Table Of ContentPreprinttypesetusingLATEXstyleemulateapjv.12/16/11
A COHERENT STUDY OF EMISSION LINES FROM BROAD-BAND PHOTOMETRY: SPECIFIC
STAR-FORMATION RATES AND [OIII]/Hβ RATIO AT 3<Z <6
A. L. Faisst1,2,†, P. Capak1,2, B. C. Hsieh3, C. Laigle4, M. Salvato5, L. Tasca6, P. Cassata6, I. Davidzon6,7, O.
Ilbert6, O. Le Fe`vre6, D. Masters1,2, H. J. McCracken4, C. Steinhardt1,2, J. D. Silverman8, S. de Barros7, G.
Hasinger9, N. Z. Scoville2,
1InfraredProcessingandAnalysisCenter,CaliforniaInstituteofTechnology,Pasadena,CA91125,USA
2CahillCenterforAstronomyandAstrophysics,CaliforniaInstituteofTechnology,Pasadena,CA91125,USA
3 AcademiaSinica,InstituteofAstronomy&Astrophysics,P.O.Box23-141,Taipei10617,Taiwan
4Institutd’AstrophysiquedeParis,CNRS&UPMC,UMR7095,98bisBoulevardArago,75014,Paris,France
5MaxPlanckInstitutfu¨rExtraterrestrischePhysik,Giessenbachstrasse1,D-85748,GarchingbeiMu¨nchen,Germany
6 6AixMarseilleUniversit´e,CNRS,LAM(Laboratoired’AstrophysiquedeMarseille)UMR7326,13388,Marseille,France
1 7INAF-OsservatorioAstronomicodiBologna,viaRanzani1,I-40127,Bologna,Italy
0 8KavliInstituteforthePhysicsandMathematicsoftheUniverse(WPI),TheUniversityofTokyoInstitutesforAdvancedStudy,The
2 UniversityofTokyo,Kashiwa,Chiba277-8583,Japan
9 InstituteforAstronomy,2680WoodlawnDr.,UniversityofHawaii,Honolulu,HI96822,USAand
n submitted to ApJ - January 28, 2016
a
J
ABSTRACT
6
2 We measure the Hα and [OIII] emission line properties as well as specific star-formation rates (sSFR)
of spectroscopically confirmed 3<z <6 galaxies in COSMOS from their observed colors vs. redshift
evolution. Our model describes consistently the ensemble of galaxies including intrinsic properties
]
A (age, metallicity, star-formation history), dust-attenuation, and optical emission lines. We forward-
model themeasuredHαequivalent-widths(EW)toobtainthesSFRouttoz ∼6withoutstellarmass
G
fitting. We find a strongly increasing rest-frame Hα EW that is flattening off above z ∼ 2.5 with
h. averageEWsof300−600˚Aatz ∼6. ThesSFRisincreasingproportionalto(1+z)2.4 atz <2.2and
p (1+z)1.5 at higher redshifts, indicative of a fast mass build-up in high-z galaxies within e−folding
- times of 100−200Myr at z ∼ 6. The redshift evolution at z > 3 cannot be fully explained in a
o
r picture of cold accretion driven growth. We find a progressively increasing [OIII]λ5007/Hβ ratio out
t to z ∼6, consistent with the ratios in local galaxies selected by increasing Hα EW (i.e., sSFR). This
s
demonstrates the potential of using “local high-z analogs” to investigate the spectroscopic properties
a
[ and relations of galaxies in the re-ionization epoch.
Keywords: galaxies: evolution – galaxies: high-redshift – galaxies: star formation
1
v
3 1. INTRODUCTION lar mass is expected for young galaxies at z ∼ 5 (see
7
Faisst et al. 2015). These recent observations have trig-
1 With current broad-band near- to mid-infrared (IR)
gered questions that have yet to be answered. For ex-
7 filtersonground-andspace-basedtelescopesweareable
ample, galaxies have been found that are more massive
0 to select galaxy samples in the very early epochs of the
thanexpectedfromhierarchicalassemblyofdark-matter
. universe. However, the study of their physical proper-
1 haloes (e.g., Steinhardt et al. 2015). The formation of
ties – essential to refine our understanding of the forma-
0 suchmassivegalaxiesathighredshiftsrequirestheirfast
tion and evolution of galaxies – is hampered by several
6 growth and therefore an increase in the specific SFR
technical problems.
1 (sSFR = SFR/M, a measure for the rate of mass build-
In the recent years, a lot of progress has been made
: up in galaxies) at z > 3 (e.g., Weinmann et al. 2011).
v in understanding galaxy formation in the early universe
i before the peak of the cosmic star-formation density at This increase is also predicted in the picture of accre-
X tion dominated galaxy growth (e.g., Dekel et al. 2009;
z ∼ 2−3. In particular, several new avenues have been
r opened by large spectroscopic and photometric cam- Tacchella et al. 2013) and recent hydrodynamical simu-
a lations (e.g., Dav´e et al. 2011; Sparre et al. 2015). Some
paigns to explore the near- to mid-IR wavelength range
studies observe these predictions (Stark et al. 2013; de
on large parts of the sky. From these, it became clear
Barros et al. 2014; Jiang et al. 2015), while other find a
that galaxies live on a so called “main-sequence” con-
considerable flattening of the sSFR at z > 3 (Gonz´alez
necting their stellar mass with their star-formation rate
et al. 2014; Tasca et al. 2015; Marmol-Queralto et al.
(SFR)outtoredshiftsashighasz ∼5(Steinhardtetal.
2015). Finally,relationsbasedonlocalgalaxies,e.g.,the
2014; Speagle et al. 2014; Tasca et al. 2015). Also, it
relations between metallicity and strong emission lines,
is suggested that galaxies grow very rapidly in the early
maynotbeapplicableanymoreathigherredshiftsdueto
universe due to high gas fractions and/or star-formation
thechangeofinternalphysicalprocessesinsuchgalaxies
efficiencies (e.g., Scoville et al. 2015; Silverman et al.,
submitted). Going in hand with the former, a marginal such as ionization or the abundance of [NII] (e.g., Mas-
ters et al. 2014; Steidel et al. 2014; Shapley et al. 2015;
flattening of the relation between metallicity and stel-
Cowie et al. 2015).
Electronicaddress: [email protected] Stellar mass and SFRs as well as metallicity and ion-
†Twitter: @astrofaisst ization parameter are the most important basic physi-
2 A. L. Faisst, et al.
cal quantities on which the above results are based on lines. Intheresultssection(Section4),wederivethered-
and the above questions depend on. While these can shiftevolutionoftheHαEW,thesSFR(z),aswellasthe
be measured reliably at low redshifts by a good multi- [OIII]/Hαratioouttoz ∼6. Theseresultsarediscussed
wavelength coverage in imaging and spectroscopy, there in Section 5 and summarized in Section 6.
are several caveats at higher redshifts. First, SFRs have Throughout this work we adopt a flat cosmology with
to be measured in the UV, as reliable estimators such as Ω = 0.7, Ω = 0.3, and h = 0.7. Magnitudes
Λ,0 m,0
the Hα emission line are out of spectral coverage. The are given in the AB system (Oke & Gunn 1983) and
UVishighlysensitivetodustattenuation(e.g.,Bouwens all masses are scaled to a Chabrier (2003) initial mass
et al. 2012b), which shows a large diversity in high red- function (IMF).
shift galaxies (see Capak et al. 2015). Second, deep ob-
served mid-IR imaging data are necessary to probe the 2. DATA&SAMPLESELECTION
old stellar populations in galaxies at z >4 and therefore 2.1. Data
allow a reliable measurement of stellar masses. Third,
Inthiswork,weusethetwosquaredegreesoftheCos-
mid-IR filters at these redshifts are contaminated by the
mic Evolution Survey (COSMOS, Scoville et al. 2007)
(unknown) contribution of strong emission lines, which
field, which are observed by a wealth of instruments in
boostthemassessignificantly(e.g.,Schaerer&deBarros
imaging as well as spectroscopy across a broad range of
2009; Stark et al. 2013; de Barros et al. 2014). Finally,
wavelengths. We make use of the following data sets.
the conversion from the observed data to these quanti-
ties (i.e., stellar mass and SFR) depends on theoretical
1. The COSMOS spectroscopy catalog, which con-
modelsoftheintrinsicpropertiesofgalaxiessuchastheir
tains more than 6000 high-quality spectra at
age, metallicity, and star-formation history (SFH), all of
1 < z < 6 (M. Salvato, private communication).
which are not known a priori for individual galaxies at
high redshifts.
2. The VIMOS Ultra Deep Survey (VUDS) spec-
In this paper, we develop a model insensitive way to
troscopy catalog, containing galaxy spectra at 2<
measure the sSFR and the emission line strength at
z <6 (Le F`evre et al. 2014).
3 < z < 6 from primary observables. Furthermore, we
demonstrate the potential of using local high-z analogs 3. The COSMOS2015 photometric catalog including
to probe the spectral properties of high redshift galaxies photometryfromtheUVtomid-IRaswellaspho-
up to z =6. In particular, we use the redshift evolution tometric redshifts and stellar masses (Laigle et al.
ofthegalaxypopulationaveragedobservednear-tomid- submitted).
IRcolortomeasuretheHαequivalentwidth(EW)from
whichwedirectlyderivethesSFR(z). Themeasurement The COSMOS spectroscopic master catalog available
of the EW has two parts, namely the measurement of to the COSMOS collaboration is a compilation of all
the observed flux/color and the estimation of the under- spectroscopic observations up to z ∼ 6 that are car-
lying continuum between 4500˚A and 6500˚A underneath ried out on the COSMOS field. The galaxy sample is
optical emission lines. The latter we forward model by selected in different ways (color, photometric redshift,
assumingintrinsicpropertiesofthegalaxies(dustatten- Lyman Break technique) and observed by several dif-
uation,metallicity,stellarpopulationage,starformation ferent instruments depending on the redshift (VIMOS,
history)andweshowthattheresultingcontinuumisvery FORS2, FMOS, MOIRCS, DEIMOS, MOSFIRE). The
insensitivetothechoiceoftheseparametersintheabove different selection techniques lead to a large coverage of
specified wavelength range. The determination of the physical properties of the galaxies, thus this sample rep-
Hα equivalent width (EW) from observed galaxy colors resents well the population of star-forming galaxies at
has been used in the past (Shim et al. 2011; Stark et al. these redshifts. For more information, we refer to the
2013; Smit et al. 2015a,b; Rasappu et al. 2015; Marmol- official COSMOS web-page2.
Queralto et al. 2015), however, mostly at discrete red- The VUDS spectroscopy catalog contains galaxies se-
shift bins and for small sample sizes. We perform here lectedbyphotometricredshiftswithafluxlimitofi =
AB
a consistent analysis across a large redshift range with a 25. The spectra are obtained with the VIMOS spectro-
muchlargersampleofspectroscopicallyconfirmedgalax- graph on the ESO Very Large Telescope (Le F`evre et al.
ies than in previous studies. Our large sample allows us 2003). For more information, we refer to Le F`evre et al.
to model the ensemble of galaxies instead of considering (2015).
single galaxies. This enables us to marginalize over the The COSMOS2015 photometric catalog contains the
(poorly known) intrinsic properties of the galaxies when photometry of the extracted galaxies on COSMOS mea-
modelingthecontinuumbelowtheopticalemissionlines. sured from the UV to the mid-IR images. The source
It also gives us a convenient way to describe the scatter extraction is based on a χ2 image determined from the
(systematic and physical) of the ensemble’s properties, Subaruz−bandandtheCOSMOS/UltraVISTAYJHK
which we can propagate through our model and investi- bands (see Capak et al. 2007; McCracken et al. 2012;
gate its effect on our results. Ilbert et al. 2013). Part of this catalog is the mid-IR
The plan for this paper is as follows. In Section 2 we data at 3.6µm and 4.5µm down to ∼ 25.5 mag (3σ in
describethesampleofspectroscopicallyconfirmedgalax- 3(cid:48)(cid:48) diameter; as of October 2015) from the Spitzer Large
iesthatisusedforthisanalysis. InSection3wedescribe Area Survey with Hyper-Suprime-Cam (SPLASH, Stein-
the modeling of the observed color vs. redshift relation hardt et al. 2014)3. These sources are extracted using
including the contribution of intrinsic parameters (age,
metallicity,SFH),dustattenuation,andopticalemission 2 http://cosmos.ipac.caltech.edu
3 http://splash.caltech.edu
Coherent study of emission lines and sSFR at 3<z <6 3
the following analysis we use the galaxy sample clear of
700 neighbors within 2(cid:48)(cid:48). The results do not significantly
600 spectroscopic change if we use the more restricted sample of galax-
500 final sample ies without contamination within 3(cid:48)(cid:48) radius (although
ber 400 the uncertainties are larger because of the strongly de-
m
u 300 creasednumberofgalaxies). Inparticular,intheredshift
n
200 range 3<z <6, we use 530 spectroscopically confirmed
100 galaxies. Figure 1 shows the stellar mass and redshift
0 distribution of our final sample of galaxies (without con-
log(M/M ) 11189012 llllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllll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l ottrsWlohaoengtimues(ttpMeihipnrnxeaoaep/ppnCMeeLotcOresr(cid:12)ttSiw)iosPoMuih∼nttrOaoharlgi9Snaeat.2c8o4l0c2ae1ta(cid:48)ix(cid:48)nhs5t)isec.ezslputUThhtd>Voehointehs3flogSa.smuFvttxeWheeRlte(laraeLyffirramstouimttgmerisldneaiensigspasegsrnoetittsfmhahselataea.metrrslSelyiutaEshmbsrmeDimoemeganifiaotassttalusutieslnridrneeeoo)dgs-ff.
ments (galaxy color) and we therefore do not use these
1.0 1.5 2.0 2.5 3.0 3.5 4.0 4.5 5.0 5.5 6.0
stellar masses in the following. They only serve to visu-
redshift
alize the expected mass range of our sample of galaxies
Figure 1. Propertiesofourfinalsample(galaxiesclearofcontam-
ination within 2(cid:48)(cid:48)). Top: Redshift distribution of our final spec- and the comparison to other studies.
troscopicsample. Bottom: Stellarmassdistributionasafunction
ofredshiftforourfinalspectroscopicsampleofgalaxies(seeLaigle 3. EMISSIONLINESTRENGTHSFROMOBSERVED
et al. submitted). The gray points show individual galaxies and COLORS
thebluesymbolsshowthemeanlog(M/M(cid:12))inredshiftbinswith In this section, we describe in detail our model includ-
68% percentiles scatter in redshift and mass shown by the error
bars. ing intrinsic galaxy properties (age, metallicity, SFH),
dust attenuation, and optical emission lines. From this
the segmentation map of the COSMOS2015 catalog and
we derive model galaxy colors as a function of redshift,
an improved version of IRACLEAN (Hsieh et al. 2012)
which are compared to the observed colors vs. redshift
in order to overcome the source confusion (blending).
evolution in our galaxy sample. This allows us, using a
We subsequently match the photometry catalog with
minimizationalgorithm, tosolveforthespectralproper-
thespectroscopycatalogswithin1(cid:48)(cid:48) radiusinordertore-
ties of the ensemble of these galaxies in specific redshift
coverthephotometryforourspectroscopicallyconfirmed
windows, which are detailed in the following.
galaxies. More than 97% of the galaxies are matched
within a radius of 0.3(cid:48)(cid:48).
3.1. Redshift windows and colors
2.2. Galaxy selection, redshift and stellar mass The idea of this paper is to constrain the emission line
distribution properties of galaxies from their observed colors. Emis-
sion lines contribute to different broad-band filters for
For the purpose of this work, we apply a very strin-
galaxies at different redshifts. This produces ”wiggles“
gent cut to our sample in terms of both spectroscopy
in the observed color-redshift evolution with respect to
as well as photometry. We only include reliable spec-
whatisexpectedfromacontinuumwithoutnebularlines.
troscopic redshifts in our sample (> 80% probability of
However, the observed color of a galaxy is not only af-
correct spectroscopic redshift) at 1 < z < 6 and re-
fected by emission lines, but also by its intrinsic proper-
move spectroscopically confirmed AGNs based on their
ties (age, metallicity, SFH) and dust attenuation. These
broad optical emission. We will use the redshift range
change with redshift and thus are degenerate with the
1 < z < 3 to verify our method by direct comparison of
effects of emission lines. In a later section we will dis-
our results to spectroscopic emission line measurements.
cuss how much these various properties affect the ob-
The measurement of colors strongly depends on source
served color of a galaxy. In order to separate the effect
confusion, which commonly is taken into account during
of emission lines, we have to calibrate our model in red-
the extraction of the galaxy photometry. For the pur-
shift ranges in which the continuum flux in broad-band
pose of this work, we add an additional security and re-
filtersisfreeofemissionlinesandthusrevealstheintrin-
move potentially blended sources in the near- to mid-IR
sic color and dust attenuation.
by directly checking their number of neighboring galax-
Figure 2 shows the location of strong emission lines
ies. For this to end, we use the high-resolution F814W
(Hα, Hβ, [OII], and [OIII]) in different near- and mid-
(I−band)imagesfromtheHubbleSpaceTelescope’sAd-
IR broad-band filters as a function of redshift. There
vanced Camera for Surveys (HST/ACS), as well as the
are several different redshift ranges (labeled for the case
COSMOS/UltraVISTAoptical/near-IRselectedcatalog.
3<z <6 and mid-IR colors):
We extract the number of companions within a certain
aperture size for each galaxy in our sample. Given the
• Redshiftrangesfreeofemissionlinesthatrevealthe
PSFaperturesizeof∼2(cid:48)(cid:48)−3(cid:48)(cid:48) inthemid-IR,weremove
intrinsic color and dust attenuation of the galaxies
all the galaxies with companions closer than 2(cid:48)(cid:48) and 3(cid:48)(cid:48)
and thus anchor our model (box labeled with “in-
in radius, respectively.
trinsic/dust”).
Our final sample at 1 < z < 6 with no contamina-
tion of neighboring galaxies within a radius of 2(cid:48)(cid:48) (3(cid:48)(cid:48)) 4 See Ilbert et al. (2006) and http://www.cfht.hawaii.edu/
contains more than 4,000 (1,500) galaxies in total. In ~arnouts/LEPHARE/
4 A. L. Faisst, et al.
Emission lines emission line free: continuum + continuum +
5.8µm
intrinsic color emission line: emission line:
[OII] [OIII] dust attenuation measure Hα(z) measure [OIII]/Hα
Hβ Hα
4.5µm
3.6µm
r
e
t
l 3.0 < z < 6.0
i
F
K
1.0 < z < 4.0
H
1.0 < z < 2.9
J
1.0 1.5 2.0 2.5 3.0 3.5 4.0 4.5 5.0 5.5 6.0
redshift
Figure 2. Contribution of different optical emission lines ([OII], green; [OIII], orange; Hα, red; Hβ, yellow) to broad-band filters as a
functionofredshift. Theblue,green,andredbandsshowthethreeredshiftwindows(redshiftsindicated)towhichweapplyouremission
line + dust model to estimate the Hα EW as well as the [OIII]/Hα line ratio. The dotted boxes show, for the 3 < z < 6 window as
textbookcase,howredshiftrangesareusedtorevealtheintrinsiccolor+dustattenuation(usedtoanchorourmodelfit),theEW(Hα)vs.
redshiftevolution,andtheratio[OIII]/Hα.
• Redshift ranges where the observed color includes The observed color of a galaxy is affected threefold by
theHαemissionlineandthusallowsustomeasure its properties: (i) By its intrinsic color (age, metallicity,
EW(Hα) (box labeled with “Hα”). SFH),(ii)bythedustattenuation,and,(iii)byemission
lines.
• Redshift ranges that allow the measurement of the In the following sections, we build up a model for the
theratio[OIII]/Hα(boxlabeledwith“[OIII]/Hα”). observed color in the different redshift windows from
these three contributions. From its comparison to the
These different redshift ranges exist for different ob- true observed colors, we can then compute the average
served colors and can be bundled in larger redshift win- emission line properties of our galaxies. It is important
dows. For the purpose of this work, we choose three tonotethattheintrinsicgalaxypropertiesaswellasdust
differentredshiftwindows,eachwithacorrespondingob- aresolelyusedtorepresentthecontinuumundertheHβ,
served color. [OIII], and Hα lines, i.e., red-ward of the 4000 ˚A break.
In particular, the fitting of our model continuum to the
(A) 1.0<z <2.9 in observed [H] − [K] color,
observedcontinuuminline-freewavelengthregionsusing
the contribution from dust as a “knob” (Section 3.2.2)
(B) 1.0<z <4.0 in observed [K] − [3.6] color, and
smooths out possible variations in the intrinsic galaxy
(C) 3.0<z <6.0 in observed [3.6] − [4.5] color. properties that are missed elsewhere in our model. This
isthemainadvantageofourforward-modelingtechnique
Each of the three redshift windows is designed to have and allows the robust estimation of the emission line
aredshiftrangefreeofemissionlinestoanchorthemodel properties as we show in the following.
to the intrinsic color. Furthermore, this choice allows us
3.2.1. The intrinsic color (age, metallicity, SFH)
to consistently model EW(Hα) across the redshift range
1 < z < 6 and the Hα/[OIII] ratio at z ∼ 2.2, z ∼ 3.3, For describing the intrinsic color of a galaxy popula-
and z ∼5.5. The redshift windows (A) and (B) are used tion, we have to assume a metallicity, stellar population
to verify our method by comparing our results to spec- age, and SFH. These are unknown a priori, therefore we
troscopic measurements. Given the strong dependence setupagridthatbracketsreasonablechoicesofthesepa-
of the 4000 ˚A Balmer break on stellar population prop- rameters. Our forward-modeling technique then allows
us to investigate the effects on the resulting observed
erties,wedonotmodelthe[OII]emissionlinehere. For-
color and the results derived in this work. As a basis
tunately,thewavelengthpartred-wardofthe[OII]emis-
we use the composite stellar population library from the
sion is relatively insensitive to the intrinsic properties of
Bruzual & Charlot (2003) with a Chabrier IMF and cre-
the stellar population as we will discuss later.
ateSEDswithdifferentSFHs, metallicities, andagesus-
3.2. Modeling the mean observed color as a function of ing GALAXEV5.
redshift
5 SeeBruzual&Charlot(2003)andhttp://www.bruzual.org/
Coherent study of emission lines and sSFR at 3<z <6 5
dust is of similar or larger amplitude (common dust ex-
tinctionsareontheorderofE(B−V)∼0.1−0.4mag in
0.4
Star−formation history our sample). Note, that the model uncertainties are sig-
constant nificantly reduced at high redshifts. First, the observed
0.3 delayed (peak at 1 Gyr) color is mostly independent of SFH for young galaxies
exp. inc. (t = 500 Myr)
with ages of less than ∼ 1 Gyr, i.e., z > 5. Second, it
m] 0.2 is expected that galaxies at high redshifts are dust poor
m5 (e.g., Dunlop et al. 2013; Bouwens et al. 2014; Capak
− [4. 0.1 Z = 0.02 et al. 2015). Third, their age and metal content is well
m] Z = 0.004 D E(B−V) = 0.2 dtiemfien.ed because of the young age of the universe at that
m3.6 0.0 e e Summarizing, we find that the expected reddening by
[ ers ers dust exceeds the effect of metallicity as well as SFH and
−0.1 univ univ ageforyounggalaxiesupto∼1Gyrinage(correspond-
−0.2 age of (z = 5) age of (z = 3) idpneegcclttionthizneg>inS5tF)rH.inF,sioacrscgiotalloiasrxttiheoescwchaiatsnhegaaetcmolonowrsetearsnirgtendoirsfihecixaftpnsot,lnyweenwteiitxah-l
108 109
age. Also the color starts to increasingly depend on the
age [yr]
assumed SFH for older galaxies and therefore lower red-
shifts. Finally, we note that a different IMF does not
Figure 3. Effect of intrinsic properties (metallicity, age, SFH)
anddustontheobserved[3.6] − [4.5]colorinthecaseofaz=5 change these conclusions. For example, using a Salpeter
galaxy. Metallicity plays only a minor role in setting the color IMFchangestheobservedcolorbylessthan0.01magat
of a galaxy (see arrow from 1/5th of solar to solar). The colored a given age.
bands show three different SFHs (constant, delayed, and expo-
nentially increasing). Clearly, dust attenuation (ranging between 3.2.2. Emission lines and dust
E(B−V)=0.1−0.4maginoursample,dependingonredshift)is
thedominantcontributortocolor,followedbytheSFHforgalaxies Besidesintrinsicproperties,dustattenuationandemis-
olderthan∼1Gyr(orz<5). Importantly,thecolorisinsensitive sionlinescontributetotheobservedcolorofagalaxy. We
toreasonableSFHsandstellarpopulationagesforyoung(<1Gyr)
derive all rest-frame UV and optical emission lines rela-
galaxiesathighredshifts.
tive to Hα, which we vary in our model. In detail, we
Galaxiesuptoz ∼4−5showatightrelationbetween parametrize the evolution of the (rest-frame) Hα EW as
SFR and stellar mass, which leads to an exponentially
increasing SFH for the average population of galaxies EW(Hα)=EW(Hα) ×(1+z)α, (1)
0
(Noeske et al. 2007; Daddi et al. 2007; Speagle et al.
2014;Steinhardt&Speagle2014;Smitetal.2015a). We where EW(Hα)0 and α are free fitting parameters.
bracket possible histories by a constant SFH, a delayed Furthermore, we assume a constant (with redshift)
exponentiallydecreasingSFH(SFR∝t/τ2×exp(−t/τ )) line flux ratio Hα/[OIII]= ξ within each of the three
p p
with a peak at τ = 1Gyr, and an exponentially in- redshift windows (see Section 3.1), because [OIII] only
p
enters in a narrow redshift range at z ∼ 2.2,3.3,5.5
creasing SFH (SFR ∝ exp(t/τ)) with an e-folding time
for windows (A), (B), and (C), respectively. In our
of τ = 500Myr. Furthermore, galaxies at high red-
shift show a considerably lower metallicity content (e.g., case, [OIII] denotes the blended doublet and we assume
Erb et al. 2006; Maiolino et al. 2008; Mannucci et al. [OIII]λ5007/[OIII]λ4960= 3.
TheHβ linefluxisdeterminedfromHαassumingcase
2009; Faisst et al. 2015). We therefore bracket the range
B recombination,
in metallicity between Z = 0.004 (1/5th of solar) and
Z = 0.02 (solar). However, because of the minor ef-
fect of metallicity on the continuum, we keep it constant f(Hα)/f(Hβ)=100.4×E(B−V)neb×(kβ−kα)×2.86, (2)
with redshift. Finally, we assume the age of the galaxy
wherek andk arethecoefficientsforagivendustat-
to be the time since the estimated start of re-ionization β α
at z = 11 (e.g., Planck Collaboration et al. 2015). Sim- tenuation curve (we assume here Calzetti et al. (2000)6)
ilar parameterizations of the galaxy’s age as a function atthewavelengthsofHαandHβ,respectively. The(stel-
of redshift (e.g., half of the Hubble time) do not change lar) dust extinction E(B−V) 7 is parametrized as ex-
stel
our results. ponentially decreasing function of redshift (e.g., Hayes
InFigure3,weshowtheeffectofintrinsicpropertiesas et al. 2011),
wellasdustattenuationontheobserved[3.6]−[4.5]color
on the example of a galaxy at z = 5. First, we empha- E(B−V)stel =E(B−V)0×e−z/zd,0 (3)
size the small effect of metallicity on the continuum: a
with E(B−V) and z as free parameters. We also
change from 1/5th of solar to solar metallicity results in 0 d,0
model weaker optical emission lines (e.g., [SII], [NII],
less than 0.05 mag change in color. Second, in the case
of an exponential increasing SFH with τ = 500 Myr,
6 Several studies indicate that high redshift galaxies follow the
the observed color as a function of age reddens less then dust attenuation curve similar to the one of the small Magellanic
0.1magforallpossibleagesofaz =5galaxy. Inthecase cloud. However, our model and data is not accurate enough to
ofaconstantanddelayedSFH,thereddeningisstronger disentangletheeffectofdifferentattenuationcurves.
7 We assume E(B−V) = E(B−V) /0.76 (Kashino et al.
due to domination by old stars with increasing age, but neb stel
2013). However,usingafactorclosetounityassuggestedbymore
less than 0.2 mag over a time of ∼ 2 Gyrs, which cor-
recent studies (Cullen et al. 2014; Shivaei et al. 2015; de Barros
responds to z ∼ 3. Compared to this, the reddening by etal.2015)doesnotaffecttheresultsofthispaper.
6 A. L. Faisst, et al.
Table 1
Summaryofobservationaldataandbest-fitmodels.
Dataandobservations Modelfit† Inputproperties
—————————————————————–
(spectroscopic) Emissionlines Dustb
———————————————————— ———————————— ————————— ————————————–
redshift color #<2(cid:48)(cid:48) #<3(cid:48)(cid:48) EW(Hα)0a α ξ E(B−V)0 zd,0 Z[Z(cid:12)] Age[yr] SFH
1.0<z<2.9 [H] − [K] 3571 1671 13.4 2.01 1.1 0.45 0.9 0.020 T(z)−T(11) constant
5.5 2.96 0.9 0.90 0.85 0.004 T(z)−T(11) constant
78.1 0.32 1.4 0.45 2.10 0.020 T(z)−T(11) exp. inc
57.2 0.72 1.0 0.85 1.40 0.04 T(z)−T(11) exp. inc
1.0<z<4.0 [K] − [3.6] 3863 1811 15.8 1.92 0.8 0.70 1.10 0.020 T(z)−T(11) constant
32.8 1.35 0.6 0.80 1.50 0.004 T(z)−T(11) constant
10.0 2.34 1.1 0.70 1.60 0.020 T(z)−T(11) exp. inc
21.0 1.71 0.7 0.90 1.80 0.004 T(z)−T(11) exp. inc
3.0<z<6.0 [3.6] − [4.5] 530 257 9.9 2.01 0.8c 0.90 1.25 0.020 T(z)−T(11) constant
13.6 1.83 0.8c 1.10 1.70 0.004 T(z)−T(11) constant
10.6 1.98 0.9c 0.80 1.60 0.020 T(z)−T(11) exp. inc.
8.5 2.10 0.8c 1.20 1.80 0.004 T(z)−T(11) exp. inc.
† The errors in the resulting EW(Hα) are estimated by a Monte-Carlo simulation to be ∼ 30%. The errors in ξ = ([OIII]/Hα)−1 are similarly
estimatedtobeontheorderof70%.
a Inangstromsandrest-frame.
b Strictlyspeaking,theseparametersdonotonlyincludedustbutalsochangesintheintrinsicSEDthatarenotincludedelsewhereinourmodel.
cThesearebestfitvalues,however,moreuncertainthanat1.0<z<2.9becauseonlypartofthewavelengthrangeincluding[OIII]andHαemission
linesiscoveredbyourdata. TheactualscatterinthesemeasurementswillbediscussedinmoredetailinSection4.3.
HeI), which we scale relative to the Hβ line fluxes ac- package8 and proceed in two steps.
cording to Anders & Fritze-v. Alvensleben (2003), as-
1. Wefitthedustattenuationasafunctionofredshift
suming reasonable 1/5th−1 solar metallicity. Although
(Equation3)inregionsdevoidofemissionlines(see
notparticularlystronginemission,theseaddupandcan
Figure 2). As mentioned above, this fit also in-
contribute up to 20% to fluxes in the broad-band filters.
cludes intrinsic changes of the SED that are not
Thecontributionsofdustandemissionlinesareadded
taken into account elsewhere in our model. The
to the intrinsic continuum described in the previous sec-
sole purpose of this is to model the continuum be-
tion. Theemissionlinesareaddedassumingafull-width-
at-half-maximum (FWHM) of 10 ˚A for the strong (Hα, low the optical emission lines (Hβ, [OIII], Hα).
Hβ, [OII], [OIII]) and 5 ˚A for weak emission lines. Note 2. WefixthevaluesofE(B−V)0 andzd,0 andfitthe
thatbecauseofthelargewidthofthebroad-bandfilters, remaining parameters EW(Hα)0, α, and ξ describ-
different (reasonable) choices of FWHM do not change ing the emission lines.
the following results. The model colors are obtained by
This procedure is important to break the degeneracies
convolution of the generated SED with the correspond-
betweentheeffectofdustattenuationandemissionlines
ing filter transmission curves (Spitzer/IRAC for 3.6 µm
on the observed color. We find that this is especially
and 4.5 µm and VISTA H and Ks band).
important at lower redshifts where galaxies show a sig-
nificantamountofdustbutmuchweakerEWscompared
to high-z galaxies.
We perform these two steps for in total four combi-
3.2.3. Fitting the observed color as a function of redshift
nations describing our intrinsic SED: 1/5th of solar and
solar metallicity and two SFHs (constant and exponen-
The top panel in Figure 4 shows the observed color in
tially increasing9). The bottom panel in Figure 4 visu-
redshift window (C), i.e., 3<z <6. The other two red-
alizes the fit for a constant SFH with solar metallicity
shift windows are shown in Appendix B. The symbols
at 3 < z < 6. The best-fit model (intrinsic + dust +
show individual galaxies (split in galaxies with no con-
tamination within in 2(cid:48)(cid:48) and 3(cid:48)(cid:48) radius, respectively) and emission lines) in solid red is shown along with the dust
reddened intrinsic color (blue dashed), and the intrin-
the blue band shows the running mean observed color
sic color (green dot-dashed). The horizontal lines label
(including 1σ scatter) as a function of redshift. In the
following, we fit this observed color vs. redshift evo-
lution with the previously described model for a given 8 https://cran.r-project.org/web/packages/minpack.lm/
index.html
combinationofmetallicityandSFH.WeuseaLevenberg- 9ThedelayedexponentiallydecreasingSFHyieldssimilarresults
Marquardt(LM)algorithm,aspartoftheR/minpack.lm thanaconstantSFHandwedonotlistithere.
Coherent study of emission lines and sSFR at 3<z <6 7
ll
2.0 ll ll
1.5 l no contamination < l2'' l data (weighted mean)
l no contamination < 3''
m] 1.0 l l ll ll l
mm[3.6m] − [4.5−−−11000.....50505 lllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllllll llllllll llllllllllllllllllllllllllllllllllllllllllllll ll
l
−2.0
3.0 3.5 4.0 4.5 5.0 5.5 6.0
l redshift ll
l
2.0
data (weighted mean) model: intrinsic + dust + emission
1.5
model: intrinsic + dust
]
m 1.0 model: intrinsic
m4.5 0.5
[
− 0.0
]
m−0.5
m3.6−1.0 free of emission lines Ha in 3.6m m Ha in Ha in
[ (fit dust extinction) 4.5m m 4.5m m and
−1.5 [OIII], Hb in
3.6m m
−2.0
3.0 3.5 4.0 4.5 5.0 5.5 6.0
redshift
Figure 4. Top: Observed color vs. redshift relation at z > 3. The open (filled) symbols denote galaxies with no contamination from
companions within a radius of 2(cid:48)(cid:48) (3(cid:48)(cid:48)) in ACS/F18W and ground based data. The blue line shows the weighted mean relation with
scatter(lightblueband)forgalaxieswithnocompanionwithin2(cid:48)(cid:48). Bottom: Thebest-fitintrinsic(blue,dashed),intrinsic+dust(green,
dot-dashed),andbest-fit(red,solid)model.
whataffectstheobservedcolorinagivenwavelengthre- Section 5.1).
gion(seealsoFigure2). Thebest-fitparametersforeach Our derived EW(Hα) in the redshift windows (A) and
redshift window and intrinsic SED are listed in Table 1. (B)areinexcellentagreementwithdirectdeterminations
fromspectroscopyatz <3obtainedbyErbetal.(2006)
4. RESULTS at z ∼ 2, Fumagalli et al. (2012) (at 1 < z < 2 as part
The model described in the previous section allows us of 3D-HST; van Dokkum et al. 2011; Brammer et al.
to fit the redshift dependence of the Hα EW as well as 2012;Skeltonetal.2014),andLamareilleetal.(2009)(at
the [OIII]/Hα line ratio from the observed color vs. red- z ∼0.5,aspartofVVDS;LeF`evreetal.2005). Together
shift evolution. Furthermore, we are able to derive the with these studies based on spectroscopic measurements
sSFR(z)fromtheformer. Theresultsaredetailedinthe oftheHαemissionline,ourresultsagreewithastrongly
following sections. increasing EW(Hα) up to z ∼ 2.5, proportional to (1+
z)1.8 (see also Fumagalli et al. 2012; Sobral et al. 2014).
4.1. The EW(Hα) out to z ∼6 This changes at higher redshifts, where we find that
the EW(Hα) is evolving less steep than expected by
Figure 5 shows the redshift evolution of EW(Hα) for
the extrapolation from lower redshifts. This result is
each of the three redshift windows at 1.0 < z < 2.9,
in good agreement with the recent study at z ∼ 4.5
1 < z < 4, and 3 < z < 6 (color-coded bands in blue,
(Marmol-Queralto et al. 2015; Smit et al. 2015a), based
green, and red). We overlay the results from the four
onsmallersamplesbutsimilargalaxyproperties. There-
combinations of metallicity and SFH to show how our
sultsofotherstudies(Shimetal.2011;Starketal.2013;
choice of intrinsic galaxy properties affects the results.
Asexpected,thedifferencesarenegligible,verifyingthat Schenker et al. 2013; Rasappu et al. 2015)10 show larger
the observed color is mainly driven by the contribution EW(Hα) on average, which we think is due to sample
of emission lines (and dust), see Figure 3. The points selection. On one hand, these galaxies are found in the
color coded in the same way show EW(Hα) in the red- fainttailoftheluminositydistributionandstellarmasses
shift ranges where it can be directly measured (see also quoted for these galaxies are 0.5dex or more lower than
Figure 2). At z > 3, we also show the results for a inoursample. Ontheotherhand,inordertobespectro-
sampleofgalaxiesselectedbyphotometricredshifts(red
squares). TheEWsareconsistentwithourspectroscopic 10 These studies do include [SII] and/or [NII] in their mea-
surements of Hα. We correct, if necessary, the contribution from
sample, suggesting that it is not severely biased towards
[SII]and[NII]byassumingaconstantfactorof15%,ifboth,and
young galaxies with enhanced star formation (see also 5%,ifonly[NII],only.
8 A. L. Faisst, et al.
T his work l Spec−z
l 1 .0 < z < 2.9 <logM> ~ 9.8 Photo−z
1000 ll 1 .0 < z < 4.0
l
l 3 .0 < z < 6.0
l
l
l
l
1.8 l
+z) l
1
) ) ~ ( l
aEW(H 100 E W(aH ll l EW(Ha ) ~ (1+z)1.3
l
Literature at similar stellar mass
SDSS Shim+11
VVDS (Lamareille+09) Stark+13
3D−HST (Fumagalli+12) l Rasappu+15
l FMOS (Silverman+15) Marmol+15
Erb+06 Smit+15
Schenker+13
10
0 1 2 3 4 5 6
redshift
Figure 5. Mean rest-frame EW(Hα) as a function of redshift. Our results in the three different redshift windows are shown in blue
(1.0<z<2.9),green(1.0<z<4.0),andred(3.0<z<6.0). ThedifferentbandsshowthefourcombinationsofmetallicityandSFHsfor
eachwindow(seetext). ThecoloredpointsshowtheredshiftswheretheHαlineisdirectlyaccessible. Theredsquaresshowthesamefora
samplebasedonphotometricredshiftsatz>3. Thesymbols(seelegend)showdifferentstudiesmeasuringEW(Hα)directlyfromspectra
(Erb et al. 2006; Lamareille et al. 2009; Fumagalli et al. 2012; Silverman et al., submitted) or from observed color or SED fitting (Shim
etal.2011;Schenkeretal.2013;Starketal.2013;Rasappuetal.2015;Marmol-Queraltoetal.2015;Smitetal.2015a). Tohomogenize
theresults,weapplyaconstantfactorof15%(5%)totheliteraturemeasurementstocorrectforthe[NII]and[SII]([NII]only)emission
lineswherenecessary.
scopicallyconfirmed,thesecontinuum-faintgalaxieshave modelsbracketingdifferentSFHs(exponentiallyincreas-
to be (strongly) Lyα emitting and therefore young with ing with τ = 5 × 108 yr and constant). These models
highstarformationasitisexpectedthatEW(Hα)ispos- arestellarmassnormalizedandwecanthereforedirectly
itivelycorrelatedwithageandSFR(e.g.,Leithereretal. convert EW(Hα) into a specific Hα luminosity without
1999; Cowie et al. 2011). Based on our minimally biased any additional measurement of stellar mass. We then
sample (see Section 5), we find that the evolution of the use the Kennicutt & Evans (2012) relation (assuming a
Hα EW is best parametrized by EW(Hα)∝(1+z)1.3 at ChabrierIMF)toconvertthespecificHαluminosityinto
z (cid:38)2.5. aspecificSFR.Figure6showstheresultingsSFR(z)de-
rived from our EW(Hα) evolution with redshift along
4.2. The sSFR at z >4
with various measurement from the literature. The or-
The Hα is a tracer for star-formation and the stellar angeshadedband(andsolidorangeline)showsthesSFR
continuumred-wardof4000˚Aisagoodtracerforstellar derivedbasedontheexponentiallyincreasingSFH,while
mass. Therefore, the Hα EW is directly proportional to the hatched band shows the case of a constant SFH. For
thesSFRofagalaxywiththenormalizationfactorsolely both we assume an age evolution corresponding to the
depending on its internal properties such as metallicity, cosmic time elapsed since z = 11. The case for a con-
ageofstellarpopulations,andSFH(e.g.,Leithereretal. stantageof500Myr(andexponentiallyincreasingSFH)
1999; Cowie et al. 2011). The ensemble approach also is shown as dashed orange line.
allows a clear determination of the average and range of Since directly calculated from the EW(Hα), the sSFR
these properties and allows their propagation (forward- evolution with redshift is different at low and high red-
modeling) to the final results. This is one big advantage shifts. While we find a strong increase of sSFR propor-
over a model based on “galaxy-by-galaxy” fitting. Fur- tional to (1 + z)2.4 at z (cid:46) 2.2, this flattens out to a
thermore, remember that our results are mostly insen- redshift dependence of (1+z)1.5 at z (cid:38)2.2.
sitive to SFH, age, and metallicities at redshifts z > 4 For comparison, the symbols show various measure-
where the age of the universe is less than 1 Gyr (see ments from the literature, which are summarized in Ta-
Figure 3). ble 2 (low redshift) and Table 3 (high redshift). These
In order to convert the EW(Hα) to sSFR, we use measurements can be broadly split into three groups:
Bruzual & Charlot (2003) composite stellar population
Coherent study of emission lines and sSFR at 3<z <6 9
sSFR from EW(Ha )
<logM> ~ 9.8
age = T(z) − T(11)
age = 500 Myrs
SFH
10 ex p inc. (t = 5e8 yr)
co n st.
] l
1
-r l l l
y
G l l
[ l l
R
l
F
S l
s 1 Speagle+14
l (1+z) fit
l time fit
l l Karim+11 l Reddy+12 Steinhardt+14
Magdis+10 Bouwens+12 deBarros+14
Daddi+07 Stark+09 l Gonzalez+14
Marmol+15 Gonzalez+10 VUDS (Tasca+15)
l
Noeske+07 Stark+13; const. EW Jiang+15 (young/old)
Dunne+09 Stark+13; evolv. EW
0.1
0 1 2 3 4 5 6 7 8
redshift
Figure 6. EvolutionofsSFRasafunctionofredshiftcomputedfromEW(Hα)(z)intherange0.5<z<8.0(orange,extrapolatedbelow
z = 0.5 and above z = 6). The orange band assumes an exponentially increasing SFH (τ = 5×108 yr) and evolving age (Hubble time
sincez=11). Thewidthofthebandinincludesarangeinmetallicity(0.2−1.0Z(cid:12))andtherangeinEW(Hα). Theorangedashedline
showsthecaseforafixesageof500Myrs. ThehatchedbandshowsthesameforaconstantSFH.Alongwiththis,weshowobservations
at low-z (Daddi et al. 2007; Noeske et al. 2007; Dunne et al. 2009; Magdis et al. 2010; Karim et al. 2011; Reddy et al. 2012), and high-z
without emission line correction (Stark et al. 2009; Gonza´lez et al. 2010; Bouwens et al. 2012a) and with emission line correction (Stark
etal.2013;Gonz´alezetal.2014;deBarrosetal.2014;Steinhardtetal.2014;Tascaetal.2015;Jiangetal.2015;Marmol-Queraltoetal.
2015). The dotted (dot-dashed) line shows the fit from Speagle et al. (2014) parametrized in redshift (time) space. We find the sSFR to
beproportionalto(1+z)2.4 atz<2.2andproportionalto(1+z)1.5 athigherredshifts,indicativeofaflattening.
(i) Measurements at z (cid:46) 3, which are based on reli- the ratio of [OIII] to Hα in three discrete redshift ranges
able SFR indicators in the far-IR or sub-mm and with- centered on z ∼2.2, z ∼3.3 and z ∼5.5 (see Figure 2).
out the problem of emission line contamination (Daddi Weareabletofitthe[OIII]/Hαratioreliablyatz ∼2.2
etal.2007;Noeskeetal.2007;Dunneetal.2009;Magdis andz ∼3.3, sincetheredshiftrangeatwhichthebroad-
et al. 2010; Karim et al. 2011; Reddy et al. 2012), (ii) band filters include Hα and [OIII] is fully covered by our
measurements at z >3 that are based on SFR from UV data (see Figure 2 and Figure 4). The right panel of
andSEDfittingaswellasstellarmassestimatesnot cor- Figure 7 shows our [OIII]λ5007/Hα ratio11 at z ∼ 2.2
rectedforemissionlines(Starketal.2009;Gonz´alezetal. and z ∼ 3.3 in blue and green, respectively, along with
2010; Bouwens et al. 2012a), and, (iii) measurements at spectroscopically determined ratios (Steidel et al. 2014;
z > 3 that are based on SFR from UV and SED fitting Sanders et al. 2015), which we find to be in excellent
as well as stellar mass estimates corrected for emission agreement with our measurements at z ∼ 2.2. The dif-
lines (Stark et al. 2013; Gonz´alez et al. 2014; de Bar- ferentbandsforeachcoloragainshowthefourcombina-
ros et al. 2014; Steinhardt et al. 2014; Jiang et al. 2015; tions of intrinsic properties in our model.
Tasca et al. 2015; Marmol-Queralto et al. 2015). Atz ∼5.5ourdataonlyincludesgalaxiesuptoz ∼5.8
and therefore does not cover the full redshift range (i.e.,
4.3. The [OIII]/Hα ratio at z ∼6
entryandexitof[OIII]inthe3.6µmband)thatisneeded
The ratio between Hα and Hβ solely depends on the to reliably constrain the [OIII]/Hα ratio. Furthermore,
dust attenuation. Since we do fit Hα and dust attenu-
ation, we can directly compute Hβ and thus separate it 11 Wesplitthe[OIII]doublet(4960˚Aand5007˚A)byassuming
from the [OIII] line. This allows us to directly measure [OIII]λ5007=1/3×[OIII]λ4960.
10 A. L. Faisst, et al.
2.0
data z ~ 0.2 LAEs (Cowie+11) 1.4
1.5 l mmooddeell ((wnoit ehm eimssiisosnio lnin leins)es) ) 00..68 ll zG ~re 0en.8 PUeSaEs L(Cs a(Hrdua+m0o9n)e+09) z = 3.3 z = 5.5 1.2 )
1ll.0 ll l aH 1 bH
lllllllllllllllllmmll[3.6m] − [4.5m]lllllllllllllllll−−−llllllllllllll11000llll.....ll50505llllllllllllllllllllllllllllllnlllloll lcolnlllltalmllllinllllatilollnlllll < ll2lllll'll'lllllllllllllllllllllllllllllllllllll lla00000g([OIII]/H).....02457 llog ([OIII]5007/ −−−000000......642024 ldllSiDstSriSb luotlcioanl z = 2.2 l SWWFMKDMBOIISSSSOSSPPSDS − (zz zMEz~~<~F11O01 z.((.S3MC5~F) 2o(IeS.Rl3hbi Elte(vaS r(e+taSr+1ntm15ed3i)aed)nres+l++111554))) −00000....02468.2llog ([OIII]5007/
l l no contamination < 3'' lo
−2.0 0 1 2 3 4 5 6
5.0 5.2 5.4 5.6 5.8 6.0
redshift
redshift
ll
Figure 7. Left: Becauseourdatadoesnothavethefullredshiftrangetoreliablymeasurethe[OIII]/Hαratioatz∼5.5,weshowhere
thescatterofthismeasurement. Thetopredlinewithlog([OIII]/Hα)=0.0showsthebest-fitmodel. Theotherlinesshowmodelswith
increased [OIII]/Hα ratios. We find log([OIII]/Hα) ∼ 0.0−0.7 at z ∼ 5.5. The data is shown in blue and the model without emission
lines is shown as red dashed line. Right: Mean dust corrected [OIII]/Hα ratio as a function of redshift. All the samples are matched in
stellarmassandthe[OIII]/Hβ ratioiscomputedassumingcaseB.Ourestimatesbasedonbroad-bandcolorsareshowninblue(z∼2.2),
green (z ∼ 3.3), and red (z ∼ 5.5) including their uncertainties (from observation and different models). Spectroscopic measurement at
lowerredshiftsareshownwithsymbols(Cowieetal.2011;Colbertetal.2013;Steideletal.2014;Mehtaetal.2015;Sandersetal.2015;
Silvermanetal.,submitted)andthedistributionofSDSSgalaxiesatz∼0isrepresentedbytheblackdashedline(opencircle: median).
All in all, we find a progressively increasing [OIII]/Hα ratio over the redshift range z ∼ 2−6. We also show high-z analogs as “Green
Peas”(greenfilledcircle,Cardamoneetal.2009)andUSELsatz∼0.8(orangeopencircle,Huetal.2009)forcomparison.
thesparsesamplingofdataattheseredshiftscontributes high-z analogs” is further discussed in Section 5.4.
to the uncertainty. Also the addition of galaxies with
photometric galaxies does not increase the sample size 5. DISCUSSION
by much at z >5.8 as shown in Appendix A. We there- We use a sample of >500 spectroscopically confirmed
forediscussthiscaseinmoredetailinthefollowing. The galaxies to derive the Hα EW, the sSFR(z), and the
leftpanelinFigure7showstheredshiftrangefromwhich [OIII]λ5007/Hα ratio at 3 < z < 6. The main idea of
we determine the [OIII]/Hα ratio at z ∼5.5. As before, our analysis is to base these measurements on primary
the points show the data and the blue band shows their observables (the observed color vs. redshift evolution)
scatter. Themodelwithoutemissionlines(butincluding and to minimize the model uncertainties. The forward
dust) is shown as red dashed line. Note, that at these modelingapproachbasedontheensembleinsteadofsin-
redshift, where the stellar ages are < 1Gyr, the model glegalaxiesallowsustomarginalizeoverarangeofSFH,
is very insensitive to SFH, age, and metallicity and pri- metallicities,andagesandtopropagatetheuncertainties
marily depends on the [OIII]/Hα ratio in this case. We in these quantities to the final result. Due to the young
show the best fit [OIII]/Hα ratio (log([OIII]/Hα) ∼ 0.0) age of the universe of less than 1 Gyr at z > 4, the dif-
along with four different models with increasing ratios ferences between different assumptions that go into the
in red. This large variation over a range of 0.7 dex in modeling of the galaxy population are less significant,
[OIII]/Hα indicates the existence of galaxies with very leading to robust results at these high redshifts (see also
strong [OIII] emission at these redshifts, in agreement Figure 3).
with the recent findings at z ∼ 6.7 using a similar tech- Intheprevioussection,wehaveestablishedthefollow-
nique (e.g., Roberts-Borsani et al. 2015). The galaxies ing results.
clusteringaroundz ∼5.65areLyαnarrowbandselected
and thus preferentially young and highly star forming. 1. The EW(Hα) increases continuously as (1+z)1.8
This could be the reason for their high [OIII]/Hα ratios. up to z ∼ 2.5 and flattens off at higher redshifts
The range of [OIII]/Hα ratios at z ∼5.5 is shown in red with a redshift proportionality of (1+z)1.3.
on the right panel of Figure 7.
We find a progressively increasing [OIII]λ5007/Hα ra- 2. The sSFR increases proportional to (1+z)2.4 at
tiowithredshiftatz >2,oncethelargeup-wardscatter z (cid:46)2.2 but shows a less strong evolution at higher
at z ∼6 is taken into account. On the other hand, there redshifts proportional to (1+z)1.5.
is not much evolution between z ∼ 1 and z ∼ 2 (liter-
ature at z ∼ 1 and z ∼ 1.5, Colbert et al. 2013; Mehta 3. We find a best-fit [OIII]/Hα ratio of z ∼ 6 star-
forming galaxies on the order of unity (similar to
etal.2015;Silvermanetal.,submitted). Theaverageline
z =2 and z =3 galaxies), however, with a scatter
ratio of local (z (cid:46)0.3) SDSS galaxies is ∼0.2 dex lower
uptoaratiooffive. Thissuggeststheprogressively
than at z = 2 and ∼ 0.2−0.8 dex lower than at z ∼ 6.
However,thedistributionof[OIII]λ5007/Hαratiosinlo- increasing [OIII]/Hα ratios at z >3.
cal galaxies (shown by the dashed density-histogram) is
Before proceeding to the discussion of these results,
broadandsub-samplesofthesegalaxiesshowsimilarline
we have to make sure that our sample is only minimally
ratios as high-z galaxies. The potential of such “local
biased.
