set a useful
limit on the scale height. We therefore examine the expected number of dwarfs at different
Galactic latitude. The purpose here is to understand how one's choice of survey field affects
the survey sensitivity to the dwarfs' scale height. We follow the same procedure described
above, except we now compute the expected number of L and T dwarfs as a function of
both scale height and field latitude. L07 relations are used to estimate the survey volume.
The top panel of Figure 2.12 presents the results for our survey depths. In the bottom
43
panel, we show the results for a one magnitude fainter survey to illustrate the effect of the
survey depths. In both CIlBeB, the dwarfs numbers are normalized at the 0° latitude and
they are shown as black (L dwarfs) and red (T dwarfs) solid lines for the scale height of 100
pc. Dotted lines indicates the results for the scale height of 400 pc.
Three HQT survey fields are indicated with blue dotted lines. It can be seen from the
top plot in Figure 2.12 that the T dwarfs number ratio varies very little between the low
and high-latitude field of the HQT survey for both scale heights. Thus the HQT survey is
not very sensitive to the scale height of T dwarfs. On the other hand, L dwarfs number
ratio is a strong function of the scale height; the number ratio is ahnost 1 for the scale
height of 400 pc, but it becomes 1.6 with the scale height of 100 pc.
This figure clearly demonstrates that by observing multiple fields at very different
latitudes, one can set better constraints on the dwarf scale height. For example, number
ratios of dwarfs in Ibl = 10° field and Ibl = SOo field is ~ 13 for the scale height of 100 pc.
This becomes 2 with the scale height of 400 pc. In case for T dwarfs, the number ratio
changes from 2.5 to 1. The bottom plot in Figure 2.12 shows the results for a survey one
magnitude deeper than ours. The number ratio becomes more steep and therefore one will
be able to set better constraints on the scale height. The T dwarf number ratio at low
(Ibl = 10°) and high-latitude (Ibl = SOO) field is 4 if the scale height is 100 pc. This becomes
1 if the dwarf scale height is 400 pc.
2.5 Conclusion
We have shown that the addition of the NB816 filter to the traditional optical (I, z') survey
for high-z quasars is beneficial for breaking a well-known color degeneracy between quasars
and foreground M and L dwarfs. Our color selection is useful as it does not require follow
up J band imaging which is still time consuming compared to wide-field optical mosaic
CCD imaging. The present Chapter discusses objects that were selected by their extremely
red I - z' and N B816 - z' colors with the purpose of identifying T dwarfs and quasars
44
at the highest redshift range probed by optical z'-band surveys (z > 6.1). Among 10
such candidates, we have successfully identified 5 or 6 T dwarfs of which three have been
spectroscopically confirmed. These T dwarfs are among the most distant spectroscopically
confirmed ultracool dwarfs. The proper motion measurements of these T dwarfs indicate
that they are old members of Galactic disks, though future follow-up imaging is necessary
given large errors in our measurements.
We have examined the expected number of L and T dwarfs from the HQT survey using a
simple analytic model in which dwarfs follow an exponential disk distribution. Our discovery
rate is consistent with the model predictions, though the uncertainty is large due to small
number statistics. However, it does seem that the previous relations between the absolute
magnitudes of dwarfs and their spectral types (e.g., Knapp et al. 2004) over-predict the
dwarf numbers. And that our results appear to match better with the recent relations of
Liu et al. (2007, in preparation). This may support the claim that the brightening at LIT
transitions is enhanced by the high binary fractions of these types. Future high resolution
imaging observations of more LIT transitions are needed to solve the uncertainty due to
the binary fraction. We have also investigated the effect of survey field latitude on the
determination of dwarf scale height. Our results show that multiple field survey that cover
wide range of latitudes will be very useful in the determination of the scale height of dwarfs.
45
5
-% > 5.7 OSOs --Ultrocool Oworfs • 4 • Stars (0 - 101) IJI + L dwarfs lIE T dwar fs
~
0 3 '" ., > ~ • • • • .., •
2 • ? m:::& ' ~ • , ! ,
~.40 • '-15 i
F • f;;J 'N
~ 7:&5 80 .... 590
6.30
.519 , , , , ,
0 , , ,
0 2 3 4 I - z' [AB)
Figure 2.8 I - z' vs. z' - J color-color diagram for candidates. Symbols are the same as in F igure 3.1, but in this color representations the shaded region only contains T dwarfs. Although a narrowband N B 816 fi lter helps to separate MI L dwarfs from high-z quasar candidates, there is sti ll a color degeneracy between quasars and T dwarfs ( ee, Figure 3.1). This plot demonstrates that the z' - J color can cleanly break this degeneracy.
46
1.5
-z > 5.7 QSOs + -- Ultrocool Dworfs + +
1.0 • Star. (0 - lot) + L dwarf. * T dwarf.
~
0.5 0
'" ., > ~
J:
.., 0.0
• •
-0.5
-1.0~~~~uw~~~~~~~~~~~~~~~~~
o 2 .3 4 5 Z' (AB] - J (Vega]
Figure 2.9 z' - J vs. J - H diagram shown with the same symbols as in Figure 3.1. Five HQT candidates with H band imaging data are plotted with numbered stars and error bars, and they lie in the shaded region of ul tracool dwarfs. Note that errors are much smaller than in the optical. T his is because ultracool dwarfs are much brighter than in the near-IR wavelenths. T he colors of z > 5.7 quasars are degenerated with M stars. T dwarfs are characterized by their red z' - J color. However in J - H color, they become increasingly blue at later spectral type.
47
2.5
-z > 5.7 OSOS + + + 2.0 - Ultrocool Dworfs
#+* • Star. (0 - Iol) + L dwarf. +. lIE T dwarf. ipt.
1.5 + 1IIit'" ...
~ • 0
'" 1.0 ., .~ '-45 > -
". '" I 0.5 ....,
0.0 ·l~ .:~ .. 311
-0.5 • TiU
- 1.0 0 2 3 4 5
z· [AB] - J [Vega ]
Figure 2.10 z' - J vs. J - f( diagram shown with the same symbols as in Figure 3.l. Only one candidate in Table 2.3 (HQT 8) was observed with K -band, and is plotted in this diagram. A similar color trend for J - K colors with z' - J is seen for ultra-cool dwarfs that were plotted in terms J - H colors in Figure 2.9, although typical J - K color values are about half a magnitude larger at the same z' - J color. T his is due to the increasing gravity/ metal licity in the later T dwarfs, which suppresses the f( band ftux. The z 1. 5.7 quasar track is
48
400 , " ' " ' " ' " ' " ' " ' " ' " ' . '
!? . ' 300 . '
5 . ' " ' • " 0 " " ...J " " ,:::::::::::::::. e; " " " ~ " 200 " c " E " #; .*
~ • .. .... .... z • , .... ..
• ....... ' "0 • .. ....... ~ • --u • ........
• ........... ~ • Q.
100 io .... ..
x "
__ 106-A ... • " __ L06-8
• • . . • __ L07 . . a a 100 200 300 400 500
Scale Height [pc]
12 , , , , -, - - ' ----
10 ---- - -- -
!? ,--5 • 0 8 .... e; ~ .. c E ~
Z "0 ---- ------.. U 4 .. Q. x __ L06-A ...
2 __ L06-8
__ L07
a a 100 200 300 400 500
Seal. H.;ghl [pc]
Figure 2.11 Expected number of L dwarfs (top) and T dwarfs (bottom) as a function of Galactic scale height). Three kinds of MI< versus spectral type relations from Liu et al. (2006) [L06-A and L06-B] and Liu et al. (2007, in preparation) [L07] are used for distance estimates. It can be seen that the differences in these relations result in a large difference in expected number of T dwarfs, but not in the L dwarfs. Our fi ndings of six T dwarfs with 1 a Poisson error range are shown by the shaded region in the right hand figu reo
49
0
" ~ "0
.~ 0 --'
0 "0
0. 10 ~ .~ 0 E (; c ~
~ D E ~
z
0
" ~ "0
.~ 0 --'
0 "0 ~ 0.10
.~ 0 E (; c
~ D E ~ z
................... ......... ............ ........ ....... ·· i····, ! · · ············ ~ .......... " ........................................ .
_ _ l dwarfs (H ". l00pc]
[H - .OOpe)
__ T dworl. [H - 1 OOpe]
[H - '00pc]
z"m = 23.3 [AB)
20 40 60 80 Galact ic l atitude (deg]
........... I ' I ........... " y ' " '" •• ••••••••• J. ' •••• • • I I ...•........... , .................• . " .....
__ l dworfs (H - 100pc)
[H - 'OOpe] __ T dwarfs [H _ 100pc)
[H - .OOpe]
20 40 60 Galactic Latitude [deg]
.... ........... .. .. .. ........
Z'.;m 24.3 [AB)
80
Figure 2.12 Expected number of L dwarfs and T dwarfs shown as a function of field Galactic latitude in degrees (Ibl) . Top panel shows the model predict ions fo r the HQT survey depths. Bottom panel shows the results for a survey that goes one magnitude deeper than the HQT survey. Numbers are normalized at a latitude Ibl = 0° . The locations of the three HQT survey fields are indicated with dashed blue lines.
50
Table 2.1. Survey Area and Ultra-cool Dwarf Number Counts
Field Name RA range DEC range Galactic Latitude Survey Area Used Area [deg] [deg2] [deg2]
A2218-wide 16:14:10.40 - 16:42:28.00 +65:10:07.30 - +66:50:09.95 39.0 3.50 1.87 A370-wide 02:36:11.13 - 02:45:19.71 -03:03:04.82 - -00:29:04.89 -53.0 5.25 3.00 SSA22-wide 22:10:28.21 - 22:21:28.39 -00:01:02.96 - +02:05:57.00 -43.0 5.25 4.42 Total 14.0 9.30
Note. - The location of the three HQT survey field is shown. "Survey Area" means the area observed with g: Subaru/SuprimeCam in three bandpasses (1, NB816, z'). "Used Area" shows the area used in this paper in which
the Sa limiting magnitudes are lAB> 24.3, N B816AB > 24.0, and z'AB > 23.3.
Table 2.2. Coefficients of Polynomical Fit to Colors of Ultra-Cool Dwarfs
Color cO cl c2 c3 c4 c5 c6 c7
I-z' 1.187e+OO 3.235e-02 -5.3000-03 2.044e-03 -4.500e-06 -1.555e-05 7.703e-07 -9.03Oe-09 i'-z' 1.52ge+OO 3.463e-02 -2.713e-03 2.77ge-03 -6.910e-05 -1.951e-05 1.292e-06 -2.226e-08 NB816-z' 6.371e-Ol 7.816e-02 -2.11ge-03 4.426e-05 1.643e-04 -2.13ge-06 -1.132e-06 4.324e-08 z'-J 2.613e+00 1.705e-01 -1.728e-02 -1.546e-03 2.201e-04 5.674e-06 -1.040e-06 2.398e-08 J-H 6.832e-Ol 4.917e-02 2.10ge-03 -2.465e-04 -2.29ge-05 -2.046e-06 9.277e-08 2.987e-09 J-K 1.197e+00 8.577e-02 7.732e-03 -1.726e-06 -1.98Oe-04 -4.205e-06 1.26Oe-06 -3.253e-08
Note. - These are the coeffiecients of the eighth-order polynomial fits for the optical and near-ffi colors of M - T dwarfs, ~ and are used in Figure 3.1 and Figure 3.3 - Figure 2.10. Note that Cousins I filter is used in the HQT survey, however we also
list Sloan i' - z' colors as a reference for interested readers. The fits are defined as
Col(Yf" = L C; X (SpT)i (2.1) i
where SpT = -10 for MO. SpT = 0 for LO, SpT = 10 for TO. and SpT = 19 for T9.
Table 2.3. Photometric Properties of Candidates
Name RA DEC z' J H 1- z' NB816- z' z' -.J J-H J-K
HQT-l 16:23:31.78 +65:34:36.39 23.3 (0.13) 22.0 (0.11) > 1.75 > 1.55 1.28 (0.17) HQT-2* 16:28:25.99 +65:25:49.07 23.3 (0.07) 20.1 (0.09) 19.6 (0.07) > 2.40 > 2.20 3.15 (0.11) 0.57 (0.11) HQT-3* 16:34:13.34 +65:18:12.65 22.7 (0.04) 19.7 (0.07) 18.8 (0.04) 2.31 (0.30) > 2.79 2.96 (0.08) 0.92 (0.08) HQT-4* 16:38:36.72 +65:47:14.97 22.7 (0.04) 19.5 (0.06) 19.0 (0.04) 2.84 (0.52) 2.20 (0.35) 3.13 (0.07) 0.49 (0.07) HQT-5 16:18:52.44 +66:37:03.83 22.8 (0.07) 20.2 (0.02) 2.32 (0.48) > 2.12 2.66 (0.08) HQT-6 02:45:01.67 -02:05:36.68 23.3 (0.07) > 21.3 > 2.50 > 2.10 2.00 (90.0) HQT-7 02:44:52.77 -02:06:49.42 23.4 (0.08) > 21.3 2.28 (0.47) > 2.00 2.10 (90.0) HQT-8* 02:37:00.57 +00:39:07.72 22.7 (0.04) 19.6 (0.07) 2.06 (0.24) 1.77 (0.24) 3.07 (0.08) 1.00 (0.09) HQT-9* 22: 19:03.09 +00:24:18.75 23.0 (0.08) 18.9 (0.01) 19.5 (0.10) 2.54 (0.14) 3.21 (0.41) 4.28 (0.03) -0.58 (0.10) HQT-IO* 22:21:26.92 +01:03:51.12 23.3 (0.12) 20.0 (0.15) 19.9 (0.16) 2.24 (0.51) 1.28 (0.50) 3.30 (0.19) 0.07 (0.22)
~ Note. - Optical magnltud ... (I, z', NB816) are given in AB system, while near-infrared magnltud ... are in Vega,-based system. Objects with asterisks
are probable T-type dwarfs from their near-IR colors, though HQT-3 could be a late L dwarf. Two sigma upper magnitude and color limits are shown for the objects that were not detected in I, NB816, or J-band.
Table 2.4. Log of Nea.r-ffi Imaging Observation
Name Filter UT Date Instrument Exp. Time [sec] Weather [seeing]
HQT-l J 2007.04.29 SubarujMOIRCS 720 Photometric [0.4"] HQT-2 J 2005.07.21 SubarujCISCO 660 Photometric [0.7"J
H 2005.07.21 SubarujCISCO 660 Photometric [0.7"] HQT-3 J 2005.07.21 SubarujCISCO 660 Photometric [0.7"]
H 2005.07.21 SubarujCISCO 660 Photometric [0.7"] HQT-4 J 2005.07.21 SubarujCISCO 660 Photometric [0.7"]
H 2007.06.12 SubarujMOIRCS 135 Photometric [0.5"] HQT-5 J 2007.04.29 SubarujMOIRCS 720 Photometric [0.4"] HQT-6 J 2005.08.30 UKIRI' jWFCAM 240 Thin cirrus [2"] HQT-7 J 2005.08.30 UKIRI' jWFCAM 240 Thin cirrus [2"] HQT-8 J 2007.01.06 SubarujMOIRCS 1080 Cirrus [1"]
Ks 2007.01.06 SubarujMOIRCS 150 Cirrus [1"] HQT-9 J 2007.06.12 SubarujMOIRCS 720 Photometric [0.5"J
H 2007.07.08 UKIRI' JUIST 1440 Photometric [O.7"J HQT-lO J 2005.05.08 UKIRI' fWFCAM 240 Photometric [2"]
H 2007.07.08 UKIRI' JUIST 1440 Photometric [O.7"J
Table 2.5. Classification and Distance Measurements
Name Class Sp Type Adopted Type Distance [PC] L()6.A L()6.B L07
HQT-l QSO HQT-2* TO - T3 T2 - T4 T2 -T4 157±48.6 115±30.3 134 ± 37.S HQT-3* L7-TO L9 - TO L9 - TO 119 ± 29.5 00.6 ± 20.4 109 ± 24.S HQT-4* TI- T3 TI - T3 133 ± 40.6 00.5 ± 25.0 114 ± 31.6 HQT-5 Lete L L7 - L9 138 ± 47.2 127 ± 4O.S 138 ± 43.0 HQT-6 TransIent or QSO HQT-7 QSO HQT-S* TO-T2 TO-T2 131.0 ± 45.3 101 ± 32.9 115 ± 36.6 HQT-9* T4 - TO T4 -T6 T4 -T6 104 ± 33.4 86.3 ± 23.6 93.5 ± 27.3 HQT-lO* T3- T5 T3 -T5 160 ± 53.3 123 ± 35.6 139 ± 42.7
Note. - The "Class" was determined from object's optical and nesr-IR colors, whereas the "Sp type" was decided by the nesr-IR spectra. We gave precedence to "Sp type" for the candidetes with nesr-IR spectra, and final spectral type for candidate brown dwarfs are Iieted In "Adopted Type". Distances are measured using relationships between absolute magnitude and spectral types from Llu et al. (2006) [LQ6.A was derived after excluding known blnarles In Knapp et al (2004) sample: LQ6.B was derived after excluding known and possible binariesJ and Liu et al. (2007, In preparation).
54
Table 2.6. Proper Motion Measurements
Name .o.T I' v ... L06-A UJ6.B L07
[yr] [1D8B/yr] [\<m/o] [\<m/a] [\<m/o]
HQT-l 2.9 HQT-2 1.1 130±97 00.8 ± 78.2 70.9 ± 56.1 82.6 ± 65.9 HQT-3 HQT-4 3.0 47±25 29.7 ± 18.2 21.5 ± 12.7 25.4 ± 15.2 HQT-5 1.9 77±33 00.5 ± 27.7 46.2 ± 24.S 00.2 ± 26.7 HQT-6 HQT-7 HQT-S 2.2 44±21 27.S ± 15.8 20.5 ± 11.1 24.0 ± 13.S HQT-9 3.6 170±19 83.6 ± 28.5 69.5 ± 20.5 75.3 ± 23.6 HQT-IO 1.7 66±28 00.0 ± 27.0 38.5 ± 19.8 43.5 ± 22.S
Note. - Proper motions and tangential velocities of HQT candidates are listed. .0. T means observation time baseline. Tangential velocities were estimated using the distances presented in Table 2.5.
55
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59
Chapter 3
Hawaii Quasars and T-dwarf {HQT} Survey II.
Constraints on the z = 6 Quasar Luminosity
Function
Abstract
In Chapter 2, we examined extremely red (I - z' > 2.0) objects found by the HQT survey.
Such objects could be either ultra-cool dwarfs or quasars at the hlghest redshift range
(z > 6.1) probed by the reddest optical filter z'. Using follow-up near-ffi data, we have
shown that more than half of them are ultracool dwarfs of late L and T. In this Chapter, we
discuss moderately red objects (0.8 < 1-z' < 2.0) selected in the 9.0 dlllf area at z'AB < 23.3
with the purpose of finding quasars at lower redshift range (5.7 < z < 6.1). Follow-up
spectroscopy with KeckjDEIMOS. however. has shown that none of the candidates is high
z quasars. Instead, majority of them are emission line galaxies at intermediate redshifts
(z < 1), whose extremely strong [0111]>.5007 and little continuum make them I-band
dropouts. Based on the non-detections of z > 5.7 quasars in the HQT survey, we determine
the upper limit on the faint-end slope of quasar luminosity function to be a < -2.1 and
< -1.1 at the 99% and 90% confidence level, respectively. This is significantly flatter than
the bright-end slope f3 = -3.1±0.4, recently reported by Jiang et al. (2007) from the SDSS
southern deep survey. We also find the characteristic magnitude of -25 < Mi450 < -24.3 at
60
99% confidence level. Our constraints are more stringent than the previous studies because
of the larger area coverage of the HQT survey at this magnitude limit, and are consistent
with the present picture of the cosmic reionization in which quasars are considered to be
the negligible contributor to the reionization.
3.1 Introduction
Identifying the objects at high-redshift (z > > 5) is a direct way of probing the sources
of UV background radiation in the early universe. and advancing our understanding of
cosmic reionization and structure formation. The last few years have witnessed tremendous
progress in the identifications of gaIaxies at z > > 5. We now know more than one hundred
spectroscopically confirmed gaIaxies at z > 5.7. Majority of them were identified through
their strong Lya emission failing into the narrowband filters (e.g., Hu et aI. 2002, 2004;
Ouchi et aI. 2005; Shimasalru et aI. 2006; Rhoads et aI. 2003; Taniguchi et aI. 2005). The
addition of the reddest optical bandpass z' has resulted in the successful discoveries of ~ 20
very luminous quasars (Ml450 < -27) in the SDSS survey from an area of ~ 7,550deg2 in
the northern Galactic cap (Fan et aI. 2000, 2003, 2004, 2006; Goto et aI. 2006). However,
studies of lower-luminosity quasars, which form the bulk of quasar population at high-z,
have been less well advanced. The major reason for this is the extreme rarity of high-z
quasars and the overwhelming number of possible contanIinants, mostly very cool dwarfs
of type M and L (e.g., Fan et aI. 2001). Separating these two populations requires J band
imaging, which is time intensive to carry out over the wide area to the depths required by
these surveys.
The recent development in wide-field near-ffi cameras, combined with the deeper optical
imaging data, have finally began to shed light on fainter, hence more "nornJai", quasars at
high-z. Just in this year, about a dozen z > 5.7 quasars which are ~ 1-2 magnitudes fainter
than the bright SDSS quasars, have been discovered. SDSS deep survey, which covers an
area of 260 deg2 to z' AB 11m = 21, have identified five quasars with Ml450 ~ -26 (Jiang et aI. ,
61
2(07). One of their quasars at z = 5.86 was independently discovered by Venemans et al.
(2007) in the UKIDSS survey using Y filter. Cool et al. (2007) found a quasar at z = 5.85
from an area of 7.7 deg2 in the NOAO Deep Wide-Field Survey Bootes Field (NDWFS;
Jannuzi & Dey 1999). Their quasar was selected from the Spitzer mid-infrared data, and
is about half a magnitude fainter than the bright SDSS quasars. The initial deep search
of z > 5.7 quasars in the Canada-France High-redshift Quasar Survey (CFHQS deep) did
not find any quasars in the 3.83 deg2 field at the magnitude limit of z'AB = 23.35 (Willott
et al. 2005). However, in more recent CFHQS survey that is shallower (z'AB = 22.5) but
covers significantly wider area than their initial search (~ 4OOdeg2), Willott et al. (2007)
discovered four quasars at z > 6. Their followup observations are still underway and
therefore no statistical analysis has been made.
Ongoing surveys mentioned above (e.g., SDSS deep; CFHQS; UKIDSS) are expected to
find more quasars at fainter limits than the SDSS main survey. However, quasars found by
these surveys are still more than two magnitudes brighter than the characteristic magnitude
(Mi450) at z = 2 (Croom et al. 20(4). The SDSS survey has shown that the quasar
number density declines from z = 2 to z = 6, suggesting a fainter M* at z = 6 (Fan et
al. 2004). Therefore current surveys are unlikely to probe the faint-end of the quasar
luminosity function at which the bulk of the quasar population reside. It is therefore
necessary to conduct a new survey for high-z quasars which goes down to fainter limits, yet
covers sufficiently large area in order to better sample quasars and to overcome the cosmic
variance that are already known to be present at z» 5 from the galaxy studies (e.g., Hu
et al. 2(05).
Finding quasars when the universe was merely one billion years old advances our
understandings of the formation of the supermassive black holes and host galaxies.
Furthermore, distant quasars serves as lighthouse of the early universe by which we can
learn about the ionization status of the IGM. We already know that quasars are unlikely
the major source of reionization from their extreme rarity shown by deep X-ray surveys
(Barger et al. 2003) and the optical surveys mentioned above. However, the exact quasar
62
contribution to the ionizing background in the early universe is still poorly constrained due
to the lack of the knowledge about the faint-end of the quasar luminosity function at high-z.
In order to move beyond the current uncertainties, we have conducted a survey for faint
quasars (Ml450 < -23.5) at z > 5.7, which we call Hawaii Quasar and T dwarf survey (HQT
survey). In Chapter 2, we described the survey data and introduced a new color selection
technique using a narrowband filter N B816. Our technique is ideal for searching for faint
high-z quasars, as it does not require follow-up near-ffi imaging which is time consuming
to carry out over the large area to the required depths. We have shown that majority of
the extremely red objects (I - :i > 2.0), which are considered to be either z > 6.1 quasars
or T dwarfs, are in fact T dwarfs. In this Chapter, we examine moderately red candidates
(0.7 < I - :i < 2.0) to search for quasars at lower redshift range (5.7 < z < 6.1) selected
by the optical :i filter.
3.2 QSO Candidate Selection in HQT Suprime-Cam field
The imaging portion of the HQT survey was carried out using a mosaic CCD camera
Suprime-Cam on 8.2-m Subaru telescope. The location of the field and the survey imaging
observations with Suprime-Cam are described in §3 of Chapter 2. As is discussed in §2 of
Chapter 2, quasars at z > 5.7 are selected as I-band dropouts due to strong IGM absorption
below Lyo emission. In the HQT survey, we use a narrowband filter N B8I6 to break a
well-known I-:i color degeneracy between high-z quasars and MIL dwarfs, which are much
more numerous than quasars.
Our color criteria for z > 5.7 quasars are expressed as:
and
1- z' > 0.7
N B8I6 - z' > 1.7
63
where magnitudes are in AB system.
These criteria essentially mean that high-z q1lllB8J'S are [ and N BSI6-dropouts (i.e.,
faint in [ and NBS16 in comparison with z'). Particularly in NBS16, they are expected
to be faint at least by 1.7 mag with respect to their z' -band magnitudes, suggesting that
we need deeper images in this bandp8BS. For example, with the magnitude criterion of
ziim' the faintest candidate will have a z'-band magnitude zfaintest = z!im' In order for this
object to satisfy the color criteria for quasars, its magnitudes should be [faintest> z!im +0.7
and NBS16/aintest > z!im + 1.7. In other words, the candidate must not be detected at
[ = [faintest and NBS16 = NBS16faintest at 20-. In Chapter 2, we selected candidate
T dwarf and z > 6.1 q1lllB8J'S at z!im = 23.3. This has restricted our search to 9.3deg2
area whose 20' limiting magnitudes are [2<1 > 24.0 (= 23.3 + 0.7) and NBS162<1 > 25.0
(= 23.3 + 1.7).
Here we select z > 5.S quasar candidates with z~B < 23.3 from the same area but
excluding one Suprime-Cam field for HQT -1 in Chapter 2, since the spectral analysis of
this source is still underway. This has reduced our survey area to 9.1 deg2. Another quasar
candidate HQT -7 selected in Chapter 2, was rejected as its magnitude is z'AB > 23.3. Since
high-redshift quasars are expected to be very faint in [ and N BSl6-band, we allowed objects
that were not detected in [ and/or NBS16 at 20' level. Every candidate which satisfies the
color criteria for quasars was inspected by eye to ensure that it is not a spurious source due
to cosmic rays, CCD defects, and spikes around bright stars. Most of our Supcime-Cam
I, NBS16, and z'-band data were obtained nearly simultaneously, allowing us to remove
transients such as Kuiper Belt objects or supernovae. After exercising these, we are left with
three quasar candidates. Their coordinates and photometric properties are summarized in
Table 3.1. Figure 3.1 shows N BS16-z' VB. [-z' diagram. Our candidate quasars are shown
as red boxes with annotated numbers corresponding to "ID" in Table 3.1. All three objects
were not detected in N BS16 at 20- level, and therefore only lower limit in N BS16 - z' were
determined. In Table 3.1, we also list one candidate quasar (HQT J222037.35+014558.05;
ID = 4) that was rejected in above procedure because of its faint magnitude (z'AB > 23.3),
64
Table 3.l. Candidate High-redshift (z > 5.7) Quasars
ID Object lAB NB816AB .iAB JMKO
1 HQT Jl64554.25+65301O.93 25.0 ± 0.42 > 24.96 23.1 ± 0.09 23.2 ± 0.21 2 HQT J221432.35+003119.05 24.8 ± 0.24 > 25.30 23.1 ± 0.06 > 23.7 3 HQT J163552.45+660706.71 23.7 ± 0.17 > 24.63 22.8 ± 0.09 4 HQT J222037.35+014558.05 24.4 ± 0.20 > 25.10 23.5 ± 0.10 > 22.5
Note. - Photometry of the HQT candidate quasars. The HQT candidate quasars designations are given as "HQT Jbbmmss,ss±ddmmss.ssn
, in which suffix is the equatorial coordinates in J2000.0. First three candidate quasars in the list (ID = 1 - 3) were chosen from the 9.3 deg2 area with the magnitude limit of .i :S: 23.3. Keck/DEIMOS optical spectra were obtained for these candidates, and we use these sample to constrain the quasar luminosity function at z ~ 6. The fourth candidate (HQT J222037.35+014558.05; ID = 4) was rejected due to its faint .i-band magnitude (.i > 23.3), but was observed with Keck/DEIMOS. Note that optical magnitudes are expressed in AB system, while J-band magnitude is in Mauna Kea Observatories (MKO) nea.r-ffi systyem (Tokunaga & Vecca 2005). which is a Vega-based photometric system.
but was later observed with Keck/DEIMOS. This candidate is shown as a pink box in
Figure 3.1. Figure 3.2 shows the images of the HQT candidate quasars in I, N B816, .i,
and J-band (from left to right). z > 5.7 quasars are 1- and N B81S-band dropouts with
weak or no detection in all but the .i-band. For the later analysis on the quasar luminosity
function, we only use the first three candidates in Table 3.1 which were selected by the
uniform criteria.
3.3 Follow-up Observations
3.3.1 Near-IR Imaging Follow-up
In other z > 5.7 quasars surveys which employ the I-band (or i'-band) dropout technique,
.i - J color has been used to remove contaminants (mainly M and L dwarfs) from the
quasar candidates. Our new color selection with a NB816 filter does not require J-band
65
3
-...
o
- 1
-z > 5_7 OSOS - Stcrs and 80s
• Stars (0 - loll + L dwarfs lIE T dwarf.
HOT candidate OSOs 00 Complete • Incomplete
o
•
, L ______________ _ ___ _
2 3 4 5 NBe16 - z' [AB)
Figure 3.1 N llS16-Zl vs [ - Z l diagram for candidate quasars. Foreground objects are shown as asterisks (T dwarfs), crosses (L dwarfs), and circles (main sequence star ). The solid blue line is the t rack of simulated quasar colors as a function of redshift. The pink line is a track of synthesized colors of stars and ultracool dwarfs (M5 - T S) . Shaded area is the selection region for very high-z (z > 6.1 ) quasars and T dwarfs, and was discllssed in the previous Chapter. Area surrounded by dotted lines is the expected color space for z > 5.S quasar. First three candidates in the Table 3.1 were selected by the uniform cri teria and are plotted as red boxes ("Complete Sample") . HQT J222037.35+01455S.05 ( "Incomplete Sample") was rejected due to its faint magnitude, but Keckj DEIMOS spectrum was obtained for this candidate, and is shown as a pink box in the present diagram. Note that none of our candidates wa~ detected in N BSl6-band at 2(7 level, and thus we could only pu t lower limits in NBS16 - z'.
66
, .... j.. ... " •
'R~! jzz~jt I
HQT J 1635+{)607
. ..:-~.
Figure 3.2 Images of the HQT candidate quasars in the I , N 8816, Z', and J fil ters (from left to right). Fields are lO"on a side and north is up and east to the left. The locations of the candidate quasar are indicated with circles. These candidates were selected because of their red N 8 816 - Z l and 1 - z' colors (i.e., very faint in I and N 8 816 with respect to Z l ),
which are suggestive of z > 5.7 quasars . The J·band image of HQT J163552.45+660706.71 was not obtained therefore is not shown in th is figure.
67
Table 3.2. Log of Nea.r-IR Imaging Observation
ID Object UT Date Instrument Exposure Weather [sec] [seeing]
1 HQT Jl64554.25+653010.93 2007.06.12 MOIRCS 720 Photometric [0.5"] 2 HQT J221432.35+003119.05 2007.06.12 MOIRCS 720 Photometric [0.5"] 4 HQT J222037.35+014558.05 2005.05.09 WFCAM 240 Photometric [2"]
data to remove such foreground dwarfs. However, it is worthwhile to check the Zl - J colors
of the candidate quasars in order to examine if they satisfy the z' - J colors of quasars.
To this end, we have obtained the J-band imaging data for our candidate quasars with
Subaru/MOIRCS and UKIRT /WFCAM. All except one candidate quasar HQT J2220+0I45
has been observed together with the extreme red objects (I - z' > 2.0) discussed in Chapter
2. The data reductions were done in the similar manner described in §3 of Chapter 2.
Table 3.2 shows the observation log. Resulting magnitudes are included in Table 3.1.
3.3.2 KeckjDEIMOS Spectroscopic Follow-up
Spectroscopic observations were obtained for our candidate quasars using the Deep
Extragalactic Imaging Multi-Object Spectrograph (DEIMOS; Faber et al. 2003) on KeckII.
The details of the observations are given in Table 3.3. We used the G830 line mm-1
grating blazed at 8640A witb an OG550 blocker and II/width slitlets. In this configuration,
the resolution is 3.3A and spectra cover a wavelength range of about 4000A with the
center roughly at 7800A. For each DEIMOS mask, observations were taken as three dither
sequences with a 1.51/separation between each dither along the slit. The spectra were
reduced following the procedure described in Cowie et al. (1996).
68
Table 3.3. KeckII/DEIMOS Spectroscopy Observations of HQT Quasar Candidates
ID Name (J2000) Redshlft Type UT Date Exposure Weather [secJ [seeingJ
1 HQT JI64554.25+65301O.93 0.804 [01lI] 2006.06.25 3600 thin cirrus [1.0"J 2 HQT J221432.35+003119.05 0.770 [01lI] 2005.10.07 3600 photometric [0.7"] 3 HQT JI63552.45+660706.71 0.782 [OIIIJ 2006.06.25 3600 cirrus [I.O"J 4 HQT J222037.35+014558.05 4.5 QSO 2006.06.26 3000 thin cirrus [0.6"J
3.4 Results
3.4.1 Strong Emission Line Galaxies at z < 1
Figure 3.3 shows 1- z' vs. z' - J diagram for four quasar candidates listed in Table 3.1.
The HQT J163552.45+660706.71lacks in the J band data, and therefore is not plotted in
this figure. An area surrounded by dotted lines is the expected color space for z > 5.7
quasars, and the candidate quasars fall into this area. This suggests that these are not
likely foreground dwarfs, and that the N B816 - z' color works just as well as z' - J color
to eliminate M and L stars.
Keck/DEIMOS spectra of the candidate quasars show, however, that none of them is
quasar at z > 5.7. Instead, they are mostly strong emission line galaxies at intermediate
redshift (z < 1). The redshifts and object types are summarized in Table 3.3. We discuss
the spectrum of each quasar candidate below.
HQT Jl64554.25+653010.93
Figure 3.4 presents the spectrum of HQT J164554.25+65301O.93 with identified emission
line feature overlaid. This candidate turned out to be a galaxy at z = 0.804 with strong
[0IlIj'>'.M959,5007 emission falling into the z'-band. However little continuum is detected.
These features makes this galaxy 1- and N B816-band dropouts, resembling the colors of
high-z quasars. Interestingly, [0IlJ.>.3727 is not detected. Its high [OIII]/[OII] ratio suggests
that this galaxy has a high ionization parameter.
69
4
~ 3 0 a> ., ~ ..,
2 ~
ro « ~
'N
0
- 1
-z > 5.7 aso --Stars and 80s
Stars + l dwarfs • T dwarfs
: ..... J:' ....
... ~r . \ .
0 2 I - Z' (AB]
•
... ... .... . -, .... . _ ....
6.30 ..,)!.40
3 4
Figure 3.3 1 - z' versus. z' - J diagram shown with the same symbols as in Figure 3.1. All quasar candidates except HQT J 1635+ 6607 (ID4) for which we do not have J band data, are plotted as red boxes ("Complete Sample") and a pink box ("Incomplete Sample") with the ID numbers defined in Table 3.1. The optically selected HQT candidate quasars fall into the area for high· z quasars, demonstrating that the N B816 - Z l color works just as well as the Z l - J color to separate quasars from foreground !VI and L dwarfs.
70
HQT J221432.35+003119.05
This candidate is also a galaxy at intermediate redshift (z = 0.770) selected because of its
strong [0111],\,\4959,5007 in the z' -band (Figure 3.5). [011]'\3737 is weakly detected.
HQT Jl63552.45+660706.71
Figure 3.6 shows that this candidate is another example of a strong emission line galaxy
(z = 0.782) mimicking the colors of high-z quasars.
HQT J222037.35+014558.05
This candidate was rejected by our magnitude cut-off criterion (zfim = 23.3), however
was observed with Keck/DEIMOS, and thus we present the spectrum as reference. This
candidate turned out to be a quasar at z = 4.5 selected as an I-band drop-out because its
strong CIV emission falls into z'-band.
71
, , 1
HOT J1645+6530 ~ -
150 - (5 "0 '" J> (5 -
1 l l l 1 -;: 100 - -:J 0 u
50 r- -
0 J ; H, 11 -
'I r 1 1
6000 7000 8000 9000 10000 Observed Wavelength [A 1
Figure 3.4 Keckj DEIMOS spectrum of HQT J 164554.25+65301O.93. The identified emission-line features are labeled and marked with the dotted lines. T he plot shows this candidate quasar is a galaxy at z = 0.804, selected as I-band dropout because of its strong [OIlIJ>.>.4959, 5007 emission lines falling into the z' fil ter. The continuum is not detected so cloes not [0 IlJ>.3727 line.
72
800 I I
'" HQT J2214+0031 <0
'" ~
f ~
- ~ (5 (5 Z "0 .c
600 t- 1 1 l l ~ -..
.. -c: 400 t- -:l 0
U .. .. ..
200 - .. -..
0 - i i. i i l 'T
I
6000 7000 8000 9000 10000 Observed Wovelength [A 1
Figure 3.5 Keck/ DEIMOS spectrum of HQT ,)221432.35+003119.05. This candidate quasar is a galaxy at z = 0.771 selected because of its strong [OIII) line emission falling into z' filter. Similarly with HQT J164554.25+65301O.93 (Fignre 3.4), continuum is not detected despite its strong emission li nes. [OII)A3727 is weak , but is detected,
73
-c: :J o
40 -
30 -
u 20-
10 ;-
Or
6000
,
HOT J 1635+6607
-
-
-
" I .. Jill . .1J •
" , ". T ( or. I ".I 'I' , , T '1 7000 8000 9000 10000
Observed Wavelength (A 1
Figure 3.6 Keck/ DE1MOS spectrum of HQT Jl63552.45+660706.71. This presents another e.xarnple of low-z strong emission line galaxy mimicking the opt ical colors of high-z quasars. This galaxy at z = 0.782
74
10
HOT J2220+0145
8
6 l' , -<: ~ 0 U 4
6000 7000 8000 9000 10000 Observed Wavelength [A 1
Figure 3.7 Keck! DElMOS spectrum of HQT J222037.35+014558.05. This candidate is a quasar at z = 4.5, selected because of its strong ClY emission li nes mimicking the I - z' color of z > 5.7 quasars .
75
3.4.2 Limits on the Quasar Luminosity Function
Little is known about the faint-end slope of the quasar luminosity function (LF) at z ~ 6,
since all known high-z quasars are very luminous (Ml450 < 25) and hence they can only
determine the bright-end slope. The HQT survey covers significantly larger area than the
other surveys which probe similarly faint depths. Thus we can use the nondetection of
high-z quasars to give more stringent constraint on the faint-end slope of the LF.
We can parameterize the quasar LF using the standard double power law form that
provides a good fit to the quasar LF at lower redshift (Croom et al. 2004):
where ~(Mi450) is the density norma1ization, Mi450 is a characteristic ("break") magnitude,
a is the faint-end slope, and {3 is the bright-end slope. For the bright-end portion of the
quasar LF, we use the recent result from Jiang et al. (2007) who constructed the quasar
LF at z ~ 6 based on the luminous quasars from the SDSS main survey (e.g., Fan et al.
2006) and five new quasars from the SDSS deep survey. Their LF is expressed by a single
power law:
~(M145(h z) = ~·1O-O.4(JH1)(Ml450+26). (3.2)
where cJ). = (5.2 ± 1.9) x 10-9 Mpc-3mag-l and {3 = -3.1 ± 0.4. Given small number
of sample quasars and narrow redshift range (5.7 < z < 6.4), they neglected the redshift
evolution and only considered luminosity dependence on the LF.
The expected number of quasars can be derived by integrating quasar LF over the survey
volume and magnitude:
1""""1-00
dV nobs(Zmin < z < z".",,< M) = ~(Ml450,Z)ddMdz, Zmin M Z
(3.3)
76
which is a function of faint-end slope ct and break magnitude Mi4llO. Upper limits for
the number of quasars in the HQT survey can be derived from the nondetections based
on Poisson statistics. According to Gehrels (1986), the Poisson single sided upper limits
for zero event (nobserved = 0) is nupper = 4.605 at a confidence level of 0.99. This means
that there is 1% probability of observing zero number of quasar when there are actually
roughly 4.6 quasars in the field. We can relate this nuyper at different confidence levels with
equation 3.4.2 to set limits on ct and Mi450.
We set the range of Mi4llO and ct using the observational evidences. The number density
of luminous quasars is known to decline at redshift beyond 3 (Fan et al. 2004). This decline
can be explained by either pure density evolution (PDE), pure luminosity evolution (PLE),
or the hybrid form of the two - luminosity-dependent density evolution (LDDE) suggested
by X-ray surveys (e.g., Ueda et al. 2(03). In either scenario, the break magnitude at higher
redshift is expected to be fainter than - or in case for the PDE scenario at least the same as
- what is found at lower redshift. Hence in our analysis, we set a brightest break magnitude
to be Mi4llO = -25.0, which is a value at z = 2 (Croom et al. 2004). As for ct, we require
it to be less than -1.0 since fainter quasars are expected to be more common than than
luminous quasars.
The result of our analysis is shown in Figure 3.8. Contours are plotted at a confidence
level of 0.99, 0.95, and 0.90. This figure shows that at Mi4llO > -25.0, the nondetection
of high-z quasar in the HQT survey limits the faint-end slope ct to be greater than -2.0
at a 99% confidence level. As comparison, we also plot the constraints from the CFHQS
Deep (Willott et al. 2(05) whose survey depth is similar to ours. It can be seen that
our constraints are significantly stronger than that from the CFHQS because of our survey
volume being about three times larger. The lower limits on ct and Mi4llO at different
confidence levels are summarized in Table 3.4.
To show our constraints more visually, we plot the quasar LF in Figure 3.9. The densities
of luminous SDSS quasars at z '" 6 are plotted as triangles ([SDSS-main] Fan et al. 2006;
[SDSS-deep) Jiang et al. 2007). The solid line shows a bright-end part of the quasar LF from
77
Table 3.4. Faint-end Slope and Break Magnitude for the z = 6 Quasar LF at Several Confidence Levels
Survey c:tmin Mi450,bright Area 90% 95% 99% 90% 95% 99% [deg2)
CFHQS Deep -2.3 -2.5 -2.8 -24.1 -24.0 -23.7 3.32 HQT -1.6 -1.8 -2.1 -24.6 -24.5 -24.3 9.0
Note. - Minimum faint-end slope c:t and break magnitude Mi450 constrained by the HQT survey and CFHQS deep survey (Willott et aI. 2005) at different confidence level. The HQT survey gives more stringent constraints due to larger survey volume.
Jiang et aI. (2007; equation ). We also plot the quasar LF at z = 2 from the Two degree
field QSO Redshift Survey (2QZ; Croom et aI. 2004) as comparison with the z = 6 LF. It
can be clearly seen that the number density of the luminous quasars decline from z = 2 to
z = 6 by more than one order of magnitude. The upper limits on the quasar density by our
survey are shown by squares with 99%, 95%, and 90% confidence level from top to bottom.
Similarly, limits from the CFHQS deep survey (Willott et aI. 2005) are shown as diamonds,
and it can be seen that lower upper-limits of our data points give stronger constraint on
the LF than the CFHQS deep data. Negative search of z > 5.7 quasar in the 2 Ms X-ray
Chandra Deep field North (CDFN) by Barger et aI. (2003) is shown as circles. Although
their depth goes down to about two magnitudes fainter than ours, due to the smaller area
coverage (160arcmin2 ), they cannot constrain the faint-end slope well. Finally, blue dotted
lines indicate possible range of the faint-end slopes c:t a.nd break magnitudes Mi450 allowed
at a confidence level of 99% by the HQT survey. Red dashed lines shows the ra.nge of Mi450
set by the constraints of Mi450 > -25 a.nd c:t < -1.0.
78
- 23.8 __ CFHQS Deep
-24.0 __ HQT
0 -24.2 of) ...
:::;; - 24.4
-24.6
- 24.8
-25 .0 - 3.0 - 2 .5
..
, ,
, ,
..'
, , ,
- 2.0 ex
\1\'> ..... I) ..
....
, , ,
.. ... ... .. .. ....
.... r:;.~';> " ' " , ,
......... <:;'fl" .' ,
......... , ,' ..... "
- 1.5 -1.0
Figure 3.8 Constraints on the faint-end slope (0) and break magnitude (Afi4W) by the HQT survey (rcd lines) and CFHQS Deep survey (black lines) assuming bright-end LF [rom Jiang et aI. (2007). For each su rvey, constraints are drawn at a confidence level of 99% (solid line), 95% (dotted line) , and 90% (dashed line). It can be seen that the HQT survey can give more significant constraints than the CFHQS deep survey.
79
I
"" o ",E 10-7 ,
u Q.
::IE :1 10-8
::IE ~
COF
i
Confidence Level = 99%
-22 -24
>-.p.S<;!DSS deep ....... ..... ...
-26
SDSS ma i
z=6S0SS (+ « L-"')
-28
Figure 3.9 Quasar luminosity fnnction at z = 6. The triangles represents the density of z > 5.7 SDSS quasars reported by Jiang et aI. (2007) and Fan et aI. (2006). The z = 6 quasar LF of Jiang et aI . (2007) is shown as a solid line (41 ()( L - 3.,). The luminosity function at z = 2 from the 2dF QSO Redshift Survey (2QZ; Croom et aI. 2004) is plotted as a dash-dot line for comparison. In the faint-end part , the upper limits on the quasar density are shown for the HQT survey as squares with 99%, 95%, and 90% confidence level from top to bottom. Likewise, upper limits from the CFHQS deep (Willott et al. 2005) and Chandm deep field north (CDFN; Barger et al. 2003) are shown as diamonds and circles, respectively. Blue dotted lines are plausible LF at faint-end constrained by the HQT upper limit at a 99% confidence level, and it is shown that the slope must be shallower than CI' = - 2.1 in order to satisfy our data point at the break magnitude Mi450 > -25.0. Vertical dotted lines show the range of possible break magnitude (-25 < Mi450 < -24.25) constrained by our data point .
80
3.5 Discussion and Conclusion
We carried out an optical wide-field survey with Subaru/Suprime-Cam to search for the
low luminosity quasars (M1450 < -23.5) at 5.7 < z < 6.4 in an area of 9 deg2. The
follow-up spectroscopy with Keck/DEIMOS hllB shown that all three candidates selected by
the uniform criteria are galaxies at intermediate redshifts (z < 1). These galaxies exhibit
extremely strong [0111].>.5007.>.4959 lines, but show little continuum. These features ma.lre
them z'-band single detected objects (i.e., I, NB816, and J-band dropouts), resembling the
optical and near-IR (z' - J) colors of high-z quasars. Such galaxies at z < 1 with strong
emission lines (e.g., Ha, [OIl], and [OIIIJ) have been known to be mistaken for high-z Lya
emitting galaxies (e.g., Stockton et al. 1999). Galaxies at z » 5 should not have flux in
the bluer bandp8BBeS (e.g., B, V, R-filter) due to strong IGM absorption below Lya, while
low-z galaxies are expected to show some stellar continuum in these bandp8BBeS. Therefore,
in current high-z galaxy surveys, deep intaging data in blue filters have been used to remove
these low-z galaxies from the candidate high-z gala.xies (e.g., Hu et al. 2004).
In other qUaBar surveys which utilize i - z' and z' - J color selections (e.g., SDSS and
CFHQS), however, these emission line galaxies have not been reported lIB contaminants.
The reaBOn for this could be that current qUaBar surveys only probe the bright end of the
qUaBar population (M1450 < -25), where luminosity range is too bright for these galaxies
to be detected. Our findings of these galaxies lIB high-z candidate quasars may address a
possible problem with the current color selection technique for low luminosity quasars, since
the colors used to select z > 5.7 quasars (I - z' and z' - J or N B816 - z' in case of the
HQT survey) cannot separate these galaxies from high-z qUaBar candidates.
The nature of these strong emission line galaxies have not been well investigated, lIB they
were often neglected by high-z researchers lIB being contantinants. Their high [OIlIJ/[OIl]
ratios suggest extremely high ionization status. Figure 3.10 shows R23 versus 032 emission
line diagnostic diagram from Kewley et al. (2002) and one of our strong [OlII] emitting
galaxies (HQT J221432.35+003119.05) is shown lIB an IIBterisk. It is clearly seen that this
81
gala.xy hOB an extremely high ionization parameter (q = 3xlOS [em/s]), and is well separated
from the stanonning galaxies found by the SDSS (dots) and GOODS (diamonds) surveys
(Kobulnicky & Kewley 2004), indicating the rareness of these populations. The SDSS
data were taken from http://www.mpa-garching.mpg.de/SDSS/DR4/SDSS-line.html The
narrow line widths (~ 150 km/s) of these objects exclude the possibility of AGNs. Using
[NeIIIlIH,6 VB. [OIIIlI[OII] diagnostics (Figure 3.11), the metallicities are meaBured to be
about 0.2 - 0.3 Z0' Together with their high H,6 equivalent widths (EWrest > 200.11),
the most natural interpretation for these galaxies is that they have just undergone major
starbursts at leaBt within a few Myrs according to the Starburst 99 gala.xy evolution model
(Leitherer et al. 1999). Furthermore, star formation rate of these emission line galaxies are
considerably high (~2O-3O Msunlyr), suggesting that their contribution to the global star
formation activity at z ~ 1 may not be negligible. Motivated by these findings, in Chapter
4, we investigate the properties of these strong emission line galaxies in more systematic
manner using the deep narrowband data from Hu et al. (2005).
The non-detections of z > 5.7 qUaBarS in the HQT survey hOB been used to set limits on
the faint-end slope (O! > -2.1) and the break magnitude (Mi450 < -24.3) at 99% confidence
level. Despite many attempts, no z > 5.7 qUaBar hOB been found so far at Ml450 < -25.0,
which is a break luminosity of qUaBar LF at z = 2. Since the HQT survey covers significantly
larger area than the other surveys for low luminosity quaBarS, our constraints are more
significant than the previous work. The dominant contribution to the energy emitted by
the qUaBar population comes from those qUaBarS around M = M·. Therefore our results
should better constrain the quasar contribution to reionization, and this will be our future
work.
82
1.0
~
'" 0.5 N
'" ~ Cl 0 ...J
0.0
-0.5
-2
Z-0.4
Z:\.O
o .
-1 o LOG(032)
soss golox;es (0.04 < z < 0.·4)
lo [;oQ[lS galaxies (0.3 < z < I .a)
HOT J2214+0031
2
Figure 3.10 Diagnost ic diagram for starforming galaxies with various ionization parameters (q [cm/ s]) and metalicities Z (solar units) taken from Kewley & Dopita (2002). Solar metallicity of 12 + log[O/ H] = 8.9 (Anders & Grevesse 1989) was used. One of our candidate quasars (HQT J2214+ 0031), which turned out to be an [OIII] emit ting galaxy at z = 0.77 by the follow-up DElMaS spectrum, is shown as an asterisk. Nearby SDSS starforming galaxies are also plotted as dots and the GOODS starforming galaxies at intermediate redshifts (0.3 < z < 1.0) are shown as diamonds.
83
2
Z=0.05
0.4
Z-20 ~ q=3. ~
15 , 0 ~
15 ~ ~
(!) 0 -' -1 SDSS galaxies
-2
-3 - 2 - 1 o LOG ([NelllJ/Hbeta)
Figure 3. 11 [OUll/[OIlI vs. [NeIIIl/H.i3 diagnostic diagram with variou ' ionizat ion panuneters and metall icities taken from Kewley & Dopita (2002). HQT J2214+0031 is shown as an asterisk and it crul be seen that this galaxy has Z ~ 0.3Z0 with extremely high ionization paramter q = 3.0 x 108 .
84
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86
Chapter 4
Extremely Low-Metallicity Galaxies at z < 1
abstract
Motivated by the findings of extremely strong emission line galaxies in the HQT survey,
we carried out more systematic studies of these galaxies using ultrardeep imaging data
obtained with Subaru/Suprime-Cam. We used two ~ 100A narrowband filters centered at
8150A and 9140A (NB816 and NB912) to search for ultra-strong emission line galaxies
(USELs) with EW(H!3) ~ 30 A. 542 candidate galaxies are found in a half square degree
survey. Followup spectroscopy for randomly selected objects in the candidate sample with
KeckII/DEIMOS shows they consist of [OIIIjA5007, [OIIjA3727, and Ha selected strong
emission line galaxies at intermediate redshifts (z < 1), and Lya emitting galaxies at
high-redshift (z » 5). We determine the H,B 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 [OIIIjA4363
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. Our
high yield rate of XMPGs suggests that the narrowband method is a powerful way to find
such populations. Moreover, the lowest meta11icity measured in our sample is close to the
minimum meta11icity found in local galaxies [12+log(O/H) ~ 7.1j, though we need deeper
87
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.
4.1 Introduction
The study of low-metallicity galaxies is of considerable 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 formation 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 40; Thuan & Izotov 48; Izotov et al. 13). More
recently, the Sloan Digital Sky Survey (SDSS) has yielded additional extremely metal-poor
galaxies (XMPGs) (12+log (O/H) < 7.65 or Z < Z0/12; Kniazev et al. 21; Izotov et al.
14). Despite enormous efforts, only a few dozen such XMPGs are known, all at redshift
z < 0.05 (e.g., 32; 15).
Historically, objective prism surveys have been used to select emission-line galaxies for
low-metallicity studies (e.g. the Hamburg QSO Survey (37) and its HSS sequel (50) that
discovered HS 2134+0400 (38) and the Kitt Peak International Spectroscopy Survey (KISS;
Salzer et al. 39; Melbourne & Salzer 29)). The advantage of using the objective prism
technique rather than the continuum selection, employed with the SDSS (21) or DEEP2
surveys (9), 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 analysis of 250,000 spectra over ~3000 deg2 for the SDSS
88
(21)), since low-metallicity populations in their first outburst have weak continua 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 historically been picked
up in high-z Lyman alpha searches (e.g., 47: 10: 11; 46: 49; 2) where they have been
considered contaminants. However, 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 Chapter.
While some spectroscopic studies have been carried out for low-redshift galaxies selected
from narrowband surveys (e.g., 28: 27), the small sample sizes have precluded any detailed
investigation of metallicity and identification of a low-metallicity population.
The narrowband method probes to much deeper limits than the objective prism surveys.
This enables probing star-forming populations out to near redshift z ~ 1 where the cosmic
star formation rates are considerably higher. Furthermore, the narrowband emission-line
selection allows us to assemble very large samples of strong-emission line objects, with a
clean selection of different line types for the construction of luminosity functions.
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 controversy about the interpretation of the low
metallicity measurements 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., 4) or, at least,
resulted in identifying low metallicity systems which are not forming their first generation
of stars. The identification of low metallicity galaxies - at the level of the XMPGs -
among massive galaxies may provides less ambiguous examples of galaxies that are genuinely
'young' and caught in the initial stages of star formation. Current efforts to identify low
metallicity galaxies from continuum selected surveys (e.g., 23: 26: 22: 9) have low-metallicity
thresholds that are higher than this - about one-third solar (in O/H). With a narrow
band selection criterion much larger emission-line samples including very low metallicity
galaxies can be identified. With these large samples it is also possible to determine whether
89
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
SupcimeCam mosaic CCD camera (30) 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 resulting 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 using multi-object masks with the DEIMOS spectrograph (6) on the
10-m Keck II telescope. Sample spectra for each class of object are shown. Flux calibration
and equivalent width distributions are presented 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 corrections) at this epoch. Analysis of
the metallicities is given in §6. Their morphologies are discussed in §7 and a final snmmary
discussion is given in §8. We use a standard Ho = 70 km s-I Mpc-I• Om = 0.3, OA = 0.7
cosmology throughout.
4.2 The Narrow Band Selection
The emission-line sample was chosen from a set of narrowband 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 A (FWHM) filters centered at nominal wavelengths of 8150 A and 9140 A 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.htmiand
are also described in Ajiki et al. (1). 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,
90
• • . ..
r -
r -
~'.-.--.-.-.'~ .. -------_. ... -........ ..................... -, /:"-i , ,
.1 \~ \" , \. ...- .\ ,.
,
5000 6000 7000 8000 9000 10000 WAVELENGTH (A)
Figure 4.1 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 A and 9140A. The present case corresponds to an Ho emission line object found in the 9140A filter (shown with the narrow solid curve). Also illustrated are the broad band V (dash-dot), R (solid), I (dashed), and Zl (dotted) filters use to measure the continuum. The spectrum shown corresponds to object 205 in Table 4.4 and is an Ho emitter at z = 0.3983. The easily visible lines are the Balmer series and the [DIlI] lines at AA5007, 4959, and 4363A.
91
Table 4.1. Narrowband Survey Area. Coverage
Field RA (J2000) Dec (J2000) (l", b") EB_VB NBS16 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" site of" gravitational lensing cluster at z ~ 0.375, and was omitted from the suvey •
• estimated using http://irsa.ipac.caltech.edu/,,pplications/DUST / based on Schlegel et aI. (42)
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
exposures 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 description of the full reduction procedure for images
is given in Capak et al. (3). All magnitudes are given in the AB system (33). 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 Lyct emitters at redshifts of z - 5.7
and z - 6.5 (11; 12) but the narrowband 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 Chapter.
The fields which we use and the area. covered (approximately a half square degree in each
bandpass) are snmmarized in Table 4.2. 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 N B816 < 25
and I - N B816 greater than 0.8. We selected galaxies in the N B912 filter with the z
92
filter DB the reference continuum bandp8B8 and included all ga1axies with N B912 < 25 and
z - N B912 greater than 1. The selection process is illustrated for a galaxy found in the
NB912 filter in Figure 4.1. Both selections correspond roughly to choosing objects with
emission lines with rest-frame equivalent widths greater than 100 A. 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). Tabulated coordinates, multi-color magnitudes, and redshifts (where
me8Bured) for these objects are summarized in Table 4.3. Very few of these objects
would be included in continuum-selected spectroscopic samples. Figure 4.2 shows the
narrowband excess as a function of narrowband (NAB) magnitude for objects with narrow
band magnitudes brighter than 24. The open symbols show the present sample while the
solid symbols show objects which would be included in an R < 24 continuum-selected
sample.
4.3 Spectra
Spectroscopic observations were obtained for 161 USELs from the sample using the Deep
Extragalactic Imaging Multi-Object Spectrograph (DEIMOS; Faber et al. 6) on Keck II
in a series of runs between 2003 and 2006. The emission line objects were included in
mDBks designed to observe high-z Lya candidates and, DB can be seen in the lower panel of
Figure 4.2, constitute an essentially random sample of the emission line galaxies.
The observations were primarily made with the G830 tlmm grating blazed at 8640 A
and used 1" wide slitlets. In this configuration, the resolution is 3.3 A, which is sufficient
to distinguish the [OIl] A3727 doublet structure. This allows us to eaBily identify [0111
A3727 emitters where often the [OIlI A3727 doublet is the only emission feature. The
spectra cover a wavelength range of approximately 4000 A and were centered at an average
93
,..... 2.5
~ All. ~
~ 2.0 D
0 0
• D
~ 0
D <1000
0 0 Z 1.5 0 o &oD6~ 0
I • • o. o.~oO ~8bD • ,. 00 ~~O~ • • . . ~ ... ~
• • .D '-D'L • III 0 u 0.5
20 21 22 23 24 N(AB)
,..... 2.5
~ WTI1I SPECTRA ~ D
~ 2.0 D - 0 -
D
~ 0 00 ~ 1.5 - D 0 _
~ • • • :0,9 DOOS .. o _DD ~~ • • • • od>o 0 DOD: <b4!!
~ 1.0 - • (J.". DPtJ 011
.: ,.D SD a 0 U 0.5
20 21 22 23 24 N(AB)
Figure 4.2 Continuum - Narrow band magnitude versus narrow band magnitude for all objects with narrow band magnitude brighter than 24. The diamonds show the narrowband excess emission magnitude 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.
94
>< ?f ::J ~
i
4000
'f ='
"" ~ :i! '"
! ! ! I
4500 5000 5500 6000 6500 REST WAVELENGTH (A)
~
i 7000
Figure 4.3 Spectrum of an Ha 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 wea.lrer lines. The more important emission line features are labeled and marked with the dotted lines.
95
- . --
..1 1 J .... .J I. .. .• .J.IJIJ L •
~ ~ ~ = g
"I 1 ' 1 % Q. - -
- -
"' -3500 4000 4500 5000 5500
REST WAVELENGTH (A)
Figure 4.4 Spectrum of an [OIIl[ 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.
96
~ §: ~ 'l!
5' .. ~ :t Q. '"
! ! I !
3500 4000 4500 5000 5500 REST WAVELENGTH (A)
Figure 4.5 Spectrum of an [OIlI] 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.
97
0 0 0 0 0 .-
='
~ ~ S-f- 2!;. !" a -
>< I r 0 f- -ti ." ... .1 L Loo1. ,~.o. ___ .1' I
I T' 0 , 0
f- -3000 3200 3400 3600 3800 4000 4200 4400
REST WAVELENGTH (A)
Figure 4.6 Spectrum of an [OIl] 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.
98
wavelength of 7800 A, though the exact wavelength range for each spectrum depends on
the slit position with respect to the center of the mask along the dispersion direction. The
G830 grating used with the OG550 blocker gives a throughput greater than 20% for most
of this range, and ~ 28% at 8150 A. 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 ss we discuss 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. (5) and the final one-dimensional spectra were extracted using a
profile weighting bssed on the strongest emission line in the spectrum. A small number
of the spectra were obtained with the ZD600 tjmm grating giving a correspondingly lower
resolution but a wider wavelength coverage. These observations were centered at 7200 A.
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 spectra are shown in Figures 4.3, 4.4, 4.5, and 4.6.
The measured redshifts are given in Tables 4.3 and 4.4. The narrow band emission line
selection produces a mixture of objects corresponding to Ha, [OIIIjM007, and [0Ilj>'3727
and, at the faintest magnitudes (> 24), high redshift Lya emitters. The number of objects
seen in each line and the redshifts where they are found are shown in Figure 4.7. The
spectroscopically identified sample from both bands contains 13 Ha, 92 [0IlIj>'5OO7, and
23 [OIIj>.3727 emitters. In the remainder of the Chapter we shall focus on the Ha and
[0111]>'5007 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 ss a function of magnitude we have
adopted a magnitude dependent weighting for each galaxy equal to the total number of
galaxies at this magnitude divided by the number of spectroscopically identified galaxies.
99
~ ., N
's <> e.n ~
~ Ii
40
30 ..
ffi
~ 20 ..
10 ~
0 0.0
10"15
10-16 b-
10-17 b-
10-18
0.0
Ha
~ ~
• • • • • • • I •
• •
[OnI]
[Om"
~ 0.5 1.0 1.5
REDSIllFf
.. • • • • • ~ ~ j
t 41 ... --:I
0.5 1.0 1.5 REDSIllFf
Figure 4.7 (a) Distribution of redshifts for the spectroscopically identified sources. [OIII] A5007 emitters are the most common. Since the focus of this Chapter is on intermediateredshift (z ;S 1) strong emission line galaxies, we did not plot the high redshift Lya galaxies (z » 5) in our sample. High-z Lya emitters are discussed in Hu et aI. (11, 12). (b) Flux versus redhift for the spectroscopically identified sample. Squares are Ha, diamonds are [OIII] A5007, and triangles are [OIl] A3727. 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.
100
However, the analysis is not particularly sensitive to the adopted scheme since the fraction
of identified galaxies is relatively constant with magnitude.
4.4 Flux Calibrations
Generally our spectra were not obtained at the parallactic angle since this is determined
by the DEIMOS mask orientation. Therefore flux calibration using standard stars is
problematic due to atmospheric refraction effects, and special care must be taken for the
flux calibration. We therefore employed three independent methods for the flux calibration.
In §4.1 we define the "primary fluxes" of the emission lines used to select the galaxies. These
primary fluxes are computed directly from the SuprimeCam broadband and narrowband
magnitudes. We use these fluxes to derive the luminosity functions of Ha and [OIIIjA5007
emitters at z < 1 (§5.1).
In §4.2 we measure line fluxes from the spectra. Relative line fluxes can be measured
from the spectra without flux calibration as long as we restrict the line measurements to
short wavelength ranges where the DEIMOS response is essentially constant. For example,
one can assume the response of neighboring lines (e.g. [OIIIjA4949 and [OHI]A5007) are the
same and therefore one can measure the flux ratio without calibration. For bright galaxies,
we can absolutely calibrate the fluxes by integrating the spectra and equating them to the
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 [OIII]A5007/[OIIIjA4959 is indeed close to 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(H,6)/f(Ha) of bright Ha emitters
are close to the Case B conditions, suggesting very little reddening.
Metallicity measurements by the direct method require four emission lines that are
widely displaced over the spectral wavelength range ([OIII]AA4959, 5007, [OHI]A4363, and
[0II]A3727). To calibrate these lines, we used neighboring Balmer lines with the assumption
of Case B conditions, and this is described in §4.4.
101
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 eN) and the corresponding
continuum magnitudes eC) 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 luminosity functions for each emitter in §5.1.
Defining the quantity
R = 1O-0.4.(N-C)
the observed frame equivalent width becomes
EW=[R-~l <1>- -
A-X
where <I> is the narrow band filter response normalized such that the integral over wavelength
is unity and A-X is the effective width of the continuum filter. The narrow band filter is
often assumed to be rectangular, in which case <I> becomes l/o-X where o-X is the width of the
narrow band, but as can be seen from Figure 4.1, this is not a very good approximation in
the present case. For very high equivalent width objects the denominator in this equation
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 [OIII]-X5007 line we must include the second member of the doublet
which also lies within the narrow band filter. We have computed these cases assuming the
flux of the [OIII]-X4959line is 0.34 times that of the [OIII]-X5007 line. Then <I> = <1>1 +0.34 x </>2
where <1>1 is the filter response at the redshifted 5007 A wavelength and </>2 is the filter
response at redshifted 4959 A.
102
50
40
~ 30
~ 20
to
;.
o 1.5
10
8
IE 6 m ::E ::::I 4 z
2
I-
I-
I-
l-
o 1.5
.;
L .;
.;
I .;
2.0 2.5 3.0 LOG REST FRAME EW ([OIll]5007)
-
-
-
I 2.0 2.5 3.0
LOG REST FRAME EW (Ha)
Figure 4.8 (a) Distribution of the rest frame equivalent widths determined from the narrow band magnitudes for the spectroscopically identified [OIlI] >'5007 sources. (b) Distribution of the rest frame equivalent widths for the Ha selected sample.
103
The distribution of the rest frame equivalent widths for the Ha and [OIlI]A5007 samples
is shown in Figure 4.8. The [OIlI]A5007 sample selects objects with rest frame equivalent
widths above about 100A while the lower redshift Ha sample selects objects with rest frame
equivalent widths above about 150A. Since the [OIIIjA5007 lines are also generally stronger
than the Ha lines the [0IIl] selection chooses less extreme objects than the Ha selection
and will include a larger fraction of galaxies at the given redshift.
The high observed frame equivalent widths make the line fluxes insensitive to the
continuum determination and these may simply be found from
[
1O-0.4N - 1O-0.4C 1 J=A 1
</>- -AA
where A is the AB zeropoint at the narrow band wavelength in units of erg cm-2 8-1 A-I.
The flux depends on the filter response at the emission line wavelength and correspondingly
is most uncertain at the edges of the filters 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 (<1». We use these primary
fluxes to construct the luminosity functions of Ha and [OIIIjA5007 selected emitters as we
discuss in §5.1.
4.4.2 Line Fluxes from the Spectra
For the short wavelength range where DEIMOS response 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 measured 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.
104
Figure 4.9 The ratio of the [OIII] .M959 line to [OIII] A5007. 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.
In follow-up work to the present analysis we have obtained much deeper spectra for
several of the objects. Up to six hours of data have been obtained in the form of multiple
one hour exposures with the DEIMOS spectrograph on Keck. The discussion of this data
will be given in Cowie et al. (2007, in preparation). We have used this data to test the noise
analysis in the individual spectra by comparing this noise with the dispersion among the line
flux measurements in each of the individual spectra comprising the multiple observations.
The two methods give consistent determinations for the noise and in all cases the newer
measurements of the line strengths in the combined deeper spectra are fully consistent with
the present observations though with much smaller error bars.
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 calibration for neighboring lines and the noise measurement by measuring the ratio
105
10.00
• • 1.00 • 0 +, • i +
0.10
0.01 L..:::------"-:':-------' 1 0-17 1 0-16 10-15
FLUX (erg cm-2 5-1)
Figure 4.10 Ratio of fluxes computed from the spectra and the broad band magnitudes versus those from the narrow band magnitudes. Ha lines are shown as solid boxes, [01lI] .>.5007 lines as diamonds and [OIl] A3727 lines as crosses.
of the [01lI] A4959/ [01II]A5007 lines. This is expected to have a value of approximately
0.34. The ratio is shown as a function of the [01lI] >'5007 flux in Figure 4.9. The measured
values scatter around the expected value and the dispersion is consistent with the noise
determination. This result supports our assumption of [01lI]>.4959/[01II]>.5007 = 0.34 in
the primary fluxes measurements described in §4.1. However, a small upward trend can
be seen in the ratio as we move towards weaker fluxes. This is consequence of the full
Gaussian fits overestimating the strength of weak lines and is the reason why we force fitted
the position and line widths of the Gaussian and allowed only the norma1ization to vary
when fitting the weak lines. We have adopted a very conservative 10% systematic error in
the [01II] 4959 line where the full Gaussian fitting procedure was used.
The brighter objects may be absolutely calibrated using the continuum magnitUdes
obtained from our Subaru data. We integrated the spectrum convolved with each
SuprimeCam filter response and set this equal to the broad band flux to normalize the
106
1.00 ~ • . - • • • -. 0.10 ~
0.01 I..;o---....... ;-::;----""'"r.r-----' 10-18 10-17 10-16 10-15
f(Ha) erg cm·2 S·1
Figure 4.11 The ratio of the 8.8/80: fluxes versus 80: 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 80: line with continuum magnitudes above 24.5 in the bandpasses corresponding to the lines.
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 (0 < 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 4.10 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.
107
4.4.3 Balmer Ratios
We now measured the Balmer ratios for the sample of objects selected in Ha and where
the continuum magnitudes were bright enough to absolutely flux calibrate the spectra. The
ratio of f{H,8)/f(Ha) is shown in Figure 4.11. The values average closely to the Case B ratio
which is shown as the solid line and at brighter fluxes the individual 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.
4.4.4 Final Flux Calibration for Metallicity Analysis
For the metallicity analysis we adopted the procedure of normalizing the longer wavelength
lines ([OIIIj,u4959. 5007. [OIII).M363) to their nearest Balmer line to determine the
unreddened fluxes. For example, in the case of the Ha emission selected galaxies, we can
measure Ha absolute fluxes by the primary fluxes method described in §4.1. We can then
derive Htl and H'Y fluxes from Ha fluxes by assuming Case B recombination [e.g., f(Ha) =
2.85 x f(HtI), f(H'Y) = 0.469 x f{H,8) at T = 104 K and Ne ~ 102 - 104cm-3; Osterbrock
34). As H.8 and [OIII)AA4959, 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:
!o(H.8) _ !(HtI) !o([OIII]A4959, A5(07) - !([OIIIjA4959, A5(07)
where !o(H.8) and !o([OIIIjA4959, A5(07) are the flux counts in the un-calibrated, reddened
DEIMOS spectra, while f(H.8) and f{[OIIIjA4959, A5007) are absolute, unreddened fluxes.
Since we know f(H.8) from f(Ha) with the Case B assumption and !o(H.8)/ !o([OIIIjA4959,
A5007) from the DEIMOS spectra, we can derive f([OIII)A 4959, A5007) using this simple
108
formula. Similarly, we can absolutely calibrate [0111] . .\4363 lines by using its neighboring
Balmer line, H'Y:
Jo(H'Y) J(H'Y) Jo([0II1]A4363) = -:CJ (=[O='==III]""A:'-':43=6=3)
where Jo(H'Y) and Jo([OIlI] A4363) are again the counts in flux uncalibrated, reddened
DEIMOS spectra, and !(H'Y) and !([OIII] A4363) are absolute fluxes.
In the case of the [0111] selected emitters, we first derive [OIlI] AA4959, 5007 absolute
fluxes using the primary fluxes method (§4.1), and then use the above formula to get absolute
fluxes of H,8, then H'Y (by the Case B ratio), and finally [OIII]A4363.
This flux calibration technique using neighboring Balmer line should work well for the
[OIII]AA4959, 5007, A4363 lines and the [NIl] lines which all lie close to Balmer lines but
may be slightly more approximate for the [SII] lines. The higher order Balmer lines are
too uncertain to apply this procedure due to inadequate SIN of the lines, and we have
used the continuum flux calibrated values with no reddening for the [OII]A3727 and [NeIll]
lines. These values will have correspondingly higher flux uncertainties. Fortunately the
[OII]A3727 line is very weak in most of the objects and the uncertainty has little effect
on the metallicity determinations. However, ionization analyses based on the (NeIll] line
should be undertaken with caution.
4.5 Star Formation History
4.5.1 Ha and [0111]).5007 Luminosity Functions
Because of the high observed frame equivalent widths the primary fluxes are insensitive
to the continuum determination. However, they do depend on the filter response 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 redshift
109
10-2
~ z=O.24
bIl 10-3 ~
~ I 10-4
Z
10-5
40 41 42 43 L(Ha) (erg S·I)
z=O.39
.
10-5
...... --~ ......... -~--........ ~-~ .......... 40 41 42 43
L(Ha) (erg S·I)
Figure 4.12 The luminosity function of Ho: 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.
110
10-2
;.... z=O.63
'iib 10-3 • • + 8 ~
... Cl Cl '8., Cl
::is
i 10-4 ~
10-5
40 41 42 43 L(Oll) (erg S·I)
10-2
;.... z=O.83
l 10-3 • • .. Cl Cl
+ ... '8.,
t ::is Cl
.8 10-4 .. E
~ = Z
10-5
40 41 42 43 L(Oll) (erg S·I)
Figure 4.13 The luminosity function of [OIII] >.5007 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.
111
range. Now the volume is simply defined by the selected redshift range for all objects
above the minimum luminosity which we take as corresponding to a flux of 1.5 x 10-17 erg
cm-2 s-1 (Figure 4.7). The luminosity function is now obtained by dividing the number
of objects in each luminosity bin by the volume. The incompleteness 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 Hll< luminosity function is shown for the two
redshift ranges corresponding to the NB816 and NB912 selections in Figure 4.12 and the
corresponding [OIII]A5007 luminosity functions in Figure 4.13.
4.5.2 Star Formation Rates
The individual objects have Hll< luminosities stretching up to about 3 x Wi! erg s-1 where,
at the higher redshifts, we use the H,8 luminosity to infer the Hll< value 8BSuming 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 (18) conversion rate for his Salpeter m8BS
function.
Since the objects are more probably caused by starbursts the true star formation rate
will depend on the evolutionary history. However, the Hll< luminosity density should give
a reasonable estimate of the universal star formation density of the objects provided ouly
that most of the ionizing photons are absorbed in the galaxies. We first formed the total
Hll< or H,8 luminosity density of the galaxies by summing over the incompleteness weighted
luminosities in each redshift interval. We only included detected objects and did not attempt
to extrapolate to lower luminosities but the result are not particularly sensitive to this
because the luminosity functions are relatively flat at the lower luminosities (Figure 4.12
and Figures 4.13). We then converted these to star formation rates with the Kennicutt (18)
conversion.
The results are shown in Figure 4.14. We first plot the rate for the total Banlples at each
redshifts shown by the open squares. We have shown star formation rates for UV continuum
112
~
'" I u D-~
I '-->-0 ~ ~
'" "-til Q.
10- 1 <-
10- 2
10- 3 ;-
10- 4
0.0
.. I I
+
'f
+
0. 5
.. : ;$ + I
+ ... ~
C • Ha measurements C • ~ • A. UV measurements
C USELs
• USELs w;lh EW(Ha»200A
• 1.0 1.5 2 .0 z
Figure 4.14 The star formation history inferred from the Ha or H,i3 luminosity density as a function of redshift. The data from our s!Unple 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 Ha rest frame equivalent widths in excesS of 200A or H,i3 rest frame equivalent widths in excess of t0k The diamonds show the UV star formation rates (uncorrected for extinction) from the ground based work of Wilson et al. (51) !Uld the triangles the Galex results of Wyder et al. (52) !Uld Schiminovich et al. (41) . He> measurements from the literature as summarized in Ly et al . (27) are shown with filled circles. In all cases the errors are ±la.
113
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 a.Iso
show the star formation rates from Ha selected samples reported in the literature and
summarized in Ly et al. (27). In order to better understand the evolution we have a.Iso
restricted our own sample to objects with rest frame equivalent widths of Ha in excess of
200A at low redshifts and HP equivalent widths in excess of 70A 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 (l+z)3 broadly similar
to other UV and optica.lly 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.
4.6 Galaxy Metallicities
4.6.1 [DIll] emitters
The spectra are of variable quality and, in order to measure the meta.llicities, we need very
high signal to noise observations. It is a.Iso important that Balmer lines are well detected
since our flux ca.librations are done using the neighboring Balmer lines (§4.4). We therefore
restricted ourselves to [OIIl] emitters whose HP fluxes are detected above 15 sigma.. Among
92 [OIlI] emitters in our total spectroscopic sample, 8 such [OIIl] emitters were chosen
in the NB912 sample, and 10 in the NB816. These emitters have H-y detected above 4
sigma.. Tables 4.5 and 4.6 give the oxygen line fluxes normalized by their HP fluxes for the
NB816 and NB912 selected emitters, respectively. 1 sigma. upper limits are listed when the
measured flux is below 1 sigma..
The most direct method to estimate the gas-phase oxygen abundance is to use the
electron temperature of the HII region. Higher gas meta.llicity increases nebular cooling,
leading to lower electron temperatures. Therefore electron temperature is a good indicator
114
of the gas metaJlicity. The electron temperature can be derived from the ratio of the
[OIII).>.4363 auroral line to [OIII)'u5007,4959. This procedure is often referred to as the
'direct' method or Te method (e.g., 43; 35; 36; 16). One well-known problem with the direct
method, however, is that [OIII)A4363 is generally very weak even in the low-metaJlicity
galaxies. For higher metaJlicity systems, the far-ffi lines become the dominant coolant and
therefore the optical auroral line is essentially not detectable. However, the majority of our
sample exhibit [OIII)A4363, already suggesting that these are metal-deficient systems. To
derive Te[OIlI] and the oxygen abundances, we used the Pagel et al. (35) calibrations with
the Te[OIl)-Te[OIII) relations derived by Garnett (7). The results are shown in Table 4.5
(for NB816 selected [OIII) emitters) and Table 4.6 (for NB912 selected [OIII) emitters).
The Izotov et al. (16) formula, which was developed with the latest atomie data and
photoionization models, gives consistent abundances within 0.1 dex. The [SII]AA6717, 6731
lines that are usually used for the determination of the electron number density, are beyond
the KeckjDEIMOS wavelength coverage for our [OIlI) emitters. Therefore we assumed no
= 100 em-3• However the choice of electron density has little effect as electron temperature
is insensitive to the electron density; indeed we get the same results even when we use no
= 1,000 em-3 •
The 100 upper and lower limits on Te[OIIl) and the oxygen abundances are also shown
in the tables. Because the [OIIl)A4363 fluxes are weak, the range of our metallicity
measurements are quite wide. However, of 18 [OIlI] emitters, even the upper metaJlicity
limits on 6 emitters satisfy the definition of XMPGs [12 + log(OjH) < 7.65; Kunth &
Ostlin 24]. All our emitters, except the ones that only have lower limits on metallicities
such as NB816 selected USEL 31 in Table 4.5, have very low metallicities - even the upper
metaJlicity limits are about 0.02 - 0.2 Z0' A few emitters may even have metaJlicities
that are comparable to the currently known most metal-poor galaxies ~ Zw 18 and
SBS0335-052W; 12 + log(OjH) ~ 7.1 - 7.2]. However our current metallicity errors
are too large to measure the baseline metaJlicity accurately and higher S jN spectra will be
Decessary for this purpose.
115
Given the weakness of the [OIII].M363 lines even in the detected systems and the
presence of galaxies without detected [OIII).M363, it is useful to also consider other estimates
of the metallicity. For most of the galaxies the line coverage is such that only the R23
indicator of Pagel et aI. (35) is useful. While there has been considerable controversy about
the R23 calibration, recent analyses of large samples of local galaxies with well determined
abundances from the direct method have shown a good empirical correlation between R23
and 12+log(O/H) at low metallicity (31i 53). Izotov & Thuan (17) have tested this and
found that the Yin et aI. formula
12 + log(O/H) = 6.486 + 1.401 x log(R23)
provided the best match to their sample. We have included the empirically determined
abundances from the Yin et aI. formula as the final column in Tables 4 and 5. The
uncertainties are computed using the statistical error for the much stronger [OIII]>'5007
line and a 10% systematic error for the [OIII]>'4959 line. The [011]>'3727 line is very
weak and the uncertainty in this line has little effect on the errors. We compare the
abundances from the two methods in Figure 11. Most of the direct and empirical abundance
determinations agree within the 1 sigma errors but in general the empirical abundance, as
might be expected, favors the upper end of the direct abundance range for the galaxies with
the lowest abundance determinations from the direct method. The empirical method places
seven galaxies in the XMPG category.
Our discovery rate of XMPGs appears to be significantly higher than those of other
surveys. Only 14 new XMPGs have been discovered from the analysis of ~530,OOO galaxy
spectra in the SDSS and they are all located nearby (z < 0.005) (SDSS DR3: Kniazev
et al. 21, SDSS DR4: Izotov et aI. 14). At higher redshift, 17 metal-poor (7.8 < 12 +
log(O/H)< 8.3) galaxies with [OIII)>'4363 detection have been found at z ~ 0.7 in the
initial phase of the DEEP2 survey of 3,900 galaxies and the Team Keck Redshift Survey of
1,536 galaxies (9). But none of these galaxies satisfies the condition of XMPGs. Therefore
116
....... ~ 150
~ ~ 100 + " 8 ~ ~ ......
50
1.0
o o o o ~
1.25~ 6
t 1.5 "! • 2.0
3.0
0.001 0.010 0.100 1.000 [011]3727 / [0111]5007
Figure 4.15 [OIII]>.4959+>'5007/[OIII]>'4363 versus [OII]>.3727/[OIII]>.5007 for the [OIII] and Ha selected emitters in Table 4.5 and 4.6. The electron temperature of the illI region is also shown.
we conclude that the narrowband method is very powerful for finding not only high-redshift
(z > > 5) galaxies, but also strong emission-line, extremely metal-deficient galaxies at
intermediate redshifts (z < I).
Figure 4.15 shows the electron temperature sensitive line ratio,
[OIII](>'4959+>'5007}/[OIII]>'4363 versus [01l]>.3727/[OIII]>.5007. H we have an estimate
of the metallicity, as in the present case, we can use the [011]>.3727/[0111]>.5007 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 density (19). For the low metallicity systems with strong [0111]>'4363 lines, we
can see from Figure 4.15 that [011]>.3727 is very weak compared to [OIII]>'5007 with values
ranging downwards from 0.3. Assuming the metallicity is less than 0.2 Z0 this would place a
lower bound of q = lOS em s-1 on the ionization parameter based on the Kewley & Dopita
(19) model. The higher metallicity objects have stronger [OII]>'3727 /[0III]>'5007 which,
117
while in part due to the metallicity, also requires these objects to have lower ionization
parameters suggesting we are seeing an evolutionary sequence.
4.6.2 Ha emitters
Among 13 Ha emitters in our spectroscopic sample, only three NB912 selected emitters have
H,8 fluxes adequate (> 15o') for the purpose of metallicity measurements. Their Te[OIIIj
and oxygen abundances were measured using the direct method described above, and are
shown in the Table 4.6 together with the data for the [OIIIj emitters. The [0II]A3727line of
NB912 selected USEL 266 is outside the Keck/DEIMOS wavelength coverage. In order to
derive an upper-limit on the metallicity, we l1B8umed [OIIjA3727/[OIII]A5007 = 1.0, which
is the maximum value in our sample (see, Fig. 4.15). All our Ha emitters are metal poor
(Zupper < 0.45 Z0), but none of them are XMPGs.
4.6.3 Composite Spectrum
AI; can be seen in Figure 4.15 the objects with low [0II]A3727/[OIIIjA5007 have relatively
uniform values of ([OIIIjA5007+A4959}/[OIIIjA4363 and similar metallicities. 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 weightings on the lines which depend on
the individual ionization parameters and metallicity but will give a rough estimate of the
average metallicity and temperature of the sample.
In Figure 4.16 we show the composite spectrum of the 8 objects with
[OIIjA3727 /[OIIIjA5007less than 0.1. The [OIIIjA4363 is now strongly detected with a value
of 16.7 ± 2.1 or eight sigma. The mean temperature is 19,500 ± 1,500 K and the average
abundance 12+log(O/H)=7.5±0.1 and the mean rest frame equivalent width ofH,8 is 57A.
The resnlts are similar to the values obtained by averaging the individual analyses of the
eight objects.
118
>< :> ti
-:
-: J I •. ..I J .J. .LO.
.. §: i % ~
=' <0 ... Q. :J: '"
i i i II- I I ,
.. i i i:
3600 3800 4000 4200 4400 4600 4800 5000 5200 REST WAVELENGTH (A)
Figure 4.16 Composite spectrum of the 8 emitters with OII],\3727/0III,\5007 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.
119
4.7 Morphologies
The morphology of the USELs may give us a clue to the mechanism of their high star
formation activity (SFR ~ a few M0/yr) and star formation history; what has triggered
the star formation - mergers, gas infaJl, or galactic winds? High resolution HST / ACS
images are available for the GOODS-North (GOODS-N) region (8) which is one of our
survey fields. There are 17 NB816 selected USELs in the GOODS-N, and 16 in the NB912
sample. Figures 4.17 and 4.18 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 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, 1-
band for NB912 emitters. High-redshift Lya emitters (z > > 5) are very red because of the
continuum absorption below Lya 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, though
USEL 118 in the NB816 sample is an XMPG based on both the direct and the empirical
methods. However, 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.
4.8 Discussion
The present emitters differ from the local dwarf IDI 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 [OIII]/[OII] ratios, and they are a relatively high
fraction (about 10% by number from Figure 4.14) of other 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
120
Figure 4.17 HST/ ACS (B, V, z') composite images of NB816 emitters in the GOODS-N fi eld with overlaid object IDs from Table 4.3 and redshifts, where known. Fields are 12"5 on a side.
121
Figure 4.18 HST / ACS (8 , V , I) composite images of 8912 emitters in the GOODS-N field with overlaid object IDs from Table 4.4 and redshifts, where known. Fields are 12':5 on a side.
122
9
,-..
t It ~~ ; \\ I 8 ""-0 ........ 0' r 0 IZw18
+ • ~BS0335 N 7 ~
• Our sample
• locol XMPGs
... z near 1 Low-Z
6 o GR8 Hosts
-12 -14 -16 - 18 -20 -22 M [AS]
B
Figure 4.19 The oxygen abundance versus the absolute rest frame B magnitude for the [alII] selected samples (red squa,-es). One sigma errors are showll for the oxygen abu ndances and one sigma lower limits are shown with upward pointing arrows. The solid line shows the (44) relation for the nearby dwarf irregulars. As with the local XMPGs (filled circles, Kniazev et al. 21 ; Kewley et al. 20 and GRB hosts (open squares, Stanek et al. 45; Kewley et aI. 20, 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 a l. 9. are shown as t riangles. 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, TIman & Izotov 48) and SBS 0335-052W (12 + log O/ H = 7.12 ± 0.03, Izotov et al. 13).
123
,..... I "-0 '-" 0" 0
+ N ~
. t t~111 8
" .-o.J ;' . •• • IZw18 • • •
7 ~B50335
• Our sample
• local XMPGs
& z near 1 low- Z
o GRB Hosts 6~ __ ~ ____ ~ __ ~~ __ ~~~~
-12 -14 - 16 -18 Me [AB]
- 20 -22
Figure 4.20 The oxygen ahundance based on the empirical method versus the absolute rest frame B magnitude for the [OllIj selected samples (r-ed squares). One sigma errors are shown for the oxygen abundances. The remaining parts of the figure follow Figure 20.
standard Salpeter [MF a forming galaxy can have equi valent widths above 30A for 109 yr
(25) wh ich would allows us to understand the observed number density of strong emitters
relative to the total galaxy population.
Tn this type of model we would expect the metallicity to !,'TOW with t ime and that higher
metalJici ty galaxies would have higher continuum magnitudes and lower equivalent widths
in H,B. We plot the absolu te rest frame B magnitudes versus the oxygen abundance deri ved
by the direct method in Figure 4.19. As with the case for the loca.! XMPGs found in the
SDSS (filled circles, Kniazev et al. 21; Kewley et al. 20) and the metal-poor galaxies
(7.8 < 12 + log Oj H < 8.3) at z ~ 1 (triangles, Hoyos et al . 9), the present emitters (red
squares) are much more luminous at a given metall icity than is found for the local dlrrs
(solid line, Skillman et al. 44) . F\lfthermore there does indeed seem to be a trend to higher
continuum luminosities at higher metallicity consistent with ongoing star formation raising
the luminosity. Recently Kewley et al. (20) reported the similarity between XMPGs and
124
long duration ORB hosts; they share similar SFRs, extinction levels, and both lie in a similar
region of the luminosity-metallicity diagram. Our sample metal-deficient galaxies, which
also lie in the region of ORB hosts, may be additional support of the connection between
XMPOs and ORB hosts. Figure 4.20 shows the same relation but the oxygen abundance
of our sample is based on Yin et al. empirical method. Essentially the same trend as
Figure 4.19 can be seen but now with much smaller error bars in oxygen abundance.
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 primarily have 12+ log( 0 /H) in the range 7.1 - 7.8. If we
assumed that the metallicity were a simple linear function 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. 25) suggesting that the
enrichment 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 H.8 equivalent width is shown in
Figure 4.21. 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 H.8 equivalent 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 metallicity is
consistent with that expected in a simple growth model. Perhaps even more importantly
as larger spectroscopic samples are obtained we should be able to determine if there is a
floor on the metallicity corresponding to the enrichment in the intergalactic gas. Within
the errors 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.
125
9~------------~------~~--~
8 - .
~ 7 :··········t·· .. ~ ....................... : 6~ __________ ~~ ________ ~ __ ~~
10 100 EW(HBeta)
1000
Figure 4.21 The oxygen abundance versus the rest frame H/1 equivalent width for the [Om] 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 I Zw 18.
126
Table 4.2. Narrowband Survey Area Coverage
Field RA (J2000) Dec (J2000) (11I,bll) EB_VIt. NB816 NB912 (arcmin') (arcmin')
SSA22 22:17:57.00 +00:14:54.5 ( 63.1,-44.1) 0.07 674 591 SSA22-'lew 22: 18:24.67 +00:36:53.4 ( 63.6,-43.9) 0.00 278 278 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 no 528
Total 1940 1675
Note. - An adjacent field to A370.new is a site of a gravitational lensing cluster at z - 0.375, and WIlB omitted from the suvey •
• estimated using http://irsa.ipac.caitech.edu/applications/DUST/ based on Schlegel et aI. (42)
4.9 Summary
We have described the results of spectroscopic observations of a narrowband selected sample
of extreme emission 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 [OIIIj.>.4363line and oxygen abundances can be measured using the direct
method. The abundances primarily lie in the 12+1og(O/H) range of 7-8 characteristic of
XMPGs.
The results suggest that at these redshifts we are seeing early chemical enrichment of
startup galaxies which are forming in relatively chemically pristine regions. As we develop
larger samples of these objects and improve the accuracy of their abundance estimates we
should be able to test this model and determine if there is a floor in the metallicity of the
galaxies.
127
Table 4.3. NB816 selected USEL Sample
No. R.A. (J2000.0) Decl. (J2000.0) N(AB) z'(AB) I R V B z .... (1) (2) (3) (4) (5) (6) (7) (8) (9) (10)
1 40.115623 -1.694778 23.92 25.28 24.84 25.44 -99.00 -99.00 -1.0000 2 40.116665 -1.617472 24.26 25.88 25.52 25.74 -99.00 -99.00 -1.0000 3 40.138332 -1.405694 23.49 24.94 24.70 25.22 -99.00 -99.00 0.6343 4 40.174706 -1.704861 24.27 25.46 25.35 25.82 -99.00 -99.00 0.6355 5 40.183208 -1.495500 24.58 25.28 25.85 26.86 -99.00 -99.00 5.6886 6 40.217041 -1.494889 24.80 26.08 26.31 26.14 -99.00 -99.00 0.2416 7 40.250916 -1.744694 23.51 24.35 24.33 24.54 -99.00 -99.00 -1.0000 8 40.276291 -1.518222 24.72 24.78 25.53 25.75 -99.00 -99.00 -1.0000 9 40.276543 -1.623361 24.55 25.36 25.36 25.76 -99.00 -99.00 -1.0000
10 40.284168 -1.453667 21.82 23.20 23.02 22.56 -99.00 -99.00 0.2480 11 40.298878 -1.447472 24.32 25.95 25.42 25.71 -99.00 -99.00 -1.0000 12 40.304039 -1.391806 20.88 22.35 22.06 22.32 -99.00 -99.00 -1.0000 13 40.307045 -1.638611 24.91 25.67 25.72 26.02 -99.00 -99.00 -1.0000 14 40.311249 -1.535195 24.08 25.93 25.49 25.71 -99.00 -99.00 0.6240 15 40.318542 -1.548278 24.60 25.03 25.50 25.96 -99.00 -99.00 -1.0000 16 40.318832 -1.430972 23.37 23.77 24.32 24.42 -99.00 -99.00 1.1804 17 40.319122 -1.446417 23.60 23.85 24.43 24.56 -99.00 -99.00 1.1873 18 40.320793 -1.778083 20.70 21.57 21.60 21.86 -99.00 -99.00 -1.0000 19 40.324123 -1.410056 24.29 24.38 25.18 25.45 -99.00 -99.00 -1.0000 20 40.326042 -1.709833 23.24 24.75 24.62 24.96 -99.00 -99.00 -1.0000 21 40.336460 -1.570278 24.26 24.73 25.06 25.33 -99.00 -99.00 -1.0000 22 40.337292 -1.388278 24.81 27.37 26.71 26.77 -99.00 -99.00 -1.0000 23 40.337582 -1.658389 24.89 25.20 25.76 26.05 -99.00 -99.00 -1.0000 24 40.340378 -1.689500 24.55 26.37 26.07 26.51 -99.00 -99.00 -1.0000 25 40.340706 -1.551944 24.99 25.70 25.93 26.35 -99.00 -99.00 -1.0000 26 40.340874 -1.516333 24.89 25.06 25.69 25.78 -99.00 -99.00 -1.0000 27 40.340958 -1.371306 22.42 23.43 23.29 23.16 -99.00 -99.00 -1.0000 28 40.341331 -1.493583 23.37 24.62 24.55 24.23 -99.00 -99.00 -1.0000 29 40.341457 -1.484222 24.48 25.46 25.49 25.73 -99.00 -99.00 -1.0000 30 40.342918 -1.599611 24.65 25.31 25.50 25.85 -99.00 -99.00 -1.0000 31 40.343082 -1.438944 23.17 24.60 24.54 24.54 -99.00 -99.00 0.6226 32 40.347584 -1.403944 24.86 27.45 26.16 26.35 -99.00 -99.00 0.6324 33 40.349873 -1.598555 23.11 23.87 24.11 24.46 -99.00 -99.00 1.1956 34 40.356415 -1.515805 24.94 26.19 26.18 26.42 -99.00 -99.00 -1.0000 35 40.372250 -1.390472 23.85 24.68 24.72 24.62 -99.00 -99.00 -1.0000 36 40.373665 -1.722583 24.93 25.85 25.74 25.88 -99.00 -99.00 -1.0000 37 40.377918 -1.701972 23.45 23.82 24.26 24.70 -99.00 -99.00 -1.0000 38 40.388081 -1.697472 23.96 24.32 24.79 24.90 -99.00 -99.00 -1.0000 39 40.388916 -1.573472 24.79 25.13 25.65 26.07 -99.00 -99.00 -1.0000 40 40.394585 -1.521500 22.73 23.69 23.84 24.06 -99.00 -99.00 0.6291 41 40.395916 -1.374472 24.04 24.55 25.07 25.06 -99.00 -99.00 1.1811 42 40.395958 -1.507833 24.27 24.79 25.14 25.32 -99.00 -99.00 -1.0000 43 40.409168 -1.763222 24.34 24.86 25.22 25.23 -99.00 -99.00 -1.0000 44 40.416420 -1.728778 24.74 25.51 25.83 26.01 -99.00 -99.00 0.6192 45 40.437538 -1.651611 24.51 25.07 25.35 25.50 -99.00 -99.00 -1.0000
128
Table 4.3--Continued
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)
46 40.439831 -1.731944 24.77 25.38 25.70 25.62 -99.00 -99.00 -1.0000 47 40.441750 -1.466444 24.99 25.69 25.81 26.14 -99.00 -99.00 -1.0000 48 40.448750 -1.715833 23.72 25.28 24.83 25.10 -99.00 -99.00 -1.0000 49 40.450581 -1.457389 22.88 23.80 23.74 23.93 -99.00 -99.00 -1.0000 50 40.466961 -1.584972 24.38 25.40 25.20 25.32 -99.00 -99.00 -1.0000 51 40.472122 -1.524972 23.51 24.77 24.47 24.66 -99.00 -99.00 0.6359 52 40.480789 -1.666417 24.81 25.22 25.61 25.73 -99.00 -99.00 -1.0000 53 40.481750 -1.620750 23.80 24.56 24.78 25.25 -99.00 -99.00 -1.0000 54 40.484291 -1.714278 24.87 25.99 25.86 26.28 -99.00 -99.00 -1.0000 55 40.494457 -1.453563 24.97 25.57 25.94 26.06 -99.00 -99.00 -1.0000 56 40.503288 -1.533778 24.61 26.95 26.27 26.72 -99.00 -99.00 -1.0000 57 40.505081 -1.376083 22.05 23.38 22.91 23.28 -99.00 -99.00 -1.0000 58 40.514919 -1.504611 23.81 24.65 24.94 25.26 -99.00 -99.00 0.6238 59 40.520126 -1.523139 23.66 24.58 24.64 24.69 -99.00 -99.00 -2.0000 60 188.733032 62.057167 22.31 23.56 23.31 23.53 23.89 24.00 -1.0000 61 188.732742 62.349419 23.29 24.38 24.14 24.41 24.79 24.90 -1.0000 62 188.745621 62.164192 21.84 23.18 22.89 23.45 24.05 23.96 -1.0000 63 188.761505 62.410778 24.47 25.57 25.30 25.64 26.48 26.29 -1.0000 64 188.763535 62.287224 24.62 26.76 30.26 27.39 27.70 29.49 5.6985 65 188.769196 62.083973 24.72 24.85 25.69 25.58 27.06 32.17 -1.0000 66 188.778275 62.170082 24.68 25.31 25.58 25.70 26.57 26.11 -1.0000 67 188.828262 62.032722 22.40 23.49 23.32 23.52 24.12 24.28 -1.0000 68 188.868088 62.415138 24.72 25.55 25.53 26.05 26.60 26.10 -1.0000 69 188.873535 62.174362 22.58 24.08 23.75 23.99 24.39 24.44 -1.0000 70 188.885742 62.364082 24.10 26.07 25.22 25.92 26.33 26.40 0.6362 71 188.888214 62.057556 24.84 25.56 25.89 25.76 26.28 26.21 -1.0000 72 188.914429 62.247555 23.19 25.11 24.60 24.89 25.18 25.41 -1.0000 73 188.918243 62.020164 24.15 25.49 25.06 25.00 25.47 25.69 -1.0000 74 188.920532 62.105362 22.54 24.37 23.75 24.22 24.49 24.50 -1.0000 75 188.928421 62.127003 24.82 26.76 26.61 26.51 26.48 26.51 -1.0000 76 188.939423 62.334141 21.64 23.08 22.73 23.08 23.38 99.00 0.6319 77 188.943375 62.339664 23.49 24.66 24.75 24.87 25.35 25.56 -1.0000 78 188.950241 62.020889 23.84 25.73 25.39 26.02 26.80 26.28 -1.0000 79 188.954468 62.326195 24.68 26.01 25.55 25.92 26.58 26.58 -1.0000 80 188.957581 62.323471 23.45 24.71 24.41 24.67 25.12 25.08 -1.0000 81 188.964874 62.320164 23.45 24.82 24.41 24.81 24.97 25.20 0.6360 82 188.980362 62.104973 24.51 25.67 25.41 25.78 26.09 25.96 0.6426 83 188.985718 62.048695 24.23 24.87 25.03 24.89 25.76 25.62 0.6367 84 188.986221 62.341640 23.05 25.10 24.11 25.20 25.43 25.45 0.6352 85 188.991531 62.436169 24.91 24.25 26.03 25.28 24.90 27.58 -1.0000 86 188.992081 62.142998 24.49 25.50 25.47 25.57 26.30 26.18 0.6327 87 189.007919 61.996750 24.78 25.08 25.63 25.55 25.54 25.59 -1.0000 88 189.014206 62.362556 24.65 26.36 26.40 26.47 26.58 28.45 -1.0000 89 189.014877 62.359585 25.00 27.98 26.29 25.58 25.83 26.04 -1.0000 90 189.028870 62.145500 24.97 25.95 26.31 25.59 26.74 26.67 -1.0000
129
Table 4.3-Continued
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)
91 189.033127 61.994053 24.63 27.51 27.84 29.34 30.44 27.06 5.7138 92 189.039078 62.435024 24.65 25.13 25.49 24.95 25.34 25.46 -1.0000 93 189.045166 62.105392 23.80 25.58 25.08 25.37 26.06 25.92 -1.0000 94 189.056168 62.129974 24.88 26.25 26.27 30.24 26.85 29.17 5.6350 95 189.085044 62.195415 22.97 23.84 23.78 23.88 24.31 24.68 0.6297 96 189.090927 62.188831 24.57 24.45 25.47 26.76 28.00 27.57 5.4610 97 189.093658 61.999249 23.07 25.39 24.70 25.29 25.47 25.47 -1.0000 98 189.097824 62.272026 23.54 24.76 24.42 25.00 25.46 25.62 -1.0000 99 189.101456 62.395695 23.98 25.29 25.40 25.81 25.98 25.99 0.6322
100 189.106201 62.002777 23.64 25.00 24.81 25.11 25.44 25.51 -1.0000 101 189.110336 62.063030 24.96 26.25 26.25 27.16 28.02 27.85 5.7010 102 189.110916 62.245529 22.73 25.58 24.47 25.06 25.72 25.61 -1.0000 103 189.130417 62.038864 24.40 24.56 25.47 25.24 25.34 25.34 1.1952 104 189.138535 62.410442 24.83 24.99 25.68 25.16 27.97 27.37 -1.0000 105 189.176178 62.135834 24.58 25.22 25.57 25.39 26.14 26.38 -1.0000 106 189.190765 62.384472 24.34 25.91 25.98 25.60 25.72 26.05 0.2384 107 189.200897 62.328556 24.34 25.61 25.66 25.28 25.89 26.05 -1.0000 108 189.204910 62.192387 24.88 24.90 25.83 26.70 27.34 27.56 -1.0000 109 189.205292 62.013584 24.97 25.39 25.89 25.24 26.12 27.89 -1.0000 110 189.215210 62.326832 24.63 25.70 27.40 29.00 27.70 29.10 5.6750 111 189.216827 62.364609 24.84 29.07 27.06 27.32 27.78 28.67 5.6873 112 189.226837 62.394279 24.85 25.03 26.19 25.95 26.34 26.17 -1.0000 113 189.291046 62.367527 24.52 24.94 25.37 25.73 27.17 28.28 -1.0000 114 189.294052 62.055916 24.19 24.11 24.99 25.04 25.36 25.62 -1.0000 115 189.294083 62.035973 24.99 25.22 26.37 25.87 25.87 26.00 -1.0000 116 189.305740 62.340336 24.79 24.73 25.61 25.21 25.52 25.62 -1.0000 117 189.305542 62.421364 24.97 25.16 25.79 25.48 25.28 26.03 -1.0000 118 189.310959 62.174469 24.15 27.20 25.06 26.14 26.00 25.97 0.6439 119 189.313705 62.007584 24.79 28.49 25.86 26.45 26.57 26.81 -1.0000 120 189.360123 62.382584 24.51 25.01 25.62 25.37 25.68 25.66 -1.0000 121 189.361160 62.262306 24.26 25.78 25.44 25.76 26.67 26.62 0.0000 122 189.375626 62.392754 22.78 24.44 24.12 24.42 24.92 24.90 -1.0000 123 189.383194 62.286167 24.28 25.13 25.20 25.52 26.92 27.70 0.6342 124 189.383987 62.268448 24.96 25.12 29.03 25.78 26.47 26.10 -1.0000 125 189.383865 62.273277 24.72 25.09 25.55 25.01 26.05 26.11 -1.0000 126 189.387344 62.406555 23.85 26.55 25.59 26.42 27.56 27.19 -1.0000 127 189.390533 62.281364 24.69 26.82 26.21 25.28 25.17 25.06 -1.0000 128 189.391907 62.282108 24.77 29.34 26.91 25.30 26.53 26.36 -1.0000 129 189.392044 62.253914 24.02 29.56 25.53 26.03 26.17 26.58 -1.0000 130 189.408920 62.382164 24.82 25.24 25.64 25.84 25.70 25.93 -1.0000 131 189.436081 62.413029 24.18 24.61 25.29 25.48 25.71 25.78 -1.0000 132 189.436249 62.195915 24.90 27.54 26.96 26.51 27.56 27.35 5.6699 133 189.438965 62.223331 24.59 24.51 25.40 26.46 27.28 28.82 -1.0000 134 189.461365 62.147663 23.51 25.79 24.88 25.33 25.92 26.21 0.0000 135 189.497375 62.344833 24.74 25.28 25.60 25.63 25.99 25.42 -1.0000
130
Table 4.3---Continued
No. R.A. (J2000.0) Decl. (J2000.0) N(AB) z'(AB) I R V B z.".e (1) (2) (3) (4) (5) (6) (7) (8) (9) (10)
136 189.508835 62.273308 24.36 24.79 25.41 25.37 25.75 25.88 -1.0000 137 189.514969 62.243473 22.39 23.57 23.56 23.93 23.93 24.12 1.1802 138 189.517288 62.246944 23.30 23.95 24.15 24.31 24.30 24.45 -1.0000 139 189.518616 62.384193 24.45 24.71 25.33 25.26 25.50 25.52 -1.0000 140 189.533279 62.349331 22.29 22.88 23.10 23.48 23.64 23.72 -1.0000 141 189.537338 62.413666 24.96 27.71 27.55 26.36 29.18 26.44 -1.0000 142 189.552673 62.398167 23.88 24.20 24.62 24.81 24.84 24.86 -1.0000 143 189.566376 62.396510 24.87 25.87 26.50 26.31 25.91 25.88 -1.0000 144 189.570587 62.057972 24.50 24.81 25.36 25.52 26.43 26.00 -1.0000 145 189.577286 62.272663 24.81 26.26 29.31 27.10 28.23 27.21 5.7275 146 189.600052 62.202751 24.94 25.64 25.95 25.24 25.63 26.24 -1.0000 147 189.607834 62.433750 24.72 25.87 25.53 26.00 29.06 27.20 -1.0000 148 189.608383 62.387306 23.96 24.27 24.90 25.08 25.20 25.23 0.0000 149 189.612579 62.207664 24.12 24.25 25.00 25.54 27.38 27.19 0.0000 150 189.629166 62.255390 23.04 23.56 23.87 24.15 24.25 24.29 1.1843 151 189.635117 62.404610 24.94 24.99 25.79 25.79 26.24 26.43 -1.0000 152 189.638306 62.205860 24.76 24.51 25.59 28.85 26.53 27.55 0.0000 153 189.653839 62.249142 23.94 24.23 25.00 25.09 25.10 25.40 -1.0000 154 189.654907 62.099976 24.14 25.59 25.26 25.57 26.36 27.11 0.0000 155 189.677246 62.143169 24.91 24.30 25.73 28.55 27.49 27.64 -1.0000 156 189.680618 62.392223 24.63 25.31 25.46 24.90 25.57 25.80 -1.0000 157 334.191620 0.370306 23.79 25.71 24.79 25.64 25.86 26.22 -1.0000 158 334.195343 0.162833 23.73 23.93 24.58 25.22 25.51 25.63 1.1911 159 334.198517 0.264389 24.99 25.11 25.88 26.42 30.98 25.40 -1.0000 160 334.208710 0.263056 24.92 25.21 25.80 25.42 24.97 24.86 -1.0000 161 334.212219 0.440861 22.47 23.45 23.35 23.83 24.35 24.22 -1.0000 162 334.210144 0.341056 24.95 25.55 25.88 26.14 26.29 26.45 -1.0000 163 334.217438 0.413917 23.85 25.70 24.97 25.69 26.93 25.94 -1.0000 164 334.216248 0.101194 22.52 24.19 23.61 24.43 24.85 24.83 -1.0000 165 334.219299 0.225444 22.67 24.41 23.94 24.65 25.13 24.96 -1.0000 166 334.222290 0.093000 22.97 23.39 23.81 24.46 24.68 24.76 1.1845 167 334.223785 0.307833 24.92 25.04 25.75 26.34 26.23 26.45 0.0000 168 334.225037 0.248528 24.35 25.08 25.19 25.23 26.37 26.63 -1.0000 169 334.228882 0.094083 23.93 26.55 26.09 99.00 99.00 27.71 5.6765 170 334.228577 0.454000 22.74 24.46 23.98 24.87 25.43 25.71 -1.0000 171 334.233612 0.411389 24.67 24.82 25.54 26.02 25.87 25.88 -1.0000 172 334.246246 0.387722 23.37 25.51 24.65 25.50 25.84 25.73 0.6177 173 334.265045 0.086528 24.47 25.56 25.57 26.34 26.23 26.67 -1.0000 174 334.266937 0.086083 24.98 25.28 25.85 26.64 25.68 26.44 0.0000 175 334.273193 0.216889 24.41 25.85 26.14 99.00 99.00 29.49 5.6703 176 334.275391 0.397806 22.32 23.07 23.36 24.17 24.60 24.63 -1.0000 177 334.283051 0.060000 24.92 25.46 25.76 26.04 27.99 25.84 -1.0000 178 334.290039 0.054444 24.04 24.85 24.84 25.10 99.00 25.29 0.6210 179 334.302704 0.252750 24.32 24.76 25.13 25.64 25.77 25.58 0.0000 180 334.336914 0.335333 24.55 27.03 26.26 27.85 27.07 27.21 5.6706
131
Table 4.3-Continued
No. R.A. (J2000.0) Decl. (J2000.0) N(AB) z'(AB) 1 R V B z.".., (1) (2) (3) (4) (5) (6) (7) (8) (9) (10)
181 334.337219 0.293806 24.10 24.83 25.65 99.00 99.00 99.00 5.6672 182 334.343658 0.078083 23.06 24.24 24.03 24.55 24.88 24.77 -1.0000 183 334.346863 0.317306 24.79 25.21 25.61 25.96 25.85 25.70 1.1919 184 334.346741 0.316694 22.57 23.12 23.38 23.92 24.11 23.88 1.1919 185 334.351868 0.197222 24.72 27.39 27.92 28.47 26.94 99.00 0.0000 186 334.359161 0.253528 24.04 26.85 25.47 26.73 29.17 27.79 0.6167 187 334.360413 0.466722 24.81 25.32 25.94 26.50 28.69 27.92 -1.0000 188 334.369751 0.321694 24.29 25.08 25.29 27.11 28.39 28.26 5.6416 189 334.373138 0.344361 24.98 99.00 27.50 26.29 27.72 27.87 -1.0000 190 334.373199 0.300861 23.93 24.26 24.75 25.23 25.59 25.10 1.1866 191 334.381714 0.160278 24.59 25.99 26.55 28.61 99.00 27.96 5.6840 192 334.383118 0.344750 24.94 26.88 27.24 27.20 27.18 99.00 -1.0000 193 334.387909 0.371278 24.36 26.17 25.95 28.08 26.66 99.00 5.6531 194 334.390808 0.450500 24.19 25.74 25.40 26.03 26.52 25.76 -1.0000 195 334.404327 0.377528 22.59 23.77 23.49 24.05 24.55 24.57 0.6286 196 334.416016 0.307361 24.86 26.68 26.49 28.45 28.16 26.76 5.7240 197 334.420380 0.404250 24.93 25.77 25.76 28.87 27.81 99.00 5.6330 198 334.418121 0.667972 23.43 25.00 24.85 25.20 25.74 25.87 0.6267 199 334.426331 0.128833 24.44 24.96 26.15 99.00 28.33 99.00 0.0000 200 334.428894 0.225694 22.87 25.09 24.39 24.96 25.44 25.02 0.6270 201 334.430969 0.219639 24.73 26.16 25.74 26.05 27.01 26.83 0.6269 202 334.433289 0.225944 22.91 23.82 23.73 24.14 24.64 24.24 0.6278 203 334.439117 0.471750 23.06 24.29 23.86 24.41 24.95 24.69 0.6175 204 334.441345 0.053583 24.41 24.48 25.23 25.38 25.27 25.43 1.1832 205 334.444305 0.374917 23.15 24.49 24.24 24.79 25.38 25.31 -1.0000 206 334.443848 0.657750 23.00 27.31 25.07 26.02 26.59 26.20 0.6276 207 334.445251 0.237194 24.39 26.76 26.02 26.36 26.33 26.31 -1.0000 208 334.449829 0.253722 22.59 25.66 24.49 25.62 26.36 25.81 0.6263 209 334.457672 0.396028 23.27 25.17 24.61 25.53 25.83 25.74 -1.0000 210 334.474579 0.741528 24.82 25.06 25.79 25.55 25.51 25.49 -1.0000 211 334.479889 0.485722 24.43 25.52 25.95 27.06 101.9 31.94 5.7265 212 334.484650 0.273250 23.73 26.26 25.06 26.41 26.59 26.44 0.6247 213 334.490366 0.825528 24.59 26.02 25.06 26.09 26.37 26.06 0.6406 214 334.502625 0.340917 24.30 25.68 25.20 25.53 26.14 26.44 -1.0000 215 334.509094 0.242111 24.12 25.92 25.99 99.00 99.00 99.00 5.6738 216 334.513123 0.088639 21.34 22.68 22.43 22.40 23.16 23.00 -1.0000 217 334.515167 0.367639 24.24 26.31 25.40 25.97 25.98 26.28 -1.0000 218 334.518890 0.239472 24.81 27.43 28.70 99.00 99.00 99.00 5.6807 219 334.533783 0.676139 24.60 25.01 25.49 25.75 25.66 25.58 -1.0000 220 334.535431 0.110722 24.96 25.93 26.53 26.96 27.67 26.77 -1.0000 221 334.543335 0.117583 23.98 25.92 25.58 26.53 27.14 27.22 0.6216 222 334.543884 0.254066 24.62 25.41 25.55 26.13 26.73 26.52 -1.0000 223 334.546783 0.183778 23.16 24.63 24.36 24.92 25.69 25.14 0.6288 224 334.550323 0.420528 24.01 24.66 24.85 25.27 25.29 24.84 1.1865 225 334.554443 0.442000 24.94 25.50 25.86 26.00 26.11 25.38 -1.0000
132
Table 4.3--Continued
No. R.A. (J2000.0) Dec!. (J2000.0) N(AB) z'(AB) I R V B z..." (1) (2) (3) (4) (5) (6) (7) (8) (9) (10)
226 334.555389 0.198833 24.29 27.20 25.94 26.66 26.98 27.71 -1.0000 227 334.556976 0.454917 23.59 24.05 24.46 24.95 24.93 24.79 1.1877 228 334.557159 0.106000 24.94 25.78 25.78 26.31 26.23 26.51 -1.0000 229 334.562775 0.664444 23.68 24.09 24.51 24.57 24.68 26.65 -1.0000 230 334.565613 0.105222 24.01 26.53 24.89 25.67 25.69 25.86 -1.0000 231 334.569214 0.432139 22.62 23.21 23.44 24.01 24.37 24.15 -1.0000 232 334.576172 0.398639 23.23 24.10 24.35 25.21 25.45 25.68 1.1658 233 334.577118 0.454083 24.77 25.61 25.99 26.23 26.93 26.46 -1.0000 234 334.581390 0.359000 24.43 24.74 25.26 25.75 26.29 25.86 -1.0000 235 334.582886 0.130917 23.25 24.48 24.15 24.94 25.40 25.14 -1.0000 236 334.589722 0.794278 24.88 25.14 25.68 25.83 26.35 26.68 1.1731 237 334.593781 0.468806 24.69 25.66 26.06 26.29 25.95 26.59 -1.0000 238 334.597961 0.159111 23.73 25.87 25.05 25.24 26.46 26.07 -1.0000 239 334.597015 0.422083 23.89 25.69 25.11 25.81 26.15 26.02 -1.0000 240 334.602356 0.346472 24;96 25.65 25.95 27.89 99.00 28.00 -1.0000 241 334.606415 0.159028 23.68 25.80 25.08 25.70 27.01 26.85 -1.0000 242 334.607574 0.728944 24.63 26.30 25.73 25.54 25.98 26.06 0.0000 243 334.636139 0.479083 24.32 25.90 25.83 27.58 29.70 27.44 5.6820 244 334.641937 0.466944 24.52 25.39 25.48 25.66 26.18 26.37 0.6255 245 334.641449 0.610944 24.90 25.17 26.17 28.82 101.9 101.63 0.0000 246 334.654999 0.586889 24.54 25.82 25.49 25.22 25.82 25.68 -1.0000 247 334.679749 0.744222 24.80 26.48 26.57 102.18 101.99 101.63 5.6560 248 334.682465 0.455139 24.73 25.32 25.71 25.36 26.07 26.89 -1.0000 249 334.689270 0.717278 23.36 24.60 24.51 24.62 25.29 25.47 -1.0000 250 334.589270 0.717278 22.90 23.98 23.80 23.91 24.56 24.56 -1.0000 251 334.692444 0.556167 24.66 24.47 25.62 25.66 26.00 25.84 -1.0000 252 334.714386 0.711056 22.47 23.86 23.36 23.67 24.11 24.15 0.6414 253 334.715454 0.731028 24.20 26.01 25.33 26.10 26.68 26.57 0.6413 254 334.723999 0.540750 23.02 25.18 24.51 24.98 25.39 25.73 0.6230 255 334.727081 0.528639 24.10 25.23 25.00 25.32 25.83 26.10 -1.0000 256 334.727264 0.455472 24.36 25.18 25.51 26.28 26.39 26.20 0.0000 257 334.735687 0.561056 22.04 21.09 23.38 22.16 21.87 22.49 -2.0000 258 334.759155 0.586722 23.18 23.11 24.14 23.63 24.25 25.11 0.6453 259 334.770844 0.514667 24.88 101.71 27.94 31.27 101.99 30.28 5.6960 260 334.776031 0.517750 24.89 24.82 25.74 26.87 27.96 25.66 -1.0000 261 334.778687 0.517583 24.45 24.28 25.25 25.43 25.89 25.34 -1.0000 262 334.788330 0.628972 24.97 25.57 25.91 25.88 101.9 26.37 -1.0000 263 334.788818 0.769639 23.16 24.04 24.05 24.08 24.46 24.73 0.2430 264 334.788116 0.589528 21.57 21.04 22.80 21.68 21.64 21.10 -1.0000 266 334.792511 0.593722 23.97 26.50 25.02 25.87 25.71 25.92 -1.0000 266 334.793549 0.602611 24.49 25.26 25.32 25.64 25.71 26.30 -1.0000 267 334.796763 0.489083 24.70 25.26 25.55 25.74 26.00 26.12 -1.0000
Note. - MagoItud ... are measured in 3" diameter apertures. Ao entry of '-99' indicates that no excess flux was measured, '99' means that there was a negative flus in the aperture. -1.0000 in the redehIft column means no spectroscopic data were obtained for the object, 0.0000 means that spectra were obtained, but the object was not identified, while -2.0000 denotee a star.
133
Table 4.4. NB912 selected USEL Sample
No. R.A. (J2000.0) Decl. (J2000.0) N(AB) z'(AB) I R V B z.".. (1) (2) (3) (4) (5) (6) (7) (8) (9) (10)
1 40.131710 -1.408361 23.86 25.08 25.89 25.97 -99.00 -99.00 0.8371 2 40.133957 -1.575250 24.76 25.92 26.13 25.84 -99.00 -99.00 1.4498 3 40.148125 -1.593611 23.21 24.62 25.53 26.11 -99.00 -99.00 0.8206 4 40.148376 -1.725417 24.76 25.97 26.98 26.39 -99.00 -99.00 0.8111 5 40.150795 -1.737583 23.31 24.35 24.88 25.32 -99.00 -99.00 0.8269 6 40.153290 -1.538833 23.64 25.00 26.18 26.87 -99.00 -99.00 0.8301 7 40.156544 -1.765833 21.93 23.05 23.51 23.74 -99.00 -99.00 -1.0000 8 40.165752 -1.580139 24.94 26.20 25.92 25.84 -99.00 -99.00 1.4482 9 40.183250 -1.389611 21.77 22.79 23.33 23.52 -99.00 -99.00 0.8325
10 40.184338 -1.596472 23.09 24.50 25.30 25.55 -99.00 -99.00 0.8291 11 40.193504 -1.690111 24.32 25.40 25.83 25.69 -99.00 -99.00 0.8266 12 40.194084 -1.373750 23.93 24.93 25.21 25.35 -99.00 -99.00 -1.0000 13 40.194710 -1.373917 24.87 25.90 26.30 26.47 -99.00 -99.00 -1.0000 14 40.196709 -1.378389 24.07 25.42 26.05 26.18 -99.00 -99.00 -1.0000 15 40.202122 -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 40.214748 -1.519944 23.14 24.40 25.38 25.94 -99.00 -99.00 0.8289 18 40.220333 -1.753778 24.34 25.41 26.40 26.29 -99.00 -99.00 -1.0000 19 40.220997 -1.388611 23.24 24.36 25.13 24.99 -99.00 -99.00 -1.0000 20 40.226917 -1.542139 23.02 24.47 25.77 25.48 -99.00 -99.00 0.8209 21 40.229126 -1.720944 24.99 27.85 27.11 102.70 -99.00 -99.00 6.4800 22 40.229332 -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 40.280624 -1.421111 24.82 25.86 26.13 25.91 -99.00 -99.00 -1.0000 25 40.290958 -1.746389 23.89 25.16 25.81 25.03 -99.00 -99.00 0.0000 26 40.323959 -1.697695 24.73 25.80 26.47 25.71 -99.00 -99.00 -1.0000 27 40.330791 -1.612417 23.03 24.68 26.06 25.47 -99.00 -99.00 0.0000 28 40.339832 -1.395361 23.87 25.54 26.98 25.45 -99.00 -99.00 0.3887 29 40.346794 -1.448389 23.93 24.98 25.37 25.48 -99.00 -99.00 0.8274 30 40.382538 -1.554111 23.52 24.87 25.60 25.08 -99.00 -99.00 0.3930 31 40.393124 -1.713750 24.94 26.83 26.70 26.50 -99.00 -99.00 -1.0000 32 40.398834 -1.466583 23.87 24.91 24.83 24.64 -99.00 -99.00 1.4590 33 40.403870 -1.530195 24.88 26.08 27.44 27.30 -99.00 -99.00 0.8223 34 40.409500 -1.369278 21.52 23.44 24.47 24.04 -99.00 -99.00 -1.0000 35 40.411705 -1.691445 24.41 25.50 25.68 25.55 -99.00 -99.00 -1.0000 36 40.425083 -1.454111 21.91 22.93 23.42 23.41 -99.00 -99.00 -1.0000 37 40.429913 -1.501139 23.13 24.21 24.76 24.92 -99.00 -99.00 0.8267 38 40.446251 -1.676111 24.91 26.09 26.37 26.04 -99.00 -99.00 -1.0000 39 40.479042 -1.534250 24.87 26.06 25.79 25.56 -99.00 -99.00 0.3861 40 40.506042 -1.755111 22.73 24.27 25.24 25.19 -99.00 -99.00 -1.0000 41 40.511585 -1.597000 24.73 25.96 25.55 25.67 -99.00 -99.00 -1.0000 42 40.518124 -1.666167 22.47 23.73 24.37 24.45 -99.00 -99.00 -1.0000 43 188.802414 62.319778 24.91 26.26 26.19 26.10 26.43 22.53 -1.0000 44 188.805588 62.377683 24.47 25.51 25.60 25.66 26.37 22.13 -1.0000 45 188.809082 62.093307 24.71 26.06 25.59 25.18 25.33 21.60 -1.0000
134
Table 4.4-Continued
No. R.A. (J2000.0) Decl. (J2000.0) N(AB) z'(AB) I R V B zSp!!c
(1) (2) (3) (4) (5) (6) (7) (8) (9) (lO)
46 188.817169 62.153473 24.91 26.46 27.75 28.42 28.59 23.85 -1.0000 47 188.833542 62.108307 25.00 26.27 25.08 25.84 25.92 22.33 -1.0000 48 188.837997 62.134029 24.90 26.87 27.82 27.49 28.70 24.06 -1.0000 49 188.841248 62.123889 24.69 25.81 25.92 26.62 26.48 23.59 -1.0000 50 188.856094 62.084747 22.73 24.35 25.13 25.14 25.56 21.75 -1.0000 51 188.892212 62.181778 24.42 25.43 25.47 25.02 25.46 21.41 -1.0000 52 188.919495 62.299862 21.33 22.42 22.78 22.36 22.68 19.17 0.3960 53 188.919373 62.255695 24.93 26.30 26.48 26.25 26.55 22.17 -1.0000 54 188.936844 62.081692 24.89 26.44 27.05 27.14 27.74 23.53 -1.0000 55 188.939911 62.118752 24.97 26.50 25.54 26.18 26.56 22.60 -1.0000 56 188.949173 62.062248 21.86 23.11 23.70 23.56 23.85 19.91 -1.0000 57 188.955826 62.066334 24.94 26.04 25.83 25.87 26.38 22.69 -1.0000 68 188.972321 62.191444 24.10 25.32 25.41 24.94 25.51 21.91 -1.0000 59 188.978271 62.142639 24.87 26.22 26.43 26.04 26.11 22.09 -1.0000 60 189.024841 62.289722 23.23 24.43 24.72 24.43 25.06 21.51 0.3898 61 189.027618 62.082359 24.45 26.22 26.13 29.43 28.12 22.90 -1.0000 62 189.046036 62.360138 24.90 26.33 25.91 26.63 28.17 23.34 -1.0000 63 189.047668 62.lO1501 24.59 26.06 25.91 25.99 26.15 22.47 -1.0000 64 189.057831 62.172222 24.85 26.29 29.31 25.95 26.99 22.85 -1.0000 65 189.074799 62.157028 24.79 25.82 25.19 25.29 25.70 21.66 -1.0000 66 189.082825 62.139248 24.90 26.07 25.56 25.76 26.25 22.06 -1.0000 67 189.119370 62.362720 23.89 25.03 25.16 25.54 25.49 21.44 -1.0000 68 189.125137 62.034111 24.73 26.52 26.33 27.12 34.42 23.51 0.0000 69 189.154999 62.188332 24.22 25.41 25.74 25.68 25.80 21.65 0.8396 70 189.164536 62.283306 24.70 25.71 25.25 25.91 26.45 22.79 -1.0000 71 189.172470 62.179943 24.97 26.39 26.86 26.65 26.68 22.95 -1.0000 72 189.189438 62.313778 24.81 26.13 26.08 26.32 26.69 24.03 -1.0000 73 189.202087 62.189083 24.65 26.36 29.91 27.17 29.50 24.18 -1.0000 74 189.206757 62.203567 24.71 25.83 25.59 25.32 25.19 21.30 -1.0000 75 189.219162 62.161583 24.98 26.04 25.14 25.05 24.90 21.11 -1.0000 76 189.226044 62.342197 23.92 24.98 25.30 25.83 25.89 22.21 -1.0000 77 189.278076 62.402336 24.59 27.20 27.42 26.50 29.04 26.79 -1.0000 78 189.296036 62.053169 24.38 25.53 26.06 26.50 26.94 22.63 0.0000 79 189.296417 62.052696 23.65 24.86 26.14 25.88 26.17 22.09 0.0000 80 189.304382 62.2686lO 23.50 24.76 25.63 25.46 25.56 21.82 -1.0000 81 189.308090 62.1186lO 22.83 24.03 24.81 24.72 25.01 21.20 -1.0000 82 189.312881 62.284142 24.89 26.61 26.56 27.01 26.91 23.12 -1.0000 83 189.314819 62.364780 24.90 26.42 25.84 25.73 25.70 21.80 -1.0000 84 189.317169 62.388390 24.55 25.70 26.04 25.70 26.06 21.99 -1.0000 85 189.346817 62.067139 22.97 24.30 24.82 24.98 25.04 21.19 0.8280 86 189.354828 62.390472 24.88 26.28 27.87 26.68 27.14 24.22 -1.0000 87 189.387344 62.265694 25.00 26.85 26.02 25.45 25.46 21.90 -1.0000 88 189.391907 62.387859 24.85 26.71 27.44 27.55 26.43 23.56 -1.0000 89 189.400620 62.083500 24.91 26.89 27.89 26.19 26.65 23.79 -1.0000 90 189.401199 62.214943 24.49 25.53 25.83 25.78 25.75 22.08 -1.0000
135
Table 4.4--Continued
No. R.A. (J2000.0) Oecl. (J2000.0) N(AB) z'(AB) I R V B Z8peC
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10)
91 189.404999 62.102139 24.24 25.28 25.49 25.41 25.23 21.28 -1.0000 92 189.421585 62.397167 24.98 26.10 25.63 25.59 25.84 21.68 -1.0000 93 189.422470 62.387444 24.92 25.98 26.37 25.53 25.89 21.58 -1.0000 94 189.431992 62.351807 22.26 23.91 25.00 24.94 24.99 21.09 -1.0000 95 189.434677 62.396168 24.63 26.00 25.58 25.51 25.51 21.38 -1.0000 96 189.440475 62.307304 24.71 26.73 26.62 26.00 26.32 22.07 -1.0000 97 189.450241 62.065529 23.81 24.88 25.92 25.80 25.87 22.30 -1.0000 98 189.478958 62.156944 24.32 25.34 25.30 25.15 25.51 22.61 -1.0000 99 189.481674 62.404053 25.00 26.29 25.60 25.77 26.20 22.34 -1.0000
100 189.522563 62.161583 23.08 24.17 24.62 24.78 24.90 21.10 -1.0000 101 189.533249 62.261414 24.98 26.00 26.05 25.67 25.94 22.94 -1.0000 102 189.535751 62.391472 23.35 25.17 25.81 26.00 26.14 22.00 -1.0000 103 189.544647 62.049469 23.51 25.04 26.55 26.20 26.37 22.43 -1.0000 104 189.568802 62.259029 23.86 25.01 25.51 25.76 25.97 21.90 0.8353 105 189.582916 62.299004 24.07 25.38 25.48 25.64 25.82 21.91 0.8190 106 189.594589 62.139641 24.77 26.28 25.95 26.16 26.02 22.94 -1.0000 107 189.612747 62.258945 24.97 26.08 26.03 26.09 26.18 22.67 -1.0000 108 189.613754 62.330719 24.88 25.96 25.60 25.60 26.71 22.78 -1.0000 109 189.615387 62.313416 24.96 27.51 26.03 26.11 26.42 22.79 -1.0000 lIO 189.617569 62.367527 24.96 26.20 26.42 26.83 26.73 24.10 0.0000 111 334.186279 0.099167 24.48 99.00 26.99 27.29 26.61 26.74 0.0000 112 334.188812 0.211222 24.76 26.32 27.14 26.91 26.70 99.00 -1.0000 113 334.190338 0.092056 24.74 26.00 25.99 26.39 26.22 26.04 0.0000 114 334.192841 0.211722 23.20 24.24 24.85 25.40 25.66 25.83 -1.0000 115 334.197296 0.394167 24.59 25.62 26.79 25.28 25.20 24.92 -1.0000 116 334.206879 0.225333 24.75 25.78 26.32 27.23 24.60 99.00 -1.0000 117 334.206909 0.328528 24.54 25.72 99.00 99.00 99.00 99.00 0.0000 118 334.207733 0.425222 23.12 24.66 25.36 25.76 26.16 26.02 -1.0000 119 334.215912 0.082083 24.82 99.00 28.08 29.25 27.52 27.93 0.0000 120 334.220551 0.282694 23.91 25.58 26.22 26.24 26.70 27.35 -1.0000 121 334.233063 0.217056 22.13 23.28 23.78 24.12 24.47 24.18 -1.0000 122 334.233368 0.216306 23.22 25.20 25.77 26.03 25.98 25.71 0.8217 123 334.242035 0.108611 24.72 25.77 26.49 27.17 26.85 27.39 0.8196 124 334.241577 0.295056 24.48 25.72 25.38 25.90 28.67 27.35 0.0000 125 334.240631 0.321000 24.91 26.73 27.17 29.56 30.64 27.16 -1.0000 126 334.244965 0.383722 24.27 25.43 26.34 26.13 26.67 26.41 -1.0000 127 334.250854 0.204667 24.83 25.92 25.48 25.89 26.93 29.16 -1.0000 128 334.248566 0.185944 24.71 26.27 26.44 26.94 26.80 26.19 -1.0000 129 334.248962 0.336250 23.98 25.02 25.84 26.29 26.37 26.25 -1.0000 130 334.258972 0.220361 22.25 23.54 24.17 24.44 24.74 24.50 -1.0000 131 334.266937 0.389944 24.89 25.91 25.83 26.24 25.28 25.67 -1.0000 132 334.276459 0.065806 24.96 26.10 26.77 99.00 28.22 99.00 0.0000 133 334.273682 0.444417 23.53 24.62 25.08 26.06 26.07 26.87 -1.0000 134 334.282776 0.153917 22.19 23.39 23.90 24.29 24.61 24.31 -1.0000 135 334.286285 0.083139 24.62 25.65 25.84 26.51 26.64 26.14 -1.0000
136
Table 4.4-Continued
No. aA. (J2000.0) Declo (J2000.0) N(AB) z'(AB) I R V B z.".., (1) (2) (3) (4) (5) (6) (7) (8) (9) (10)
136 334.288544 0.401167 23.66 25.08 25.57 26.08 27.04 25.91 -1.1JOOO 137 334.289459 0.459583 23.37 24.99 25.73 26.19 26.34 26.35 -1.1JOOO 138 334.295624 0.378667 24.99 26.10 25.76 26.21 26.71 25.97 -1.1JOOO 139 334.299133 0.461778 22.60 23.91 24.51 24.92 25.23 25.07 -1.1JOOO 140 334.306274 0.207806 24.95 26.21 25.36 25.84 26.39 27.89 O.IJOOO 141 334.314606 0.232694 24.07 25.17 25.37 25.82 26.04 25.40 -1.1JOOO 142 334.312714 0.430556 24.85 26.21 26.17 26.95 26.44 26.21 -1.1JOOO 143 334.318298 0.278333 22.81 24.42 25.42 25.59 27.11 25.88 0.8186 144 334.321899 0.158528 24.83 26.10 26.49 26.53 26.68 26.36 -1.1JOOO 145 334.321381 0.169139 24.13 25.34 26.11 26.70 26.68 26.30 -1.1JOOO 146 334.328796 0.130306 23.73 25.11 26.04 26.47 26.40 26.12 0.8217 147 334.335785 0.119333 22.80 23.87 24.36 24.85 25.22 25.13 0.8216 148 334.337494 0.075500 23.71 24.97 25.36 26.13 26.35 26.35 0.8259 149 334.343201 0.262417 24.56 26.68 26.18 26.94 26.85 26.70 0.8342 150 334.342255 0.429583 24.85 26.02 26.07 26.24 26.25 26.03 -1.1JOOO 151 334.340485 0.246972 21.56 23.11 24.27 24.41 24.62 24.41 -1.1JOOO 152 334.347046 0.361917 24.97 26.28 26.93 27.30 99.00 28.50 -1.1JOOO 153 334.347046 0.361917 22.38 23.43 23.96 24.37 24.77 24.64 -1.1JOOO 154 334.347870 0.272389 21.91 23.41 24.32 24.49 24.81 24.59 -1.1JOOO 155 334.357880 0.050472 23.99 25.00 25.22 25.83 99.00 99.00 -1.1JOOO 156 334.357422 0.188667 23.93 25.00 25.82 26.53 26.92 27.27 0.0000 157 334.358795 0.221861 23.55 25.25 25.79 26.39 27.25 26.10 -1.1JOOO 158 334.362213 0.438750 24.37 25.57 26.27 26.87 26.76 27.81 -1.1JOOO 159 334.363953 0.197278 22.73 23.84 24.05 24.01 24.40 24.47 0.3954 160 334.365082 0.146944 23.49 24.87 25.65 25.93 26.96 26.16 -1.1JOOO 161 334.374329 0.077667 24.11 25.24 25.85 26.02 26.24 26.01 -1.1JOOO 162 334.377350 0.187500 22.63 23.74 24.24 24.70 25.10 24.76 -1.1JOOO 163 334.383057 0.262278 22.58 23.69 24.27 24.74 24.98 24.70 0.8338 164 334.384857 0.153694 23.24 24.87 26.01 26.37 27.46 26.04 -1.1JOOO 165 334.385132 0.267861 24.97 26.02 25.82 26.00 26.41 26.43 -1.1JOOO 166 334.392883 0.300056 24.21 25.59 26.40 27.15 27.85 27.17 -1.1JOOO 167 334.391113 0.440472 24.98 26.12 25.38 25.83 26.25 26.75 -1.1JOOO 168 334.393585 0.133917 23.88 25.01 25.04 25.79 26.02 25.71 -1.1JOOO 169 334.400269 0.341667 23.17 24.25 24.44 24.81 25.08 24.93 - 1.1JOOO 170 334.400940 0.208083 24.65 25.75 25.83 26.23 26.99 26.74 1.4471 171 334.405212 0.343194 24.20 25.39 25.38 25.69 26.23 25.74 -1.1JOOO 172 334.407562 0.638278 23.75 24.97 25.37 25.79 26.49 26.68 0.8262 173 334.411072 0.445222 24.56 25.74 25.69 25.75 25.83 25.33 -1.1JOOO 174 334.409973 0.152639 24.82 27.21 99.00 99.00 99.00 29.19 6.4825 175 334.417572 0.137333 24.SO 26.41 28.55 27.75 27.78 99.00 O.IJOOO 176 334.417572 0.132167 24.34 25.60 25.54 25.93 26.14 25.42 -1.1JOOO 177 334.417664 0.176389 24.92 26.14 26.56 26.63 27.18 26.33 -1.1JOOO 178 334.416382 0.295222 24.28 25.70 26.19 26.76 26.02 26.22 -1.1JOOO 179 334.417511 0.208222 24.63 25.64 25.66 25.95 25.82 25.86 -1.1JOOO ISO 334.421295 0.526306 24.37 26.96 102.35 102.18 101.99 101.63 6.5047
137
Table 4.4----COntinued
No. R.A. (J2000.0) Decl. (J2000.0) N(AB) z'(AB) I R V B z.,.,c (1) (2) (3) (4) (5) (6) (7) (8) (9) (10)
181 334.422882 0.406028 23.77 24.93 24.70 24.81 25.43 25.72 -1.0000 182 334.429260 0.302361 24.16 25.95 99.00 99.00 99.00 99.00 6.4700 183 334.430176 0.204556 24.19 25.32 25.67 26.11 26.21 26.18 -1.0000 184 334.432739 0.308333 23.56 24.73 24.96 25.47 25.38 25.37 -1.0000 185 334.438249 0.468556 23.56 24.65 24.84 25.44 25.91 25.50 -1.0000 186 334.437714 0.251944 24.97 26.00 26.40 26.47 27.09 26.95 0.0000 187 334.437531 0.464917 22.45 23.54 23.78 24.25 24.62 24.37 -1.0000 188 334.441467 0.460444 23.18 24.71 25.30 25.61 25.46 25.65 0.0000 189 334.442108 0.214167 23.28 24.47 25.56 25.65 26.35 26.04 0.8152 190 334.443359 0.303500 23.11 24.74 25.57 25.72 26.19 25.47 0.8182 191 334.445374 0.250750 23.23 24.78 25.41 25.47 25.71 25.70 0.8221 192 334.445465 0.355389 24.01 25.72 25.59 26.14 26.71 26.03 -1.0000 193 334.453400 0.180944 24.76 26.39 26.14 26.48 26.28 27.78 -1.0000 194 334.456696 0.807056 24.85 27.69 102.35 102.18 101.99 101.63 0.0000 195 334.454742 0.208917 24.30 25.59 25.61 25.99 25.87 25.53 -1.0000 196 334.458710 0.355278 24.04 25.20 25.23 25.33 26.09 25.21 -1.0000 197 334.460144 0.057361 24.10 25.52 25.96 26.13 26.96 27.80 0.8207 198 334.465179 0.435194 24.98 26.74 25.97 26.52 26.81 26.85 -1.0000 199 334.465454 0.133528 24.96 26.17 26.06 26.56 27.41 26.90 -1.0000 200 334.466522 0.312611 20.84 22.69 23.71 23.75 23.93 23.54 -1.0000 201 334.468140 0.601722 24.55 25.70 26.16 25.92 26.28 26.26 -1.0000 202 334.471039 0.149167 23.60 24.66 25.06 25.42 25.90 25.17 -1.0000 203 334.474854 0.291944 24.30 25.32 25.31 25.87 26.66 26.14 -1.0000 204 334.474457 0.469000 24.77 25.94 24.89 25.06 25.76 25.64 -1.0000 205 334.476166 0.686500 22.35 23.68 24.32 23.31 24.06 24.51 0.3981 206 334.477966 0.130000 25.00 26.53 25.89 26.31 26.88 26.40 -1.0000 207 334.477264 0.060083 23.66 24.77 25.43 25.95 26.78 26.02 0.8225 208 334.477753 0.470667 24.62 25.63 25.48 25.67 25.70 26.03 -1.0000 209 334.480377 0.502056 24.30 25.62 26.85 26.58 26.92 27.52 0.8233 210 334.483337 0.593694 24.99 26.72 26.97 27.14 28.08 28.38 0.8270 211 334.482147 0.255861 24.86 26.49 27.81 29.77 99.00 29.35 0.0000 212 334.483337 0.357194 24.78 25.88 25.36 26.01 26.36 26.61 -1.0000 213 334.483551 0.238139 24.74 26.23 25.91 25.91 26.76 27.09 -1.0000 214 334.486786 0.155917 24.97 26.03 26.31 26.63 26.46 26.65 -1.0000 215 334.485626 0.289028 24.84 26.26 25.45 25.70 26.12 25.81 -1.0000 216 334.486708 0.305222 24.84 27.27 28.97 27.36 99.00 27.14 0.0000 217 334.490645 0.613556 24.54 25.56 25.78 26.02 26.29 26.30 0.0000 218 334.506470 0.236556 23.95 25.48 26.62 26.50 99.00 27.27 0.0000 219 334.505402 0.282750 21.86 23.05 23.73 23.98 24.33 24.06 0.8206 220 334.506073 0.312167 24.96 26.15 25.70 26.09 27.98 26.88 -1.0000 221 334.506470 0.326278 24.89 26.01 25.21 25.69 99.00 31.08 -1.0000 222 334.505066 0.473056 23.85 24.88 25.40 25.94 26.41 26.41 0.8233 223 334.507080 0.274083 22.86 24.14 24.68 25.34 99.00 25.87 0.8209 224 334.506805 0.327611 24.59 25.67 25.37 26.01 99.00 26.11 -1.0000 225 334.507874 0.812278 24.62 28.39 27.99 27.12 26.65 27.45 6.5180
138
Table 4.4-Continued
No. aA. (J2000.0) Decl. (J2000.0) N(AB) z'(AB) 1 R V B zspec
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10)
226 334.510406 0.374194 24.84 26.00 25.82 25.87 25.87 26.05 -1.0000 227 334.513153 0.468778 24.72 26.03 26.45 25.81 26.12 26.24 -1.0000 228 334.515198 0.462917 24.81 25.84 25.05 25.45 25.64 25.75 -1.0000 229 334.513153 0.468778 24.75 26.09 25.50 26.06 26.13 27.32 -1.0000 230 334.516235 0.425500 24.88 26.00 25.75 25.99 25.72 26.06 -1.0000 231 334.516235 0.379750 24.99 26.03 25.83 25.83 25.79 26.39 -1.0000 232 334.523346 0.325583 23.33 24.44 24.44 24.84 25.11 24.70 0.8420 233 334.520691 0.446389 23.42 24.81 24.87 25.28 25.78 25.72 -1.0000 234 334.523499 0.428639 24.96 26.23 25.70 26.55 26.05 26.20 -1.0000 235 334.526031 0.361667 24.74 26.05 25.89 26.99 27.19 26.65 - 1.0000 236 334.525177 0.359083 24.66 25.90 26.02 27.44 -99.00 27.10 0.0000 237 334.537323 0.349167 22.68 23.76 23.86 24.10 24.19 24.01 1.4864 236 334.537323 0.490111 22.99 24.00 24.44 24.61 25.02 25.06 0.8180 239 334.642847 0.440722 21.10 22.13 22.29 22.77 23.12 23.04 0.8273 240 334.541168 0.398083 23.22 24.85 25.23 25.66 26.12 25.60 0.8293 241 334.541168 0.398083 24.77 26.01 25.92 26.73 26.60 27.60 -1.0000 242 334.543762 0.062000 22.70 24.06 24.64 25.01 25.28 25.01 -1.0000 243 334.543762 0.311417 24.88 26.30 25.29 25.68 26.28 26.97 -1.0000 244 334.545105 0.067000 24.79 26.95 99.00 99.00 99.00 99.00 0.0000 245 334.548492 0.052167 24.70 25.78 25.35 25.39 25.46 99.00 -1.0000 246 334.548462 0.335722 24.80 26.69 26.01 26.38 26.61 25.89 -1.0000 247 334.548187 0.398528 24.67 25.83 25.83 26.97 99.00 99.00 -1.0000 248 334.575500 0.534444 24.54 25.61 25.86 25.92 26.38 26.51 -1.0000 249 334.580109 0.606139 22.97 24.13 24.87 24.93 25.17 25.29 0.8274 250 334.583466 0.681194 24.29 25.39 26.21 26.09 26.21 26.21 0.8295 251 334.591522 0.653972 24.95 26.26 26.84 26.75 27.16 27.42 -1.0000 252 334.605072 0.570444 23.69 24.70 25.30 25.24 25.57 25.60 0.8294 253 334.614990 0.423056 24.92 27.47 26.56 27.11 31.25 27.91 -1.0000 254 334.615173 0.791028 24.89 28.16 102.35 102.18 101.99 101.63 6.5050 255 334.625763 0.424972 24.60 25.67 25.99 25.77 26.04 26.08 -1.0000 256 334.623962 0.524361 23.15 24.24 24.61 24.05 24.64 24.98 -1.0000 257 334.627350 0.752278 24.39 25.75 26.14 26.19 27.05 26.94 -1.0000 258 334.638947 0.521500 24.05 25.18 25.72 25.67 26.00 26.05 -1.0000 259 334.647644 0.507778 23.63 25.13 25.34 25.23 25.75 25.92 -1.0000 260 334.654755 0.814722 24.58 25.70 25.76 25.72 26.20 26.90 0.8289 261 334.666840 0.422056 24.21 25.25 25.13 25.14 25.26 25.07 1.4627 262 334.670563 0.618389 24.42 25.52 26.00 25.95 26.16 26.41 0.8252 263 334.668457 0.684972 24.94 26.78 26.48 26.62 27.32 26.73 -1.0000 264 334.670563 0.646028 24.89 25.93 26.43 26.38 26.45 27.22 0.0000 265 334.687927 0.550639 24.93 26.20 26.60 26.50 26.73 26.98 0.0000 266 334.721222 0.693750 22.68 24.08 24.76 23.56 24.34 24.68 0.3987 267 334.749329 0.561972 23.59 24.62 25.20 25.29 25.40 25.56 0.8269 268 334.750763 0.582028 23.94 25.60 26.20 26.14 26.65 26.54 0.8276 269 334.751465 0.471639 24.54 25.55 25.74 26.23 27.36 26.58 0.0000 270 334.776215 0.789361 23.24 24.48 25.38 25.29 25.50 25.52 0.8176
139
Table 4.4-Continued
No. R.A. (J2000.0) Decl. (J2000.0) N(AB) z'(AB) 1 R V B zspee
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10)
271 334.788086 0.441750 23.95 25.18 25.44 25.60 25.96 25.80 0.8320 272 334.792847 0.570306 24.93 25.94 25.99 25.98 26.33 26.19 -1.0000 273 334.794342 0.735417 23.72 24.80 25.15 25.33 25.61 25.78 0.8270 274 334.796967 0.674806 24.63 25.72 26.44 26.16 26.43 26.75 -1.0000 275 334.797302 0.724806 24.89 26.30 26.44 25.97 26.97 27.68 -1.0000
Note. - Magnitudes are measured In 3" diameter apertures. An entry of '-99' Indicates that no excess Rux was measnred, '99' means that there was a negative Rus In the aperture. -1.0000 In the redshift colnmn means no spectroscopic data were obtained for the object, 0.0000 means that spectra were obtained, but the object was not identified, while -2.0000 denotes a star.
140
..... "'" .....
Table 4.5. Line fluxes and Oxygen Abundance for NB816 selected USELs
Object f([OIII]M007) f([OlIIj.).4959) f([OlIIj.M363) f([OIIj>.3727) Te[OIII] 12+log(O/H) 12+Iog(O/H) •
[OIII] emitters 31 513.6 ± 24.0 222.3 ± 11.4 < 6.60 54.4 ± 4.62 < 1.19 > 8.09 7.74 ± 0.02 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 7.84 ±0.02 51 401.5 ± 12.6 146.9 ± 5.52 9.40 ± 4.50 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 7.89 ± O.or
118 492.6 ± 29.7 193.9 ± 12.9 34.4 ± 12.11 11.3 ± 2.61 2.16 < 3.08 < 4.32 6.93 < 7.16 < 7.44 7.67±0.03 195 335.0 ± 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 7.53 ± 0.03 206 597.1 ± 19.5 204.1 ± 7.41 21.6 ± 9.20 < 1.72 1.48 < 1.97 < 2.48 7.42 < 7.55 < 7.84 7.75±0.02 208 658.0 ± 30.9 249.8 ± 12.3 15.6 ± 8.00 < 22.1 1.16 < 1.56 < 1.93 7.67 < 7.85 < 8.22 7.84 ± 0.03 223 242.9 ± 15.3 83.3 ± 7.53 23.7 ± 18.9 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 7.72 ± 0.01
Note. - Only emitters with> 150- Hi3 fluxes are 1isted. All fluxes are normalized by their f(Hi3) and multiplied by 100. 1<7 upper limite are 1isted for [011]>.3727 flux below au and [OIII].M363 below 1<7. The [OIIjfluxes are relatively uncertain 88 diBcussed in the text. The [011]>.3727 llnes of ID 41 and ID 223 are beyond the DEIMOS wavelength coverage. The unite of Te[OIIIj are Hr' [Kj •
• Empirically determined meta1Ucity using R23 based on Yin et aI. (53) shown for comparison. This method generally favors measurements at the upper end of the direct abundance range for galaxies with the lowest abundance determinations. In computing this we have assumed that the sources lie in the low meta1Ucity regime.
Table 4.6. Line Fluxes and Oxygen Abundances for NB912 selected USELs
Object f([OllljA5007) f([OIIIjA4959) f([OIII]A4363) f(jOIIjA3727) T.jOIII] 12+log(O/H) 12+log(O/H) a
[OIllj 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 7.71 ±0.02 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 7.79±0.03 9 442.3 ± 15.3 154.7 ± 6.42 < 12.3 157.2 ± 6.68 < 1.70 > 7.70 7.72±0.02
10 490.0 ± 11.9 178.7 ± 5.14 13.7 ± 4.40 128.6 ± 5.3 1.42 < 1.69 < 1.97 7.59 < 7.72 < 7.89 7.74 ± 0.02 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 7.44±0.03 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 7.55±0.02
239 202.4 ± 10.2 75.6 ± 6.40 < 8.20 190.6 ± 10.2 < 2.08 > 7.34 7.43±0.02 270 351.7 ± 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 7.50 ± 0.02
Ha 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 7.73±0.02 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 7.81 ±0.02
..... 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 4.5 but for NB912 emitters. The [0II]A3727 line of ID 266 Is beyond the DEIMOS wavelength coverage .
• Empirically determined metallicity using R23 hased on Yin et aI. (53) shown for comparison. ThIs method generally favors measurements at the upper end of the direct abundance range for galaxIes with the lowest abundance determinations. In computing this we have assumed that the sources lIe in the low metallicity regime.
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5.1 Summary
Chapter 5
Summary and Future Work
We have presented the results from the Hawaii Quasar and T dwarf survey (HQT survey),
which is a wide-field, red optical survey carried out with the Subaru/Suprime-Cam. The
aims of our survey are: (1) to identify low-luminosity quasars (Ml450 < -22.5) at high-z
(z > 5.7), which are considered to form the bulk of quasar population at high-redshift, and
(2) to identify faint, field T dwarfs and understand their number density and distribution
in the Galaxy.
We developed a new color selection technique to break a well-known 1- z' color
degeneracy between high-z quasars and foreground M and L dwarfs, which are much more
numerous than quasars. In an area of 9.3 deg2 at z'AB < 23.3, we have successfully
identified six new T dwarfs that are confirmed by near-ffi imaging and/or spectroscopy.
With estimated distances of 60-170 pc, these are among the most distant spectroscopically
confirmed brown dwarfs to date. Our finding rate of T dwarfs is consistent with the
model predictions in which the distribution of dwarfs follow the exponential Galactic disk.
However, the large brightening amplitude (~ 1 mag) previously reported for the L/T
transition objects appear to overpredict the numbers. Our results agree well with the
recent relations of Lin et al. (in preparation) and indirectly suggests a high binary fraction
rate for these transition objects.
147
No z > 5.7 quasars have been found in the HQT survey, supporting the current view
of high-z quasars being negIigible contributor (~ 10%) to the cosmic reionization. The
nondetections of low luminosity quasars in our survey BUggeBts that there is a break in the
quasar luminosity function (LF) at -25 < Mi4&J < -24.3 at 99% confidence level. Current
major quasar surveys are too shallow to probe the faint end of the quasar LF, though they
have successfully identified about two dozen luminous quasars (e.g., SDSS main - Fan et
al. 2006; SDSS deep - Jiang et al. 2007). Nondetections of z > 5.7 quasars in deep surveys
(e.g., CDF-N - Barger et al. 2004; CFHQS-deep - Willott et al. 2005), however, could
not set the limits tight enough to see the "break" due to their small area coverage. With
our area coverage and depths, we were able to see the "break" in the quasar LF at z ~ 6
for the first time, and set limits on the faint-end slope Q to be Q ~ -2.1 at 99% confidence
level.
The strong emission line galaxies at z < 1 were selected as candidate quasars in the
HQT survey. This may raise a problem with the current quasar color selection technique
when it is applied to search for low luminosity quasars - since the colors currently used to
select quasars (I -z', z' -J, or NB816-z') cannot differentiate high-z quasars from these
low-z emission line galaxies. Such low-z gala.xies have been known to be contaminants for
high-z galaxy surveys and thus imaging in the blue bandpasses have been used to separate
them from high-z objects. However. in the current high-z quasar surveys which only probe
extremely bright quasars (e.g., SDSS) or cover small area (e.g., CFHQS), these emission line
galaxies have not been found as candidate quasars, and thus they have not been considered
to be contaminants. Therefore, our findings may tell us that one should be aware of these
low-z emission line galaxies when conducting surveys for low luminosity quasars.
The nature of these emission line gala.xies, however, appear to be of great interest.
They have high ionization (q > 3.0 x lOS), star formation rates (10 - 2OM0/year), and H,8
equivalent widths (EWrest > 2OOA). Together with their low metal contents (Z < 0.2Z0 ),
they appear to be young systems which have just undergone major starbursts within a few
Myr.
148
In order to systematically search for these galaxies, we have used narrowband selected
samples from the high-z Lya galaxy survey fields with ultradeep multiwave1ength data. We
termed these galaxies as Ultra-Strong Emission Line galaxies (USELs) (EW(Hp) ~ 3OA).
Keck/DEIMOS spectra of 161 USELs have shown that their star formation density is about
5 - 10% of the value found from UV continuum objects at z = 0 - 1, suggesting that
they are significant contributor to cosmic star formation at z < 1. Many of the USELs
show [OIII].M363 auroral lines, which can be detected only in the very low-meta1licity
galaxies. Indeed, our spectral analysis have shown that at least about a dozen USELs
have oxygen abundances satisfying the criteria of eXtremely Metal Poor Galaxies (XMPGsj
Z < 1/12Z0 ). Since known XMPGs are all within z < 0.05, XMPGs found by our survey are
the most distant known today and our high yield rate of XMPGs suggest that narrowband
method is powerful in finding such populations.
5.2 Future Work
5.2.1 A new color criteria
The findings of z < 1 galaxies as candidate high-z quasars calls a necessity of a new color
technique which can separate high-z quasars from these low-z galaxies. We plan to obtain
deep imaging data in bluer bandpasses (B, V) to see if the addition of blue imaging help
to clarify this problem.
5.2.2 Quasar contribution to cosmic reionization
Based on the nondetection of z > 5.7 quasars in the HQT survey, we gave constraints on
the faint-end slope a and break magnitude Ml450 of the quasar LF at z ~ 6. The lack
of low luminosity quasars in the previous surveys have already shown that contribution of
quasars to the cosmic reionization is less than 9% (Shan!mr & Mathur). However, tighter
constraints on the quasar LF determined by our survey are expected to give better estimates
149
on the quasar contribution to reionization. Our next step is to determine the exact quasar
contribution to the ionizing background in the early universe.
5.2.3 Multiwavelength Studies of USELs at z < 1
About a dozen USELs with Keck/DEIMOS spectra are located in the GOODS field. The
wealth of existing deep, multiwavelength data in GOODS field allow us to investigate these
galaxies in more detailed. As a first step, we plan to estimate the stellar masses of these
galaxies using K-band data. It will be also interesing to study the dust content of these
galaxies using Spitzer data. Since we can estimate the age of galaxy from its Balmer line
equivalent widths, we may be able to probe the evolution of dust content as a function of
time after starbursts.
150