# Spectral Energy Distributions of low-luminosity radio galaxies at z$\sim
1-3$:
a high-zview of the host/AGN connection.

Ranieri D. Baldi11affiliation: SISSA, via Bonomea 265, 34136 Trieste, Italy
22affiliation: Space Telescope Science Institute, 3700 San Martin Drive,
Baltimore, MD 21218, USA , Marco Chiaberge22affiliation: Space Telescope
Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA
33affiliation: INAF-Istituto di Radio Astronomia, via P. Gobetti 101, I-40129
Bologna, Italy 44affiliation: Center for Astrophysical Sciences, Johns Hopkins
University, 3400 N. Charles Street Baltimore, MD 21218 , Alessandro
Capetti55affiliation: INAF-Osservatorio Astronomico di Torino, Strada
Osservatorio 20, 10025 Pino Torinese, Italy , Javier Rodriguez-
Zaurin66affiliation: Insituto de Astrofisica de Canarias, Via Lactea s/n, La
Laguna 38200, Spain; Departamento de Astrofisica, Universidad de La Laguna, La
Laguna 38206, Spain 22affiliation: Space Telescope Science Institute, 3700 San
Martin Drive, Baltimore, MD 21218, USA , Susana Deustua22affiliation: Space
Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA
,William B. Sparks22affiliation: Space Telescope Science Institute, 3700 San
Martin Drive, Baltimore, MD 21218, USA<EMAIL_ADDRESS>
###### Abstract

We study the Spectral Energy Distributions, SEDs, (from FUV to MIR bands) of
the first sizeable sample of 34 low-luminosity radio galaxies at high
redshifts, selected in the COSMOS field. To model the SEDs we use two
different template-fitting techniques: i) the Hyperz code that only considers
single stellar templates and ii) our own developed technique 2SPD that also
includes the contribution from a young stellar population and dust emission.
The resulting photometric redshifts range from z$\sim$0.7 to 3 and are in
substantial agreement with measurements from earlier work, but significantly
more accurate. The SED of most objects is consistent with a dominant
contribution from an old stellar population with an age $\sim 1-3\times$ 109
years. The inferred total stellar mass range is $\sim$10${}^{10}-10^{12}$ M⊙.
Dust emission is needed to account for the 24 $\mu m$ emission in 15 objects.
Estimates of the dust luminosity yield values in the range $L_{\rm dust}\sim
10^{43.5}-10^{45.5}$ erg s-1. The global dust temperature, crudely estimated
for the sources with a MIR excess, is $\sim$ 300-850 K. A UV excess is often
observed with a luminosity in the range $\sim 10^{42}-10^{44}$ erg s-1 at 2000
Å rest frame.

Our results show that the hosts of these high-z low-luminosity radio sources
are old massive galaxies, similarly to the local FRIs. However, the UV and MIR
excesses indicate the possible significant contribution from star formation
and/or nuclear activity in such bands, not seen in low-z FRIs. Our sources
display a wide variety of properties: from possible quasars at the highest
luminosities, to low-luminosity old galaxies.

###### Subject headings:

Galaxies: active – galaxies: high-redshift – Galaxies: nuclei – Galaxies:
evolution – Galaxies: photometry – Infrared: galaxies

##  1\. Introduction

The search for radio-loud Active Galactic Nuclei (AGN) is one of the widely
used tools to study the distant universe (z $>$ 1). Indeed the first objects
found at z $>$ 1 were radio-loud quasars (see Stern & Spinrad 1999 and
references therein) and a large number of radio galaxies ($\sim$300) are known
to exist at high redshifts (e.g., Miley & De Breuck 2008).

Flux-limited samples of radio galaxies such as the 3CR and its deeper
successors 6C and 7C catalogs are affected by a tight redshift-luminosity
correlation. This, alongside with the steep luminosity function of radio
sources, gives rise to a selection bias which results in the presence of high
power sources predominantly at high redshifts only and low power sources
exclusively at low redshift. Therefore, our knowledge of radio galaxies and
their hosts at high redshifts is exclusively based on studies of powerful
(edge-brightened, Fanaroff & Riley 1974) FR IIs. Thus, the properties of the
low-luminosity AGN population in the early Universe is essentially unknown.
The missing pieces of the puzzle might be obtained by studying a sample of low
luminosity (’edge-darkened’, Fanaroff & Riley 1974) FR Is at high redshifts.
The lower power of these objects with respect to that of high-z FR IIs might
allow us to analyze with more clarity the properties of the hosts at high
redshifts. Furthermore, in terms of host galaxy properties, low power sources
are both more abundant and most likely more similar to quiescent galaxies (in
terms of host galaxy properties) than FR IIs.

Only a few FR Is at high redshift are present in the 7C sample (Heywood et
al., 2007) and two were possibly found in HDF North (Snellen & Best, 2001).
Chiaberge et al. (2009) obtained the first sizeable sample of 37 FR I
candidates at z $\gtrsim 1$, located in the 2 deg2 area of the sky observed by
the COSMOS survey (Scoville et al., 2007). In this paper we perform a detailed
analysis of the photometric properties of the host galaxies of these radio
sources, by taking advantage of the large multi-wavelength coverage provided
by the COSMOS collaboration. This allows us to derive their spectral energy
distributions (SEDs) from $\sim$ 0.15 $\mu$m to 24 $\mu$m. We then model the
obtained SEDs with synthetic stellar population templates with the aim of both
inferring the properties of each galaxy and of deriving an accurate estimate
of the photometric redshifts.

Mobasher et al. (2007) and Ilbert et al. (2009) performed a similar study with
the aim of determining the photo-z for the whole COSMOS sample. However, as
already noted by Chiaberge et al. (2009), the faint counterparts of these
radiogalaxies might be easily misidentified or be missed from the COSMOS
catalog. Therefore a careful object-by-object study of the sample is necessary
to correctly identify the genuine counterparts of the radio sources at all
wavelengths. This allows us to obtain reliable SEDs and perform a robust study
of the hosts of these objects. In this work, using the multi-band data
provided by the COSMOS survey (described in Section 3), we obtain the
photometric data from the catalog checking visually the correct identification
of each object of the sample (Section 4). In Section 5 we describe the codes
used to model the SEDs, Hyperz and 2SPD. We present the results obtained from
the SED modeling in Section 6: the photometric redshifts, the radio power
distribution, and the host properties, such as stellar ages, masses and dust
and UV components. In Section 7 we summarize the results and we discuss our
preliminary findings.

We adopt a Hubble constant of H0 = 71 km s-1 Mpc-1, $\Omega_{m}$ = 0.27 and
$\Omega_{vac}$ = 0.73, as given by the WMAP cosmology (e.g., Spergel et al.
2003; Jarosik et al. 2011). All the magnitudes are in AB mag system, if not
otherwise specified.

##  2\. Sample

Chiaberge et al. (2009) selected a sample of 37 high-z low-power radio
galaxies at z $\sim$ 1-2 in the Cosmic Evolution Survey (COSMOS) field
(Scoville et al., 2007). This is the first sizeable sample of low-power radio
sources at mid- to high-redshifts. This sample has been obtained with a four-
steps multi-wavelength selection process, which is described in detail in
Chiaberge et al. (2009). Here we briefly summarize the selection procedure
which depends on the following assumptions: i) the FR I/FR II break luminosity
at 1.4 GHz does not change with redshift, and ii) the host properties of
distant low-luminosity sources are similar to those of high-power FR IIs in
the same redshift range.

The first step consists of selecting radio sources from the FIRST catalog
(Becker et al., 1995) whose 1.4 GHz total flux corresponds to the luminosities
expected by FR Is at 1$<z<$2\. The second step is based on a radio
morphological classification: the radio sources which show clear “edge-
brightened” structures are rejected in order to exclude the bona-fide FR IIs
from the sample. The third step implies the optical identification of the
radio sources in the COSMOS optical images. The objects associated with host
galaxies brighter than $i$ $<$ 22 were rejected to exclude lower redshifts
starburst galaxies. As a final step, u-band dropouts are also rejected since
these are objects most likely located at redshift higher than z=2. The
resulting sample consists of 37 FR I candidates. A posteriori, the photometric
redshifts range of most of them is between $\sim$1 and $\sim$2 (Mobasher et
al., 2007; Ilbert et al., 2009), with the exception of 3 objects (namely, 7,
27, and 66111The ’naming convention’ of the sources used in the paper is based
on the listing number in the COSMOS catalog.) that we exclude from any further
analysis because out of the redshift range of interest. We are then left with
a final sample of 34 objects. The list of the radio sources is presented in
Table 1.

Table 1Radio positions, redshifts and radio flux of the sample ID | RA | DEC | zphot,Ilbert | zspec | FNVSS | FFIRST
---|---|---|---|---|---|---
1 | 150.20744 | 2.2818749 | 0.92${}^{+0.02}_{-0.06}$ | 0.8827a-0.8823b |  $<$2.5 | 1.79
2 | 150.46751 | 2.7598829 |  |  | 2.6 | 1.08
3 | 150.00253 | 2.2586310 | 1.96${}^{+0.36}_{-0.41}$ |  | 5.2 | 4.21
4 | 149.99153 | 2.3027799 | 1.45${}^{+0.08}_{-0.15}$ |  | 7.5 | 5.99
5 | 150.10612 | 2.0144780 | 1.84${}^{+0.22}_{-0.12}$ |  | 3.4 | 1.30
11 | 150.07816 | 1.8985500 | 1.31${}^{+0.60}_{-0.24}$ |  |  $<$2.5 | 1.13
13 | 149.97784 | 2.5042069 | 1.09${}^{+0.06}_{-0.06}$ |  | 2.4 | 1.51
16 | 150.53772 | 2.2673550 | 0.95${}^{+0.05}_{-0.02}$ | 0.9687b | 4.4 | 5.70
18 | 149.69325 | 2.2674670 | 0.92${}^{+0.01}_{-0.09}$ |  | 5.1 | 4.39
20 | 149.83209 | 2.5695460 | 1.00${}^{+0.02}_{-0.05}$ |  |  $<$2.5 | 1.33
22 | 149.89508 | 2.6292144 |  |  |  $<$2.5 | 2.74
25 | 150.45673 | 2.5597000 | 1.12${}^{+0.29}_{-0.03}$ | 0.7917b∗ | 2.7 | 2.18
26 | 149.62114 | 2.0919881 | 1.20${}^{+0.06}_{-0.02}$ |  | 3.2 | 1.88
28 | 149.60064 | 2.0918673 |  |  | 2.4 | 1.77
29 | 149.64587 | 1.9529760 | 1.59${}^{+0.45}_{-0.30}$ |  | 2.3 | 2.12
30 | 149.61542 | 1.9910541 | 0.90${}^{+0.31}_{-0.03}$ |  | 2.4 | 1.26
31 | 149.61916 | 1.9163600 | 0.91${}^{+0.02}_{-0.05}$ | 0.9132a-0.9123b | 4.1 | 3.71
32 | 149.66830 | 1.8379777 |  |  | 3.1 | 1.31
34 | 150.56023 | 2.5861051 | 1.42${}^{+0.65}_{-0.32}$ |  | 4.5 | 5.25
36 | 150.55662 | 1.7913361 |  |  | 3.3 | 3.19
37 | 150.74336 | 2.1705379 | 1.27${}^{+0.09}_{-0.02}$ |  | 2.2 | 1.87
38 | 150.53645 | 2.6842549 | 1.12${}^{+0.10}_{-0.05}$ |  | 11.6 | 10.01
39 | 149.95804 | 2.8288901 | 1.08${}^{+0.03}_{-0.03}$ |  |  $<$2.5 | 1.37
52 | 149.90590 | 2.3964710 | 0.74${}^{+0.02}_{-0.03}$ | 0.7417b |  $<$2.5 | 1.54
70 | 150.61987 | 2.2894360 | 2.21${}^{+0.57}_{-0.37}$ |  | 4.5 | 3.90
202 | 149.99506 | 1.6324950 | 1.19${}^{+0.24}_{-0.14}$ |  | 3.3 | 1.08
219 | 150.06444 | 2.8754051 | 1.04${}^{+0.01}_{-0.09}$ |  |  $<$2.5 | 1.85
224 | 150.28999 | 1.5408180 | 1.10${}^{+0.07}_{-0.03}$ |  | 3.2 | 3.31
226 | 150.43864 | 1.5934480 | 1.76${}^{+0.61}_{-0.14}$ |  |  $<$2.5 | 1.19
228 | 149.49455 | 2.5052481 | 1.89${}^{+0.74}_{-0.59}$ |  | 3.7 | 2.04
234 | 150.78925 | 2.4539680 | 1.10${}^{+0.12}_{-0.03}$ |  | 5.2 | 4.43
236 | 150.70554 | 2.6296339 |  | 2.132c | 7.0 | 7.10
258 | 149.55934 | 1.6310670 | 0.90${}^{+0.02}_{-0.02}$ | 0.9009a | 3.7 | 2.24
285 | 150.72131 | 1.5823840 | 1.21${}^{+0.06}_{-0.08}$ |  | 3.5 | 2.95

##  3\. Data

###  3.1. UV, optical, and IR photometric data

The photometric data used to build the SEDs of our sources are taken from the
COSMOS survey (Scoville et al., 2007). The survey comprises ground based as
well as imaging and spectroscopic observations from radio to X-rays
wavelengths, covering a 2 deg2 area. Given the high sensitivity and resolution
of these data, COSMOS provides samples of high redshift objects with greatly
reduced cosmic variance as compared to earlier surveys.

Ground-based UV, optical, and IR observations and data reduction are presented
in Capak et al. (2007), Capak et al. (2008) and Taniguchi et al. (2007). A
multiwavelength photometric catalog was generated using SExtractor (Bertin &
Arnouts, 1996). The COSMOS catalog is derived from a combination of the CFHT
$i^{*}$ and Subaru $i^{+}$ images, to which the authors refer as ’I-band
images. The catalog includes objects with total (”mag-auto”) I $<$ 25 and
searches for counterparts in a radius of 1″around the I-band detection. At
fainter magnitudes the catalog begins to be incomplete and have more spurious
detections, and photometric redshifts are poorly constrained.

For this study we use the COSMOS Intermediate and Broad Band Photometry
Catalog 2008 (Capak et al., 2008)222http://irsa.ipac.caltech.edu/cgi-
bin/Gator/nph-dd which provides the multiwavelength magnitudes of our sources
from FUV to K bands. Narrow band filters are not considered due to the
possible strong contamination of emission lines from the AGN and their low
signal-to-noise ratio.

HST (Koekemoer et al., 2007) and Spitzer data (both IRAC and MIPS, Sanders et
al. 2007) are also included in the COSMOS survey. The later is presented in
two separate catalogs. S-COSMOS IRAC 4-channel Photometry Catalog June 2007 is
for IRAC data. S-COSMOS MIPS 24 Photometry Catalog October 2008 and S-COSMOS
MIPS 24 $\mu$m DEEP Photometry Catalog June 2007 are for MIPS data at 24
$\mu$m with two different flux limits, 0.15 and 0.08 mJy, respectively.

For the sake of clarity, we summarize here the available datasets from the
COSMOS survey used in this work (Table 2):

  1. 1.

GALEX: UV data were taken using GALEX (Martin et al., 2005). The observations,
performed as part of the GALEX Deep Imaging Survey, are in the NUV and FUV
bands, respectively with an angular resolution of 5$\farcs$6 (NUV) and
4$\farcs$2 (FUV) (Morrissey et al., 2007).

  2. 2.

HST: the COSMOS HST data are single-orbit F814W ACS images. They have the
highest angular resolution ($\sim 0\farcs 09$, Koekemoer et al. 2007) among
the COSMOS images.

  3. 3.

SUBARU: The Suprime-Cam instrument on the Subaru telescope observed the COSMOS
field in six broad bands ($B_{J}$, $g^{+}$, $V_{J}$, $r^{+}$, $i^{+}$,
$z^{+}$) with an angular resolution of $\sim$0$\farcs$2 (Taniguchi et al.,
2007).

  4. 4.

CFHT: the Canada-France-Hawaii Telescope (CFHT) provides $u^{*}$ and $i^{*}$
images using Megacam (Boulade et al., 2003) and $K$ images with Wircam.

  5. 5.

UKIRT: near infrared Wide Field camera (WFCAM) on the United Kingdom Infrared
Telescope (UKIRT) provides the $J$-band images (Casali et al., 2007).

  6. 6.

NOAO: $K_{S}$ data are taken at the Kitt Peak National Observatory (KPNO)
telescope with FLAMINGOS and The Cerro Tololo Inter-American Observatory
(CTIO) telescope (Capak et al., 2007). These telescopes belong to the National
Optical Astronomy Observatory (NOAO).

  7. 7.

Spitzer: Spitzer cycle-2 S-COSMOS is an infrared imaging survey of the COSMOS
field (Sanders et al., 2007). They obtained observations with the IRAC camera
in four channels, at 3.6, 4.5, 5.6, and 8 $\mu$m, and with MIPS in the 24, 70,
160 $\mu$m band. We only consider the data at 24 $\mu$m, since no object is
detected at longer wavelengths (with the exception of object 37 which is also
detected at 70 $\mu$m).

Table 2COSMOS broad bands and their properties. Filter | Telescope | $\lambda_{eff}$ | FWHM | sensitivity
---|---|---|---|---
$FUV$ | GALEX | 1538.6Å | 230.8Å | 25.7
$NUV$ | GALEX | 2315.7Å | 789.1Å | 26.0
$u^{*}$ | CFHT | 3911.0Å | 538.0Å | 26.5
$B_{J}$ | Subaru | 4439.6Å | 806.7Å | 27.0
$g^{+}$ | Subaru | 4728.3Å | 1162.9Å | 27.0
$V_{J}$ | Subaru | 5448.9Å | 934.8Å | 26.6
$r^{+}$ | Subaru | 6231.8Å | 1348.8Å | 26.8
$i^{*}$ | CFHT | 7628.9Å | 1460.0Å | 24.0
$i^{+}$ | Subaru | 7629.1Å | 1489.4Å | 26.2
$F814W$ | HST | 8037.2Å | 1539.0Å | 27.2
$z^{+}$ | Subaru | 9021.6Å | 9021.6Å | 25.2
$J$ | UKIRT | 12444.1Å | 1558.0Å | 23.7
$K_{S}$ | NOAO | 21434.8Å | 3115.0Å | 21.6
$K$ | CFHT | 21480.2Å | 3250.0Å | 23.7
IRAC1 | Spitzer | 35262.5Å | 7412.0Å | 23.9
IRAC2 | Spitzer | 44606.7Å | 10113.0Å | 23.3
IRAC3 | Spitzer | 56764.4Å | 13499.0Å | 21.3
IRAC4 | Spitzer | 77030.1Å | 28397.0Å | 21.0
MIPS1 | Spitzer | 23.68$\mu$m | 4.7$\mu$m | 29.6

###  3.2. Radio data

Chiaberge et al. (2009) selected the sample in the COSMOS field using radio
fluxes at 1.4 GHz from the FIRST survey (Becker et al., 1995). The data are
obtained with the VLA in B configuration with an angular resolution of
$\sim$5″and reach a flux limit of $\sim$1 mJy and are listed in Chiaberge et
al. (2009).

In addition, the NVSS survey (Condon et al., 1998) provides 1.4 GHz radio data
for our sample, but at lower resolution ($\sim$45 ″) and with a higher flux
density limit ($\sim$2.5 mJy) than the FIRST survey. These differences imply
that seven of our objects are missed in the NVSS catalog. Nonetheless, the
NVSS data are useful since they are more sensitive to diffuse low surface
brightness radio emission that the FIRST data. The radio fluxes are taken from
the NVSS archive333http://www.cv.nrao.edu/nvss/NVSSlist.shtml. A NVSS/FIRST
comparison indicates that the flux ratio between the two catalogs is usually
between $\sim 1$ and 2, and never larger than $\sim$ 3\. Table 1 shows the
NVSS and FIRST radio fluxes of the objects.

###  3.3. Spectroscopy

The COSMOS survey provides spectroscopic data from the Very Large Telescope
(VLT) (zCOSMOS, Lilly et al. 2007) and from the Magellan (Baade) telescope
(Trump et al., 2007)

zCOSMOS is a large-area redshift survey which consists of two parts. The
first, namely zCOSMOS-bright, considers a magnitude-limited (I-band mag
$<$22.5) sample of $\sim$20,000 galaxies located within the central 1 deg2.
Spectra cover the wavelength range 5500 Å $<\lambda<$ 9000 Å. The second,
namely zCOSMOS-deep, is an ongoing survey (not yet public) of $\sim$10,000
blue galaxies in the same field filtered with a color selection to be in the
range of 1.4 $<$ z $<$ 3.0. The spectra cover the wavelength range of 3600 Å
$<\lambda<$ 6800 Å.

The Magellan survey presents spectroscopic redshifts for the first 466 X-ray
and radio-selected AGN candidates. The wavelength coverage of these spectra is
$\sim$ 5500-9200Å. Their redshift yield is 72% for iAB $<$ 24 and $>$90% for
i${}_{AB}<$22.

In this work we use only spectroscopic measurements with a confidence level
greater than 99%. Six objects of the 34 objects are included in the
spectroscopic surveys described above with the required quality (Table 1).

##  4\. Multi-band counterparts identification

The process of counterparts identification of our radio galaxies in the COSMOS
catalog suffers from the limitations typical of any multiband survey (such as
misidentification of targets with close neighbor or the contamination by
nearby bright sources). We then prefer to perform our multi-band counterpart
identification on each source by visually inspecting its multiband images,
rather than blindly use the data provided by the COSMOS catalog.

We start the process looking for a I-band counterpart to the radio source in
the COSMOS catalog by adopting 0$\farcs$3 as search radius. 29 identifications
are found in our sample, most at distances smaller than 0$\farcs$1\. Three
objects (22, 28, and 32) are clearly visible in the I-band images, but they
are not found in the COSMOS catalog, since they are below its detection
threshold. In addition, there are two objects (2 and 36) for which no I-band
counter part is found.

For the I-band detected sources, we search in the COSMOS broadband catalog
that provides photometry (from the FUV to the K band) over an aperture of 3″
diameter. The COSMOS catalog associates its counterparts in the remaining
bands with the brightest and closest sources to its I-band position within a
radius of 1″. We extend this method including also the Spitzer/IRAC and MIPS
catalogs, by using a larger search radius (2″) due to their coarser
resolutions with respect to the COSMOS broadband images. For the 5 sources not
present in the COSMOS catalog (because they are too faint in the I-band) we
note that in all cases a clear counterpart is instead present at longer
wavelength providing us with a robust multiband identification.

Figure 1.— Images of object 28 taken from the COSMOS survey. From upper left
to lower right panels we show the images from Subaru $i^{+}$, Subaru $z^{+}$,
UKIRT J, CFHT K and the four channels of Spitzer/IRAC. The size of the image
is 5$\farcs$5 $\times$ 5$\farcs$5\. We overlap the radio contours (green
lines) at 1.4 GHz from VLA COSMOS data on the frames to show the position of
the radio source.

During the visual check of the individual counterparts, two main problems
emerge in the COSMOS broadband catalog: i) the presence of nearby sources,
within the 3″ radius used for the aperture photometry, contaminating the broad
band measurements of the genuine emission from the radio-galaxies of our
sample; ii) the counterparts to the i-band object does not always correspond
to the same object over the various bands. Object 28 is a clear example of
this situation (see Fig. 1): from the J band to 8$\mu$m the brightest source
in the field is coincident with the radio source, while in the $i^{+}$ and
$z^{+}$ bands this is out-shined by a nearby object, causing an erroneous
identification in the COSMOS catalog across the various bands. Furthermore,
the presence of this neighbor also causes a strong (and dominant short-ward of
the J band) contamination to the genuine flux of the radio-galaxy. Other
examples of contamination and mis-identification are reported in Fig.2.

In order to amend these problems, we first perform a new 3″ aperture
photometry properly centered on the position of the radio source, to isolate
the genuine emission of the counterparts. In case of contamination from nearby
source(s), we subtract from the flux resulting from the photometry centered on
the radio source the emission from the neighbor(s), limiting to the fraction
that falls into the 3″ aperture. When we cannot separate a source from a close
source or when the counterpart is not visible in a given band, we measure an
1-$\sigma$ upper limit to the flux.

At 24 $\mu$m, when the counterpart of a source is not seen at the radio
position, we use the detection limit of MIPS catalogs as upper-limit to the
source flux. However, when a nearby source contaminates the emission of our
target, we prefer to measure the upper-limit directly on the image at the
radio position. In addition, in the NUV/FUV band, if the source is not
detected in GALEX we prefer not to include upper limits in our analysis,
because its corresponding flux is substantially higher than those at larger
wavelengths.

Figure 2.— Left panel: Subaru $i^{+}$ image of object 34 with superposed the
radio contours (green) and the 3″ aperture centered on the radio source
(blue). This is a case in which, in addition to the radio-galaxy, other
objects lie inside the integration aperture of the target, contaminating the
COSMOS catalog flux measurement. Middle and right panels: Subaru $z^{+}$ and
CFHT $u^{*}$ images of object 5 with superposed the radio contours. This is a
case of misidentification of a source with a neighbor included within the 1″
from the target) in different bands. The size of the images is 2″$\times$ 2″.

In some cases the GALEX photometry returns apparently incorrect results. In
Fig. 3 we show as an example the HST and GALEX images of object 13; this
source is labeled as detected by the COSMOS catalog in both the FUV and NUV
GALEX bands, while clearly there is no emission above the background level.
The opposite is seen in MIPS measurements, when clearly visible objects are
missing from the catalog (see Fig. 4). In these cases, we do not consider the
GALEX data, while for MIPS we provide a new estimate to the 24 $\mu$m flux.

We also find an apparent error in the reported UKIRT J-band photometric
points. We obtain our own measurements on a sample of stars present in the
COSMOS field (Wright et al., 2010) using the PSF-matched and PSF-original
J-band images444The PSF-matched images are obtained by convolving each PSF-
homogenized image with a Gaussian kernel that produced the same flux ratio
between a 3″and 10″aperture for a point source to avoid PSF-matching problems
in the multiband catalog. The PSF-original images are the pure images without
any PSF matching. Both the images are provided by the COSMOS catalog. and we
compare them to the magnitudes given by the catalog. This test reveals that
the zero-point mag of the PSF-matched J-band images is higher than the
corresponding value of the PSF-original J-band images, taking into account of
the effect of the PSF convolution. This systematic difference of the zero-
point magnitudes is $\Delta$J = 0.90$\pm$0.06, intriguingly similar to the
offset between Vega and AB mag system for the UKIRT J-band (JAB = JVega \+
0.94 for UKIRT, Hewett et al. 2006). We obviate this problem by performing
3″-aperture photometry on the J-band counterparts on the UKIRT images smoothed
with a Gaussian kernel to a 1.2-1.5″ FWHM to reproduce the effect of the PSF
matching.

Figure 3.— From left to right: HST/ACS F814W, GALEX NUV and FUV bands images
of object 13. The radio source (whose position is marked by the circles at the
images center) does not apparently correspond to any UV emission, contrarily
to the detection found in the catalog. The size of the images is
7$\farcs$5$\times$7$\farcs$5\.

Figure 4.— Images of object 70 from ACS (left) and MIPS 24$\mu$m (right). This
provides an example of a source clearly emitting at 24$\mu$m, but not present
in the COSMOS Spitzer/MIPS catalog. The position of the radio source is
represented by the green radio contours from VLA COSMOS. The size of the
images is 28″$\times$28″.

Once we obtain the correct aperture photometry of the various multi-band
counterparts of the radio galaxies, we apply the appropriate aperture
corrections. While this is not needed for the Subaru, CFHT, and UKIRT, for
GALEX data we apply an aperture correction by multiplying the total flux by a
factor of 0.759 (Capak et al., 2007). Similarly, since we select from the IRAC
catalog the 2$\farcs$9 aperture photometry (because it is closest to the 3″
aperture of the COSMOS broadband catalog among the different optional
apertures), the IRAC aperture flux was converted to total source flux by using
the aperture correction factors (from the IRAC manual) of 1.19, 1.27, 1.48,
and 1.27 for the four channels. The MIPS data instead already include the
aperture correction. Since the CFHT telescope is more sensitive and with
higher resolution than the images from the NOAO telescopes, this usually
results in far smaller errors for the K-band magnitudes. In such cases, we
prefer to use only the CFHT K-band data. The results of the corrected
3″-aperture photometric measurements are tabulated in Tables 6 and 7.

##  5\. SED fitting

The SEDs are derived by collecting multiband data from the FUV to the MIR.
Since not all of the objects are detected in the entire set of available
bands, the number of detections used to constrain the SED fitting ranges from
15 to 19.

The synthetic stellar templates used to model the observed SEDs are the
Bruzual & Charlot (2009) (priv. comm.) and Maraston (2005) templates. These
templates are defined with different Initial Mass Function (IMF) (Salpeter,
1955; Kroupa, 2001; Chabrier, 2003) with solar metallicity (Z⊙ = 0.02). These
libraries contain composite stellar population (CSP) computed with different
star formation histories: a constant star-forming system (with constant star
formation rate of 1 M⊙/yr); a single burst of star formation; and ten
$\mu$-models555These templates correspond to synthetic spectra computed with
exponentially decaying star formation rate, $\psi(t)$ $\propto$ e(-t/τ) where
$\tau$ is the star formation timescale. with time-decays of 0.1, 0.2, 0.3,
0.6, 1.0, 2.0, 3.0, 5.0, 10.0, and 15.0 Gyr. These templates include 221
tracks of ages from 0.1 Myr to 20 Gyr and cover the wavelength range from 91 Å
to 160 $\mu$m. Fig. 5 shows the stellar templates used for SED modeling.

Figure 5.— Examples of SED templates from BC09 (Bruzual & Charlot, 2009) (left
panel) and MA05 (Maraston, 2005) (right panel) in units of L☉ $\AA^{-1}$. The
templates correspond to different ages and different star formation histories.
Arranging the BC09 and MA05 templates with increasing luminosity at 3550 Å
(following the dashed line), the different plotted models correspond to
$\mu$=15 Gyr decaying CSP of 1.1 Myr (red line); single stellar population of
20 Gyr (gray line); $\mu$=0.1 Gyr decaying CSP of 10 Gyr (dark green line);
$\mu$=10 Gyr decaying CSP of 2.8 Myr (light green line); $\mu$=0.3 Gyr
decaying CSP 5.25 Gyr (orange line); $\mu$=5 Gyr decaying CSP of 6.9 Myr (blue
line); $\mu$=3 Gyr decaying CSP of 17 Myr (yellow line); $\mu$=0.6 Gyr
decaying CSP of 1.7 Gyr (pink line); $\mu$=2 Gyr decaying CSP of 38 Myr
(fuchsia line); and $\mu$=1 Gyr decaying CSP of 0.2 Gyr (light blue line). We
invite the reader to download the color version of the plot to distinguish the
different lines.

The Bruzual & Charlot (2009) models differs from those of Maraston (2005)
mainly because of a different recipe for the inclusion of thermally pulsing
asymptotic giant branch (TP-AGB) phase. These stars significantly contribute
to the IR emission for ages higher than $\sim$1 Gyr. The resulting effect is
that the Bruzual & Charlot templates are bluer and give slightly older ages
than Maraston (2005) models.

###  5.1. SED fitting method: Hyperz

We firstly estimate the photometric redshift for our sample by using the
template fitting technique, Hyperz (Version 1.3 , obtained from M. Bolzonella
in priv. comm.), described in detail by Bolzonella et al. (2000). This SED
fitting procedure is based on reproducing the overall shape of the SEDs and
recognizing strong spectral properties, such as the 4000 Å Ca break or Lyman
break at 912 Å. This code consists of a convolution of templates, which
represent the rest-frame SEDs for galaxies with different star formation
histories, with the filter response functions given by the COSMOS survey.

The basic assumption of this SED modeling is that the host galaxies of the
radio sources in our sample are dominated by stellar emission. Thus the MIPS
24 $\mu$m photometric point is excluded because this emission is obviously not
of stellar origin. Furthermore, object 236, a spectroscopically-confirmed QSO
(see Sect. 5.3), is excluded from this analysis. The code also takes into
account extinction (by using the reddening law from Calzetti et al. 2000)
which is applied to the templates. The reddened templates are then shifted in
wavelength searching for the correct redshift. The best fit is obtained
through $\chi^{2}$ minimization. The grid spacings in redshift and in $A_{V}$
are 0.01 and 0.1, respectively. The Hyperz code gives the probability
P($\chi_{rid}^{2}$) associated with $\chi_{rid}^{2}$ (i.e., the reduced
$\chi^{2}$, which is defined as $\chi^{2}_{min}$/$\nu$, where $\nu$ is the
number of degrees of freedom). Because we are essentially interested in the
redshift (the other parameters are affected by the degeneracy), the degrees of
freedom are Nfilters \- 1.

Similarly to Ilbert et al. (2009), we increase the flux errors by 10% when the
probability associated with the $\chi_{rid}^{2}$ is less than 0.01. This
factor does not shift the best-fit photo-z value but broadens the $\chi^{2}$
peak and derived redshift uncertainty. The program provides the redshift of
the object measured at the different confidence intervals defined by the
values of $\chi_{rid}^{2}$. We conservatively choose the solution at a 99%
confidence level. Bolzonella et al. (2000) discussed the effects of degeneracy
in the parameter space defined by star formation history, age, and reddening.
However, they show that this does not strongly affect the value of the
photometric redshift. Hyperzmass, a code which works in a similar way to
Hyperz, returns the mass of the stellar population corresponding to the best
fit and we adopt again the 99%-confidence level to quote the errors on the
mass.

Fig. 15 (left panels) shows the SED fit obtained with Hyperz. The resulting
properties from the modeling are listed in Table 3. Inspection of this figure
indicates that Hyperz does not always return satisfactory results; indeed, in
several cases, it cannot reproduce the bluest part of the spectrum (possibly
due to a young stellar component) or the substantial SED “bump” at long
wavelengths, clearly suggestive of dust emission and often extending even to
the IRAC 3.6 $\mu$m measurements.

###  5.2. ‘2SPD’ fitting technique

In order to improve the quality of the SED fitting we developed a code that
includes 2 Stellar Populations and Dust component (2SPD). More precisely, we
take into account two different stellar populations, typically, one younger
and one older (YSP and OSP, respectively). We model the dust component with a
single (or, in a few cases, two) temperature black-body emission. Thus, at
this stage, we can include the 24 $\mu$m MIPS flux in the fitting process.

The synthetic models used are the same as described in the previous section,
but limiting to those with single stellar population with ages ranging from 1
Myr to 12.5 Gyr. We adopt a dust-screen model for the extinction normalized
with the free parameter $A_{V}$, and the Calzetti et al. (2000) law.

The code searches for the best match between the sum of the different
components (young and old stellar templates and dust emission) and the
photometric points minimizing the appropriate $\chi^{2}$ function. 2SPD
returns the following free parameters: $z$, $A_{V}$, the age of the two
stellar populations, the temperature of the dust component(s), and the various
normalization factors. From these quantities we measure the stellar mass
content of the two stellar populations at 4800 Å rest frame. Similarly to the
case of Hyperz, caution should be exerted before associating these values to
physical quantities because of degeneracy in the parameter space, apart from
the photometric redshifts.

To estimate the errors on the photo-z and mass derivations, we measure the
99%- confident solutions for these quantities. This is computed by varying the
value the parameter of interest until the $\chi^{2}$ value increases by
$\Delta\chi^{2}$ = 6.63, corresponding to a confidence level of 99% for that
parameter.

Fig. 15 (righ panels) shows the outputs of 2SPD code, while Table 4 presents
the resulting parameters of the fit.

The dust emission is usually poorly constrained. In many cases no excess above
the stellar component is required at $\lambda<24\mu$m. In this case dust is
needed only to account for the 24 $\mu$m data-point and often this measurement
is an upper limit. The values reported in Table 4 correspond to the best
fitting model for each galaxy and must be interpreted as approximate rather
than real measurements of the dust emission properties. We will return to the
dust properties in more detail in Sect. 6.4.

###  5.3. COSMOS photometric redshifts

As a comparison with our results we also collect the photometric redshifts
from the COSMOS redshift catalog. Mobasher et al. (2007) provide photometric
redshifts for 860,000 galaxies in the COSMOS field with $i^{+}$ $<$ 25\. Their
technique is based on a $\chi^{2}$ template-fitting procedure applied to the
SEDs derived with up to 16 photometric points from the $u$ to the K band. To
evaluate the reliability of the derived photometric redshifts, they compare
them with spectroscopic redshifts from a sample of 868 galaxies in zCOSMOS
with z $<$ 1.2. The rms scatter between photometric and spectroscopic
redshifts is $\sigma_{(zphot-zspec)/(1+zspec)}$ = 0.031 with a small fraction
of outliers (2.5%).

Ilbert et al. (2009) substantially improve this analysis by including data
from 30 broad, intermediate, and narrow bands covering the wavelength range
from UV to MIR. Redshifts are computed for 607,617 sources (with $i^{+}<$ 26).
The method used by these authors accounts for the contribution of emission
lines to the SEDs. Comparison with 4,148 zCOSMOS spectroscopic redshifts
indicates a dispersion of $\sigma_{(zphot-zspec)/(1+zspec)}$ = 0.007 at
$i^{+}_{AB}<$ 22.5. Nevertheless, the accuracy is strongly degraded at
$i^{+}>$ 25.5. Table 1 shows the photometric redshifts measured by Ilbert et
al. (2009).

Most of the radio galaxies of our sample are included in the COSMOS
Photometric Redshift Catalog Fall 2008, which provides the photometric
redshifts measured by Ilbert et al. (2009). We adopt again a 0$\farcs$3 search
radius and find 29 objects. We report the results in Table 3. One of these
objects (236) turns out to be a spectroscopically-confirmed QSO at z = 2.132
(Prescott et al., 2006). The attempt of Ilbert et al. (2009) to estimate its
photometric redshift failed because the templates used by these authors were
not suitable to fit its AGN-dominated SED. Similarly, Mobasher et al. (2007)
provided for this object a tentative photometric redshift, which is
z=1.23${}^{+0.03}_{-0.17}$, which is clearly inconsistent with the
spectroscopic value.

##  6\. RESUTS

The SED modeling process has been performed for all the objects (except for
the spectroscopically-confirmed QSO, object 236), by using the two template-
fitting techniques (Fig 15). As expected, 2SPD code has turned out to more
reliably model the SEDs of our sources than Hyperz. This is crucial in UV
band. In fact, when significant UV emission appears in the SED, such as, in
objects 31, 258 and 285, Hyperz technique struggles to fit simultaneously such
a UV component and the remaining part of SED, which is well represented by an
OSP. In this case, the best fit result from Hyperz is obtained using an old
composite stellar population. This is due to the fact that the $\chi^{2}$ is
lower if only the the old component of the SED is fitted, rather than the UV
data points which are less well represented. Conversely, as expected, 2SPD
technique is more efficient to treat this case, since it just adds a small
fraction of YSP to a dominant OSP to reproduce the UV emission, by departing
from the concept of single star formation history.

Table 3Hyperz SED fitting ID | redshift | CSP Template | Log M∗
---|---|---|---
| zphot,Hyperz | zphot,2SPD | type | SFH | Age | AV | $\chi_{rid}^{2}$ | Log M⊙
1 | 0.85${}^{+0.05}_{-0.08}$ | 0.88${}^{+0.04}_{-0.05}$ | BC | ssp | 1.434 | 1.00 | 0.73 | 10.98${}^{+0.07}_{-0.04}$
2 | 1.31${}^{+0.14}_{-0.07}$ | 1.33${}^{+0.10}_{-0.09}$ | BC |  $\tau$=0.3 | 0.5088 | 2.9 | 1.43 | 11.00${}^{+0.09}_{-0.12}$
3 | 2.33${}^{+0.18}_{-0.18}$ | 2.20${}^{+0.32}_{-0.44}$ | Ma |  $\tau$=0.3 | 1.0152 | 1.0 | 11.50 | 10.43${}^{+0.07}_{-0.02}$
4 | 1.44${}^{+0.14}_{-0.30}$ | 1.37${}^{+0.10}_{-0.06}$ | BC |  $\tau$=0.1 | 0.3602 | 2.4 | 0.30 | 11.10${}^{+0.16}_{-0.09}$
5 | 1.93${}^{+0.06}_{-0.11}$ | 2.01${}^{+0.22}_{-0.35}$ | Ma | const | 2.0 | 2.3 | 9.57 | 11.68${}^{+0.21}_{-0.15}$
11 | 1.55${}^{+0.28}_{-0.15}$ | 1.57${}^{+0.14}_{-0.09}$ | Ma |  $\tau$=1.0 | 1.0152 | 0.8 | 0.51 | 11.05${}^{+0.11}_{-0.06}$
13 | 1.12${}^{+0.02}_{-0.01}$ | 1.19${}^{+0.08}_{-0.11}$ | BC | ssp | 0.1805 | 2.4 | 12.60 | 10.91${}^{+0.03}_{-0.03}$
16 | 1.04${}^{+0.06}_{-0.15}$ | 0.97${}^{+0.12}_{-0.07}$ | BC | ssp | 0.1805 | 2.2 | 0.78 | 10.58${}^{+0.10}_{-0.05}$
18 | 0.92${}^{+0.02}_{-0.02}$ | 0.92${}^{+0.14}_{-0.11}$ | BC | ssp | 0.1805 | 2.1 | 6.63 | 10.48${}^{+0.03}_{-0.04}$
20 | 0.80${}^{+0.20}_{-0.05}$ | 0.88${}^{+0.02}_{-0.02}$ | BC |  $\tau$=0.3 | 1.7 | 1.4 | 0.76 | 10.85${}^{+0.15}_{-0.08}$
22 | 1.21${}^{+0.09}_{-0.05}$ | 1.30${}^{+0.05}_{-0.04}$ | BC | ssp | 0.5088 | 3.8 | 3.83 | 10.52${}^{+0.87}_{-0.08}$
25 | 1.37${}^{+0.09}_{-0.12}$ | 1.33${}^{+0.11}_{-0.13}$ | BC | ssp | 0.1805 | 2.2 | 2.21 | 11.08${}^{+0.04}_{-0.11}$
26 | 1.09${}^{+0.05}_{-0.10}$ | 1.09${}^{+0.12}_{-0.07}$ | BC |  $\tau$=0.1 | 0.7187 | 1.8 | 1.75 | 11.29${}^{+0.01}_{-0.05}$
28 | 2.61${}^{+0.37}_{-0.22}$ | 2.90${}^{+0.20}_{-0.26}$ | Ma | ssp | 1.0152 | 0.64 | 3.31 | 11.73${}^{+0.03}_{-0.04}$
29 | 1.58${}^{+0.37}_{-0.17}$ | 1.32${}^{+0.23}_{-0.24}$ | Ma |  $\tau$=0.3 | 0.7187 | 1.0 | 0.89 | 10.36${}^{+0.14}_{-0.04}$
30 | 0.99${}^{+0.37}_{-0.11}$ | 1.06${}^{+0.11}_{-0.07}$ | BC |  $\tau$=0.3 | 1.7 | 1.8 | 1.14 | 11.10${}^{+0.14}_{-0.18}$
31 | 0.81${}^{+0.21}_{-0.18}$ | 0.88${}^{+0.03}_{-0.05}$ | BC |  $\tau$=0.3 | 0.7187 | 2.0 | 1.40 | 10.66${}^{+0.09}_{-0.04}$
32 | 3.11${}^{+0.21}_{-0.14}$ | 2.71${}^{+0.38}_{-0.34}$ | BC | ssp | 0.01 | 3.0 | 2.98 | 10.83${}^{+0.03}_{-0.08}$
34 | 1.50${}^{+0.56}_{-0.29}$ | 1.55${}^{+0.41}_{-0.19}$ | Ma |  $\tau$=15 | 2.0 | 2.0 | 0.60 | 10.85${}^{+0.64}_{-0.37}$
36 | 0.91${}^{+0.13}_{-0.08}$ | 1.07${}^{+0.10}_{-0.04}$ | BC |  $\tau$=0.6 | 0.0151 | 4.6 | 1.82 | 10.05${}^{+0.11}_{-0.21}$
37 | 2.04${}^{+0.15}_{-0.24}$ | 1.38${}^{+0.43}_{-0.42}$ | Ma | ssp | 0.0132 | 1.6 | 11.73 | 11.86${}^{+0.13}_{-0.03}$
38 | 1.34${}^{+0.14}_{-0.42}$ | 1.30${}^{+0.17}_{-0.28}$ | Ma | ssp | 0.01995 | 2.8 | 1.23 | 10.82${}^{+0.13}_{-0.17}$
39 | 0.80${}^{+0.23}_{-0.05}$ | 1.10${}^{+0.05}_{-0.05}$ | BC | const | 1.7 | 3.7 | 8.72 | 10.80${}^{+0.15}_{-0.12}$
52 | 0.75${}^{+0.12}_{-0.11}$ | 0.74${}^{+0.18}_{-0.19}$ | BC |  $\tau$=0.3 | 0.5088 | 2.0 | 1.60 | 10.54${}^{+0.19}_{-0.09}$
70 | 2.39${}^{+0.52}_{-0.14}$ | 2.32${}^{+0.53}_{-0.20}$ | Ma |  $\tau$=15.0 | 1.434 | 1.4 | 1.16 | 11.17${}^{+0.19}_{-0.12}$
202 | 0.95${}^{+0.44}_{-0.24}$ | 1.31${}^{+0.09}_{-0.12}$ | BC |  $\tau$=0.3 | 1.0152 | 2.6 | 0.83 | 10.60${}^{+0.28}_{-0.14}$
219 | 1.04${}^{+0.07}_{-0.15}$ | 1.03${}^{+0.02}_{-0.04}$ | BC | ssp | 0.1278 | 2.6 | 1.98 | 10.96${}^{+0.12}_{-0.01}$
224 | 1.07${}^{+0.11}_{-0.14}$ | 1.10${}^{+0.10}_{-0.04}$ | BC |  $\tau$=0.3 | 1.0152 | 2.0 | 0.85 | 10.92${}^{+0.12}_{-0.21}$
226 | 1.98${}^{+0.16}_{-0.27}$ | 2.35${}^{+0.63}_{-0.31}$ | Ma | ssp | 0.01 | 2.0 | 9.57 | 10.08${}^{+0.02}_{-0.05}$
228 | 1.30${}^{+0.09}_{-0.06}$ | 1.31${}^{+0.05}_{-0.07}$ | BC |  $\tau$=0.1 | 0.0263 | 4.2 | 1.25 | 10.35${}^{+0.20}_{-0.40}$
234 | 1.02${}^{+0.08}_{-0.04}$ | 1.10${}^{+0.14}_{-0.08}$ | BC | ssp | 0.1278 | 3.0 | 1.66 | 10.74${}^{+0.08}_{-0.01}$
258 | 0.83${}^{+0.11}_{-0.06}$ | 0.96${}^{+0.19}_{-0.13}$ | BC |  $\tau$=3.0 | 6.5 | 1.0 | 5.27 | 10.90${}^{+0.09}_{-0.07}$
285 | 1.22${}^{+0.13}_{-0.06}$ | 1.10${}^{+0.13}_{-0.08}$ | BC | const | 0.0151 | 2.8 | 4.56 | 10.03${}^{+0.18}_{-0.01}$
Table 42SPD SED fitting ID | redshift | YSP | OSP | log M∗ | Dust | IR excess | UV
---|---|---|---|---|---|---|---
| zphot | Age | AV | $f_{YSP}$ | Log M∗ | Age | AV |  | Tdust | Ldust | LIRexc. | $\alpha_{8-24}$ | LUV
1 | 0.88${}^{+0.04}_{-0.05}$ | 0.02 | 0.80 | 1.5% | 0.49% | 4.0 | 0.15 | 10.08${}^{+0.04}_{-0.04}$ | 140 | 2.4 |  $<$43.31 |  |
2 | 1.33${}^{+0.10}_{-0.09}$ | 0.05 | 0.90 | 22.1% | 1.3% | 2.0 | 0.98 | 11.00${}^{+0.04}_{-0.04}$ | 156 | 7.4 |  $<$43.65 |  | 42.53
3 | 2.20${}^{+0.32}_{-0.44}$ | 0.03 | 0.12 | 6.5% | 0.69% | 1.0 | 0.00 | 10.59${}^{+0.08}_{-0.10}$ | 201-347 | 250-59 | 45.27 | -0.01 | 43.21
4 | 1.37${}^{+0.10}_{-0.06}$ | 0.009 | 1.00 | 6.9% | 0.24% | 3.0 | 0.57 | 11.16${}^{+0.04}_{-0.03}$ | 100 | 5.2 |  $<$43.55 |  | 42.71
5 | 2.01${}^{+0.22}_{-0.35}$ | 0.02 | 1.80 | 53.1% | 5.9% | 3.0 | 0.93 | 11.49${}^{+0.04}_{-0.03}$ | 121 | 144.0 | 44.76 |  $>$0.49 |
11 | 1.57${}^{+0.14}_{-0.09}$ | 0.007 | 1.87 | 17.2% | 0.62% | 3.0 | 0.30 | 10.98${}^{+0.10}_{-0.05}$ | 98 | 15.4 |  $<$43.80 |  |
13 | 1.19${}^{+0.08}_{-0.11}$ | 0.03 | 1.17 | 16.2% | 3.2% | 1.0 | 0.58 | 10.72${}^{+0.04}_{-0.03}$ | 155-406 | 7.0-6.2 | 44.23 | -0.41 | 42.65m
16 | 0.97${}^{+0.12}_{-0.07}$ | 0.006 | 1.16 | 14.8% | 0.23% | 2.0 | 0.52 | 10.74${}^{+0.06}_{-0.06}$ | 158 | 4.8 | 43.60 |  $>$-0.12 | 42.42m
18 | 0.92${}^{+0.14}_{-0.11}$ | 0.004 | 2.55 | 50.3% | 13.5% | 0.2 | 1.35 | 10.02${}^{+0.08}_{-0.08}$ | 116 | 13.9 |  $<$43.96 |  |
20 | 0.88${}^{+0.02}_{-0.02}$ | 0.002 | 1.26 | 3.7% | 0.07% | 3.0 | 0.42 | 11.03${}^{+0.02}_{-0.03}$ | 173 | 1.4 |  $<$43.02 |  | 42.25m
22 | 1.30${}^{+0.05}_{-0.04}$ | 0.04 | 1.18 | 10.1% | 0.35% | 2.0 | 1.50 | 11.16${}^{+0.02}_{-0.03}$ | 150 | 11.8 | 43.87 |  $>$0.66 |
25 | 1.33${}^{+0.11}_{-0.13}$ | 0.002 | 1.17 | 4.1% | 0.28% | 0.4 | 1.05 | 10.75${}^{+0.04}_{-0.05}$ | 138 | 8.8 | 43.75 |  $>$0.26 | 42.87m
26 | 1.09${}^{+0.12}_{-0.07}$ | 0.006 | 1.36 | 7.4% | 0.17% | 1.0 | 0.98 | 11.12${}^{+0.04}_{-0.04}$ | 128 | 3.9 |  $<$43.45 |  | 42.50m
28 | 2.90${}^{+0.20}_{-0.26}$ | 0.001 | 1.97 | 24.7% | 3.9% | 0.9 | 0.00 | 11.38${}^{+0.04}_{-0.04}$ | 146 | 47.4 | 44.28 |  $>$0.22 |
29 | 1.32${}^{+0.23}_{-0.24}$ | 0.004 | 0.38 | 11.9% | 0.18% | 0.4 | 0.94 | 10.03${}^{+0.05}_{-0.05}$ | 83 | 16.4 |  $<$43.68 |  | 42.83
30 | 1.06${}^{+0.11}_{-0.07}$ | 0.001 | 1.77 | 3.0% | 0.13% | 2.0 | 0.73 | 11.03${}^{+0.05}_{-0.05}$ | 130 | 3.5 |  $<$43.47 |  |
31 | 0.88${}^{+0.03}_{-0.05}$ | 0.007 | 0.23 | 7.5% | 0.06% | 2.0 | 0.23 | 10.75${}^{+0.03}_{-0.03}$ | 164 | 2.6 |  $<$43.35 |  | 42.86
32 | 2.71${}^{+0.38}_{-0.34}$ | 0.002 | 0.98 | 10.6% | 0.26% | 2.0 | 0.39 | 10.98${}^{+0.04}_{-0.04}$ | 205-452 | 51-58 | 44.96 | -0.98 | 43.05
34 | 1.55${}^{+0.41}_{-0.19}$ | 0.005 | 0.55 | 11.0% | 0.05% | 3.0 | 0.51 | 10.99${}^{+0.07}_{-0.07}$ | 140 | 19.7 |  $<$44.04 |  | 42.82
36 | 1.07${}^{+0.10}_{-0.04}$ | 0.008 | 2.30 | 42.4% | 2.0% | 3.0 | 1.07 | 10.83${}^{+0.02}_{-0.03}$ | 186 | 4.4 | 43.46 |  $>$0.20 |
37 | 1.38${}^{+0.43}_{-0.42}$ | 0.001 | 0.00 | 1.0% | 0.09% | 0.3 | 0.09 | 10.61${}^{+0.11}_{-0.11}$ | 213-614 | 84-22 | 45.03 | -0.25 | 44.10
38 | 1.30${}^{+0.17}_{-0.28}$ | 0.005 | 0.95 | 27.0% | 0.58% | 0.9 | 0.86 | 10.65${}^{+0.07}_{-0.07}$ | 260 | 6.2 | 43.77 | 0.45 | 43.03
39 | 1.10${}^{+0.05}_{-0.05}$ | 0.007 | 1.86 | 15.0% | 0.83% | 1.0 | 0.97 | 10.88${}^{+0.03}_{-0.03}$ | 80 | 5.3 |  $<$43.20 |  |
52 | 0.74${}^{+0.18}_{-0.19}$ | 0.004 | 0.66 | 13.6% | 0.11% | 2.0 | 0.34 | 10.78${}^{+0.10}_{-0.10}$ | 189 | 2.6 | 43.33 | 0.61 | 42.96
70 | 2.32${}^{+0.53}_{-0.20}$ | 0.009 | 0.48 | 12.8% | 0.87% | 0.4 | 0.74 | 10.65${}^{+0.07}_{-0.03}$ | 180 | 27.5 | 44.07 |  $>$-0.50 | 43.54
202 | 1.31${}^{+0.09}_{-0.12}$ | 0.005 | 0.14 | 1.2% | 0.004% | 3.0 | 0.14 | 10.86${}^{+0.03}_{-0.06}$ | 124 | 7.4 |  $<$43.66 |  | 42.21
219 | 1.03${}^{+0.02}_{-0.04}$ | 0.02 | 0.83 | 6.2% | 0.69% | 0.3 | 1.47 | 10.67${}^{+0.03}_{-0.03}$ | 112 | 2.8 |  $<$43.32 |  | 42.69m
224 | 1.10${}^{+0.10}_{-0.04}$ | 0.004 | 0.63 | 3.8% | 0.04% | 1.0 | 0.76 | 10.71${}^{+0.03}_{-0.04}$ | 155 | 4.5 |  $<$43.53 |  | 42.53m
226 | 2.35${}^{+0.63}_{-0.31}$ | 0.001 | 0.23 | 2.8% | 0.12% | 0.9 | 0.00 | 10.59${}^{+0.07}_{-0.09}$ | 212-467 | 88-36 | 45.01 | -0.46 | 43.34
228 | 1.31${}^{+0.05}_{-0.07}$ | 0.002 | 2.29 | 37.1% | 0.72% | 2.0 | 2.21 | 11.06${}^{+0.03}_{-0.03}$ | 133 | 6.9 |  $<$43.64 |  |
234 | 1.10${}^{+0.14}_{-0.08}$ | 0.006 | 2.30 | 46.3% | 2.2% | 2.0 | 1.05 | 10.83${}^{+0.04}_{-0.04}$ | 104 | 5.8 |  $<$43.52 |  |
258 | 0.96${}^{+0.19}_{-0.13}$ | 0.001 | 0.08 | 1.6% | 0.05% | 0.08 | 1.72 | 10.64${}^{+0.08}_{-0.09}$ | 134 | 3.0 |  $<$43.41 |  | 43.27
285 | 1.10${}^{+0.13}_{-0.08}$ | 0.001 | 0.13 | 2.1% | 0.05% | 0.04 | 1.56 | 10.43${}^{+0.04}_{-0.04}$ | 135-448 | 1.9-1.8 | 43.73 | -0.22 | 43.08
Table 5’Final’ redshifts and radio properties of the sample ID | z | LNVSS | LFIRST | class
---|---|---|---|---
1 | 0.88s |  $<$31.78 | 31.93 | LP
2 | 1.33 | 32.02 | 32.4 | LP
3 | 2.20 | 33.1 | 33.19 | HP
4 | 1.37 | 32.77 | 32.86 | HP
5 | 2.01 | 32.47 | 32.89 | HP
11 | 1.57 |  $<$32.18 | 32.52 | LP
13 | 1.19 | 32.02 | 32.22 | LP
16 | 0.97s | 32.38 | 32.27 | LP
18 | 0.92 | 32.22 | 32.28 | LP
20 | 0.88 |  $<$31.66 | 31.93 | LP
22 | 1.30 |  $<$32.38 | 32.34 | LP
25 | 1.33 | 32.29 | 32.39 | LP
26 | 1.09 | 32.02 | 32.25 | LP
28 | 2.90 | 32.99 | 33.13 | HP
29 | 1.32 | 32.28 | 32.31 | LP
30 | 1.06 | 31.83 | 32.11 | LP
31 | 0.91s | 32.14 | 32.18 | LP
32 | 2.71 | 32.8 | 33.17 | HP
34 | 1.55 | 32.84 | 32.77 | HP
36 | 1.07 | 32.23 | 32.25 | LP
37 | 1.38 | 32.26 | 32.33 | LP
38 | 1.30 | 32.94 | 33.00 | HP
39 | 1.10 |  $<$31.90 | 32.16 | LP
52 | 0.74s |  $<$31.54 | 31.73 | LP
70 | 2.32 | 33.12 | 33.18 | HP
202 | 1.31 | 31.98 | 32.46 | LP
219 | 1.03 |  $<$31.96 | 32.09 | LP
224 | 1.10 | 32.28 | 32.27 | LP
226 | 2.35 |  $<$32.61 | 32.94 | HP
228 | 1.31 | 32.25 | 32.51 | LP
234 | 1.10 | 32.41 | 32.48 | LP
236 | 2.13 | 33.29 | 33.29 | HP
258 | 0.90s | 31.90 | 32.12 | LP
285 | 1.10 | 32.23 | 32.31 | LP

###  6.1. Photometric redshifts

The first test on the accuracy of our photo-z derivation is a comparison with
the spectroscopic redshifts, available for 6 objects (namely 1, 16, 25, 31,
52, and 258). They are all compatible with each other within the errors with
the only exception of object 25. For this object the photometric redshifts
measured with Hyperz and 2SPD (z = 1.37${}^{+0.09}_{-0.12}$ and z =
1.33${}^{+0.11}_{-0.13}$, respectively) are significantly different from the
redshift inferred from its Magellan spectrum (z = 0.7917). However, the
Magellan spectrum does not match with the COSMOS measurements, showing a large
offset, of $\sim$0.8 dex. Apparently, the object observed by Magellan is not
the radio galaxy 25 and we do not consider its spectroscopic-z as reliable.
Reassuringly, we checked that the spectra and photometric data-points agree
for the other 5 objects with available spectra.

Figure 6.— Spectrum of object 25 (red solid line) from the Magellan survey
(Trump et al., 2007) which provides a spectroscopic redshift of 0.7917. We
overlap the synthetic SED (green dashed line) derived from Hyperz at redshift
z = 1.37, fitting the photometric points. The wavelengths on the top
corresponds to the observed wavelengths, and at the bottom to rest frame.

We then compare the photometric redshifts derived with the two SED fitting
techniques. The overall range of $z$ is in both cases from $\sim$0.7 to $\sim$
3\. The median photo-z are 1.21 and 1.30 for Hyperz and 2SPD respectively.
Generally, the photometric redshifts derived with the two methods are
consistent with each other within the errors (Fig. 7, left panel). The
normalized redshift differences ($\Delta z/(1+z)$) are smaller than 0.08 for
all but 5 objects that reach $\Delta z/z\sim 0.12-0.15.$, and a single strong
outlier (object 37, $\Delta z/z\sim 0.28$). For this galaxy, the Hyperz fit to
the SED is particularly weak as it is not able to fit its photometric points
at both FUV and MIR wavelengths.

Figure 7.— Comparison of the photometric-z measured with 2SPD with those
obtained with Hyperz (left) and by Ilbert et al. (2009) (right).

Furthermore, we compare the photometric redshifts measured with our method and
with the template-fitting technique of Ilbert et al. (2009). Objects 2, 22,
28, 32, and 36 are not involved in this comparison since they are not included
in the Ilbert et al. (2009) sample. Fig. 7 (right panel) shows the comparison
between the photo-z resulting from 2SPD and from Ilbert et al. (2009). Again,
the two methods yield similar photometric redshifts with $\Delta
z/(1+z)\lesssim 0.11$ for all but 2 objects, namely 226 and 228 for which we
obtain $\Delta z/(1+z)=0.18$ and 0.25 respectively. In the case of object 226,
similarly to the object 37 discussed above, the SED obtained with 2SPD reveals
a strong excess at both UV and MIR wavelengths. Object 228 is instead a case
of misidentification in the COSMOS multiband catalog. This result stresses the
importance of our detailed work of identification to infer the genuine
properties of the galaxies.

Therefore, unless a spectroscopic redshift is available, we generally prefer
to use the photometric redshift obtained with 2SPD because our code also
includes the dust component which turns out to be important in the SED
modeling. Table 5 shows the redshifts used throughout all the paper.

###  6.2. Radio power distribution

Chiaberge et al. (2009) selected the sample in the COSMOS field looking for
radio sources with radio luminosities below the FR I/FR II break. Since we now
have improved estimates of the photometric redshifts of the objects, we can
derive a more accurate value for their radio power. We show in Fig. 8 the
histograms of K-corrected (using a radio spectral index $\alpha$ = 0.8) radio
powers at 1.4 GHz measured with NVSS and FIRST. The resulting 1.4 GHz
luminosities are in the range 10${}^{31.5}-10^{33.3}$ erg s-1 Hz-1. The local
FR I/FR II break luminosity (Fanaroff & Riley, 1974), converted from 178 MHz
to 1.4 GHz adopting $\alpha$ = 0.8, is L${}_{1.4\,\rm GHz}\sim 10^{32.6}$ erg
s-1 Hz-1. Therefore, the radio distribution of our sample actually straddles
the FR I/FR II break. Nonetheless we remind the reader that radio sources with
a clear FR II morphologies were previously excluded by Chiaberge et al.
(2009).

Figure 8.— Distribution of the K-corrected total radio luminosity (in erg s-1
Hz-1) at 1.4 GHz from the FIRST data (upper panel) and NVSS data (lower
panel). The dashed lines corresponds to the local FR I/FR II break used to
separate the sample in high power (HP, shaded histograms) and low power (LP)
sources.

As already discussed by Chiaberge et al. (2009), the high-frequency radio
observations available for this sample (their rest frame frequency are in the
$\sim$ 3 - 5 GHz range) are not ideal to observe and detect their total radio
structures because they miss diffuse/steep-spectrum radio emission. Lower
frequency radio observations are needed for a more accurate measurement of
their intrinsic total radio power distribution.

Since the radio power distribution is rather broad (covering two orders of
magnitudes), we separate the sample into two groups, including sources of high
and low power (HP and LP, respectively) using the local FR I/FR II break as
divide (see Table 5). This operative definition enables us to explore the role
of radio power in defining the overall properties of these galaxies. Two third
of the sample shows radio luminosities below the break. The LP and HP classes
are, not surprisingly, roughly separated also in redshifts (Fig. 9), since our
sample is limited in radio flux, and indeed the HP sources generally lie at z
$\gtrsim$ 1.5.

Figure 9.— Redshifts versus the K-corrected FIRST radio luminosity at 1.4 GHz
(in erg s-1 Hz-1). The empty points are the LPs and the full points are the
HPs.

###  6.3. Host galaxy properties inferred from SED modeling

We now focus on the properties of the host galaxies inferred from SED modeling
and in particular on their stellar populations.

Figure 10.— Left panel: distribution of stellar masses (in M⊙) of our sample
obtained with 2SPD (upper panel) and with Hyperz (lower panel). Right panel:
comparison between the stellar masses (in M⊙) inferred from SED modeling using
Hyperz and 2SPD techniques. The black histogram (points) represents the HPs,
while the empty one the LPs.

One of the key and more robust result is the measurement of the stellar
content, M∗. The inferred mass range is $\sim$10${}^{10}-10^{12}$ M⊙. The
masses measured with Hyperzmass and with 2SPD are compared in Fig. 10. The two
techniques return generally consistent values of the stellar masses for most
sources (the median values are in both cases $\sim$ 7 $\times$ 1010 M⊙), but
with a few evident outliers. The stellar masses derived from single- and two-
stellar components are know to differ significantly (e.g. Papovich et al.
2006). Furthermore, the inclusion of dust components by 2SPD has two effects
by changing i) the stellar mass since thermal emission largely contributes to
observed infrared emission, and ii) the photometric redshift. The largest
outlier is object 37 whose SED is poorly fitted by Hyperz with a very high
mass 7.0 $\times$ 1011 M⊙, while a strong dust component and a mass 10 times
lower are required by 2SPD. At the opposite end of the mass distribution, we
find two objects (36 and 226) in which again 2SPD finds a strong dust
component, but where the Hyperz is a factor 3 to 7 lower than that obtained
with 2SPD. Apparently, the presence of dust components has a stronger impact
on the stellar mass estimate than the inclusion of a young stellar population.

Considering the radio power of the sources (see Fig. 11, left panel), HPs show
slightly larger stellar mass content than LPs. The median values (obtained
from 2SPD code) for LPs is 5.8 $\times$ 1010 M⊙, while for HPs is 8.7 $\times$
1010 M⊙. However, the two distributions are very broad and show an almost
complete overlap. A very similar result is found looking for differences in
mass between objects at different redshifts (see Fig. 11, right panel).

Figure 11.— Stellar masses (in M⊙) measured with 2SPD in relation with the
K-corrected FIRST radio powers (in erg s-1 Hz-1) (left panel) and redshifts
(right panel) of the sample. Empty points are LPs, while filled points are
HPs.

Due to the degeneration inherent to the modeling of SEDs the remaining results
of the stellar populations are relatively unconstrained. Overall, if we
concentrate in the whole sample, the SED fitting shows that the hosts are
mainly dominated by a massive OSP with an age of $\sim$ 1$-$3 $\times$ 109
years. Although less reliable, the YSPs required to fit the UV excesses show
ages of $1-30$ Myr, with a contribution to the total mass of the galaxy of
$\lesssim 1\%$ for most sources. The flux contribution at 4800 Å (rest frame)
of the YSP is $<$30% for most of the objects.

###  6.4. Dust emission

As discussed in Section 5.2, the parameters related to the dust emission must
be taken as approximate rather than a real measurements of the dust component.
Nonetheless, it is still significant that dust emission is required to
adequately model the SEDs of 15 objects due to the detection of emission at 24
$\mu$m. In addition significant excesses above the stellar emission are
observed also at shorter infrared wavelengths in 8 of these galaxies. In order
to explore the dust properties we estimated the residuals between the best
fitting model including only the stellar component and the data-points,
looking for an excess at long wavelengths due to dust. We then integrated the
residuals to obtain the total dust luminosity in the range covered by the
Spitzer data, i.e. $\sim$ 3 - 26 $\mu$m. The integration is performed by
assuming that the spectrum is represented by a multiple step function (see
Fig. 12).

Figure 12.— Residuals obtained by subtracting the stellar component from the
SED in the 5-bands Spitzer data which we ascribe to the dust emission. We show
the results for object 3 (black) and 5 (red) as an example. The step function
connecting the data-points is used to obtain the IR excess flux, i.e., the
total dust flux.

The resulting dust luminosities (see Table 4), estimated as IR excess
luminosities, are in the range $L_{\rm dust}\sim 10^{43.5}-10^{45.5}$ erg s-1.
A trend links radio and dust emission (see Fig. 13, left panel) with most
(7/9) HP radio-sources showing a significant dust emission with luminosities
larger than $L_{\rm dust}\sim 10^{44}$ erg s-1. The LPs are instead (with only
2 exceptions) of lower Ldust and in many cases (16 galaxies) only upper limits
can be derived.

We check the statistical significance of such a trend, by using a censored
statistical analysis (ASURV, Lavalley et al. 1992) which takes into account of
the presence of upper limits. Using the generalized Kendall’s $\tau$ test
(Kendall, 1983), the probability that a fortuitous correlation appears is
0.0003. However, the common dependence of the two luminosities on redshift
might play an important role. Therefore, we perform a partial correlation
analysis (Akritas & Siebert, 1996) to examine the linear relation between the
luminosities excluding the dependence on redshift. Operatively, we use the
partial Kendall’s $\tau$ test, whose null hypothesis is the absence of the
correlation excluding the redshift variable. The partial Kendall’s coefficient
$\tau$ is 0.20 and the standard deviation is 0.081. The null hypothesis is
then rejected at the level of 0.05.

We also measured the spectral index of the residuals between 8 $\mu$m and 24
$\mu$m, $\alpha_{(8-24)}$. Taking into account only significant excesses at 8
$\mu$m (residuals larger than 3 $\sigma$), this value can be estimated in 8
cases, with values spanning between $\alpha_{(8-24)}\sim$ 1 and -1. For the
remaining objects with only a 24 $\mu$m detection, the limit to the 8 $\mu$m
flux translates into a lower limit of $\alpha_{(8-24)}\gtrsim 1$. This
parameter can be crudely related to the overall dust temperature. By assuming
a single black-body dust component, the values of $\alpha_{(8-24)}$ translate
into a temperature range of 500-850 K and 300-550 K for $\alpha_{(8-24)}=-1$
and $1$, respectively. The derived temperature depends on redshift, with the
lower (upper) values of T being derived for $z=0.75$ ($z=2$).

Fig. 13 (middle panel) shows a broad relation between the IR spectral index
and the IR excess luminosities. This would mean that the decrease of
$\alpha_{(8-24)}$ is ascribable to the increase of the high-temperature dust
component. Since a large radio power also implies a large dust luminosity,
this may suggest a possible AGN nature of the dust heating source for the
objects with largest IR excess.

Figure 13.— Infrared excess luminosity ($\>{\rm erg}\,{\rm s}^{-1}$) versus:
(left panel) K-corrected FIRST radio luminosity ($\>{\rm erg}\,{\rm s}^{-1}$),
and (right panel) spectral index from 8 to 24 $\mu$m estimated from the IR
excess in the SED above the stellar emission. Empty points are the LPs and
full points are the HPs.

###  6.5. UV excess

As noted in Sect. 5.1, the Hyperz code often does not reproduce satisfactorily
the bluest part of the SED and we chose to model such part of spectrum with a
young stellar population, included by our 2SPD. The inclusion of this
component does not alter significantly the photometric redshift or the
galaxies mass, but it is clearly of great importance to understand the nature
of these radio galaxies. Nonetheless, inspection of the SEDs obtained with
2SPD indicates that the UV excesses (above the contribution of the old stellar
component) are usually very poorly constrained. Furthermore, the very stellar
origin is not granted and the UV excess might hide an AGN contribution.

It is necessary to introduce a model-independent criterion to assess which
sources actually show an UV excess and to estimate its luminosity. We visually
inspected all SEDs, searching for sources with a substantial flattening in the
SED at short wavelengths or with a change of the slope between the OSP and the
emission in the UV band. Fifteen sources show a clear UV excess (namely object
2, 3, 4, 29, 31, 32, 34, 37, 38, 52, 70, 202, 226, 258, and 285) and
additional 7 objects show a marginal UV excess (namely object 13, 16, 20, 25,
26, 219, and 224). In the remaining galaxies the SEDs in UV bands drop sharply
and are well reproduced by the emission from OSPs. In order to quantify the UV
contribution, we measure the flux at 2000 Å in the rest frame, LUV, from the
best fitting model, for those objects showing the UV excess. The UV
luminosities range is $10^{42}-10^{44}$ erg s-1. HPs show larger UV
luminosities than LPs by a factor 2.5, even though the strongest UV excess is
seen in the LP radio galaxy 37. If we concentrate the sources which do not
show an UV excess, the 2000 Å luminosity is $\lesssim 10^{43}$ erg s-1
(mostly, $\sim 10^{42}$ erg s-1).

In order to explore the nature of the UV emission, it is useful to compare the
UV luminosity with the multiband properties obtained in the previous sections.
No relations appear to link the UV excess to the redshift and the mass of the
galaxy. Conversely, the UV luminosity appears to be linked to the radio power
and IR excess luminosity (Fig. 14).

For the radio and UV luminosities the linear correlation coefficient is
$r=0.404$ and the probability of the presence of a fortuitous relation is P=
0.062. The partial Kendall’s coefficient $\tau$ is 0.17 and the standard
deviation is 0.15: the null hypothesis, i.e. the absence of the correlation,
can not be rejected.

For the IR-UV relation, the generalized Kendall’s $\tau$ test returns a
probability of no correlation of P= 0.0072. We obtain a partial Kendall’s
coefficient $\tau$ = 0.41 and a standard deviation of 0.069: the null
hypothesis, i.e. the absence of the correlation, is rejected at the level of
0.05.

Figure 14.— UV luminosity ($\>{\rm erg}\,{\rm s}^{-1}$) measured at 2000 Å
versus: (left panel) K-corrected FIRST radio luminosity ($\>{\rm erg}\,{\rm
s}^{-1}$) , and (right panel) infrared excess luminosity ($\>{\rm erg}\,{\rm
s}^{-1}$) . Empty points are the LPs and full points are the HPs.

##  7\. Discussion and Conclusions

We used the multiwavelength data provided by the COSMOS survey to analyze the
first sizeable sample of 34 low-power radio galaxies at high redshifts (z
$\gtrsim$ 1), selected by Chiaberge et al. (2009). We performed a careful
visual inspection of all their multiwavelength counterparts to identify their
genuine emission and to infer their SEDs. Those would have been compromised in
case of a blind use of the COSMOS catalog. Taking advantage of the FUV-MIR
spectral coverage, we modeled their SEDs using two different template fitting
techniques: i) the Hyperz code that uses a single stellar population and ii)
our own code 2SPD that includes also dust component(s) and a young stellar
population. We analyzed the properties of the SEDs of the sample. We here
summarize the main results and briefly discuss them:

  1. (i)

The photometric redshifts of these radio sources range from $\sim$0.7 to 3\.
The photo-z measured with the two techniques are generally consistent with
each other, with those measured by Ilbert et al. (2009), and with the
available spectroscopic redshifts. In addition, we measured the photometric
redshifts for five objects which were not included in the Ilbert et al. (2009)
sample.

  2. (ii)

The new and accurate measurements of the photometric redshifts enable us to
infer the radio power distribution of these objects. The resulting K-corrected
1.4 GHz luminosities are in the range 10${}^{31.5}-10^{33.3}$ erg s-1 Hz-1,
straddling the FR I/FR II break luminosity (L${}_{1.4\,\rm GHz}\sim 10^{32.6}$
erg s-1 Hz-1) as defined for low redshifts galaxies.

  3. (iii)

The resulting stellar masses are mostly confined in the range
$\sim$1010.5-11.5 M⊙, with a median value of 7 $\times$ 1010 M⊙.

  4. (iv)

The SED of most objects of the sample is consistent with the presence of a
dominant contribution from an old stellar population with an age $\sim
1-3\times$ 109 years. However, significant excesses are often observed at the
shortest and/or longest wavelengths.

  5. (v)

A dust component is needed to account for the 24 $\mu$m emission in 15
galaxies and significant excesses above the stellar emission are observed also
at shorter infrared wavelengths in 8 of these galaxies. Estimates of the dust
luminosity yield values in the range $L_{\rm dust}\sim 10^{43.5}-10^{45.5}$
erg s-1. The overall dust temperature, estimated for the 8 radio-galaxies with
a substantial dust excess at $\lambda\lesssim 8$$\mu$m, is in the range
300-850 K.

  6. (vi)

Inspection of the SED obtained with 2SPD indicates that the UV excesses (above
the contribution of the old stellar component) are often present
(significantly in 15 sources and marginally in 7 sources), but they are
usually weakly constrained. The UV luminosities measured at 2000 Å (rest
frame) is in the range $10^{42}-10^{44}$ erg s-1.

  7. (vii)

Although the censored analysis does not provide significantly high statistical
parameters, we can tentatively confirm the presence of positive links between
the dust emission with both the radio and UV luminosities. For these
relations, the possibility of a common luminosity dependence on distances is
rejected at the level of 0.05.

The selection performed by Chiaberge et al. (2009) with the aim of searching
for FR I candidates at z $\sim 1-2$ turned out to be successful. In fact, our
work confirms i ) their location in the range of redshifts aimed with the
selection process, although extending slightly below 1 and up to z $\sim$ 3
and ii) the low radio luminosity of the sample, generally consistent with
those of the local FR I, although 1/3 of the sample exceeds the local FR I/FR
II luminosity break. The extension of their radio power to larger luminosities
does not necessarily imply a FR II nature for those sources. In fact the
sources with a clear FR II morphology were excluded from the analysis. In
addition, in the local Universe, the radio distribution of FR Is is broad and
overlaps with the FR II distribution (Zirbel & Baum, 1995). Furthermore, this
overlap increases at higher radio frequencies.

Overall, the hosts of these high-z low-luminosity radio-sources are similar to
those of the local FR I which usually live in red massive early-type galaxies
(e.g., Zirbel 1996; Best et al. 2005; Baldi & Capetti 2008; Smolčić 2009;
Baldi & Capetti 2010).

However, the SED modeling reveals that additional components to the old
stellar population have to be included to account for the emission at the SED
extremes, i.e. in either the UV or in the MIR band, in most of the sources of
the sample. This behavior is not seen in the low-z FR I that are generally
faint in UV (both from the point of view of star formation and nuclear
emission, Chiaberge et al. 2002; Baldi & Capetti 2008). Similarly, their MIR
luminosities exceed by a very large factor (between 30 and 3000 for the MIR
detected sources) the typical low-z FR I luminosities ($\sim 10^{42}\>{\rm
erg}\,{\rm s}^{-1}$, Hardcastle et al. 2009). Conversely, the UV and MIR
properties are somewhat similar to those of local FR IIs, which show bluer
color (e.g., Baldi & Capetti 2008; Smolčić 2009) and large dust amount (e.g.,
de Koff et al. 2000; Dicken et al. 2010) than FR Is.

The origin of the MIR and UV emission can not be firmly established based on
the available data. The estimate of the dust temperature (possible for only 8
objects) is in the range expected from dust heating from a quasar-like nucleus
($T\gtrsim 300$ K, e.g. Siebenmorgen et al. 2004; Ogle et al. 2006) and far
larger than the dust associated with star formation (e.g., Hwang et al. 2010;
Sreenilayam & Fich 2011; Boquien et al. 2011; Patel et al. 2011). The high end
of the MIR luminosity reaches values similar to those found in high power
radio-galaxies and QSOs. However, dust emission is seen only in less than half
the sources of the sample.

Similarly, a UV component in excess to the old stellar population is firmly
detected in the same number of sources (but there is not a one-to-one
correspondence between UV and MIR emission). The observed UV luminosities are
much larger that the faint non-thermal UV-nuclei seen in FR I and, instead,
similar to those FR IIs (Chiaberge et al., 1999, 2002), dominated by the
accretion disk emission.

Summing up, we find that the sources of the sample display a wide variety of
properties, despite the relatively narrow range in radio luminosities. The 8
objects with the strongest MIR excess, and with high dust temperature
indicative of a quasar-like nature, also show a significant UV excess. To this
group we must obviously add the object 236, the spectroscopically confirmed
QSO. We must note, however, that the SED of these 8 sources is much redder
than that of object 236 and that their HST images do not provide any evidence
for the presence of bright unresolved nuclei, as described by Chiaberge et al.
(2009) based on visual inspection of the data.

At the opposite end of the radio luminosity distribution, there are 7 galaxies
not showing any UV excess nor any 24 $\mu$m emission (and this number raises
to 11 by including also the objects with only marginal UV excesses).

The remaining galaxies, amounting to about half of the sample, show an excess
only in either the UV or MIR band. For those objects which do not show an UV
excess or are not detected at 24 $\mu$m, the limits to the UV/MIR ratio are
broadly consistent with the values obtained for the remaining galaxies of the
sample. In particular we do not find objects with a high MIR luminosity
without an UV excess that might be expected in the case of an obscured QSO.
Unfortunately, with the available information we cannot distinguish between an
origin related to star formation or to an active nucleus. In addition,
relatively large amounts of dust, suggested by the dust luminosities, indicate
that possible obscuration may prevent us from detecting UV emission.

A further detailed analysis of the nuclear properties is needed to understand
which type of AGN are associated with these high-z radio galaxies. This will
include their X-ray emission (Tundo et al., 2011) and the radio core flux,
available from the COSMOS/VLA data, but which we defer to a future study. This
might also provide new insights on the controversial FR I-QSO association
(e.g., Falcke et al. 1995; Baum et al. 1995; Cao & Rawlings 2004; Blundell &
Rawlings 2001, see also Blundell 2003 for review on this subject).

Summarizing, this work validates the first sizeable sample of low-luminosity
radio galaxies at high redshifts, 0.7 $\lesssim$ z $\lesssim$ 3\. This opens
the possibility to perform a detailed comparison of the host and nuclear
properties of these sources with those of i) the local low-luminosity radio-
galaxies, ii) the powerful radio-galaxies in the same redshift range, and iii)
the population of non active galaxies at $z\sim 1-3$. These issues will be
addressed in a forthcoming paper.

R.D.B. acknowledges the financial support (grant DDRF D0001.82439) from Space
Telescope Science Institute, Baltimore. We are grateful to M. Bolzonella, C.
Maraston, and J. Pforr. which significantly help with the SED modeling. We
also thank the referee and A. Celotti for their contributions to improve the
paper.

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$\begin{array}[]{ccc}\vspace{2em}\includegraphics[scale={0.33},angle={90}]{1cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed1_2sp_dustnew.ps}\\\
\vspace{1em}\includegraphics[scale={0.33},angle={90}]{2cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed2_2sp_dustnew.ps}\\\
\vspace{1em}\includegraphics[scale={0.33},angle={90}]{3cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed3_2sp_2dustnew.ps}\\\
\vspace{1em}\includegraphics[scale={0.33},angle={90}]{4cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed4_2sp_dustnew.ps}\end{array}$

Figure 15.— Left panels: SEDs of the sample (green line) fitting the
photometric points, as result of Hyperz. We plot the photometric point at 24
$\mu$m (Spitzer/MIPS), if available, but is not included in the modeling.
Right panels: SEDs of the objects fitted with 2SPD. The total model is the
green line, the YSP is the blue line, the OSP is the red line, and the dust
component(s) is the light blue line. For both the panels, the wavelengths on
top of the plots correspond to observed wavelengths, while those on bottom are
at rest frame. For object 236 (QSO) we only show the photometric points.

$\begin{array}[]{ccc}\vspace{2em}\includegraphics[scale={0.33},angle={90}]{5cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed5_2sp_dustnew.ps}\\\
\vspace{1em}\includegraphics[scale={0.33},angle={90}]{11cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed11_2sp_dustnew.ps}\\\
\vspace{2em}\includegraphics[scale={0.33},angle={90}]{13cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed13_2sp_2dustnew.ps}\\\
\vspace{1em}\includegraphics[scale={0.33},angle={90}]{16cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed16_2sp_dustnew.ps}\end{array}$

Figure 15.— Continued

$\begin{array}[]{ccc}\vspace{2em}\includegraphics[scale={0.33},angle={90}]{18cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed18_2sp_dustnew.ps}\\\
\vspace{1em}\includegraphics[scale={0.33},angle={90}]{20cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed20_2sp_dustnew.ps}\\\
\vspace{1em}\includegraphics[scale={0.33},angle={90}]{22cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed22_2sp_dustnew.ps}\\\
\vspace{1em}\includegraphics[scale={0.33},angle={90}]{25cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed25_2sp_dustnew.ps}\end{array}$

Figure 15.— Continued

$\begin{array}[]{ccc}\vspace{2em}\includegraphics[scale={0.33},angle={90}]{26cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed26_2sp_dustnew.ps}\\\
\vspace{1em}\includegraphics[scale={0.33},angle={90}]{28cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed28_2sp_dustnew.ps}\\\
\vspace{1em}\includegraphics[scale={0.33},angle={90}]{29cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed29_2sp_dustnew.ps}\\\
\vspace{2em}\includegraphics[scale={0.33},angle={90}]{30cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed30_2sp_dustnew.ps}\end{array}$

Figure 15.— Continued

$\begin{array}[]{ccc}\vspace{2em}\includegraphics[scale={0.33},angle={90}]{31cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed31_2sp_dustnew.ps}\\\
\vspace{1em}\includegraphics[scale={0.33},angle={90}]{32cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed32_2sp_2dustnew.ps}\\\
\vspace{2em}\includegraphics[scale={0.33},angle={90}]{34cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed34_2sp_dustnew.ps}\\\
\vspace{1em}\includegraphics[scale={0.33},angle={90}]{36cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed36_2sp_dustnew.ps}\end{array}$

Figure 15.— Continued

$\begin{array}[]{ccc}\vspace{2em}\includegraphics[scale={0.33},angle={90}]{37cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed37_2sp_2dustnew.ps}\\\
\vspace{1em}\includegraphics[scale={0.33},angle={90}]{38cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed38_2sp_dustnew.ps}\\\
\vspace{1em}\includegraphics[scale={0.33},angle={90}]{39cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed39_2sp_dustnew.ps}\\\
\vspace{1em}\includegraphics[scale={0.33},angle={90}]{52cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed52_2sp_dustnew.ps}\end{array}$

Figure 15.— Continued

$\begin{array}[]{ccc}\vspace{2em}\includegraphics[scale={0.33},angle={90}]{70cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed70_2sp_dustnew.ps}\\\
\vspace{1em}\includegraphics[scale={0.33},angle={90}]{202cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed202_2sp_dustnew.ps}\\\
\vspace{1em}\includegraphics[scale={0.33},angle={90}]{219cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed219_2sp_dustnew.ps}\\\
\vspace{2em}\includegraphics[scale={0.33},angle={90}]{224cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed224_2sp_dustnew.ps}\end{array}$

Figure 15.— Continued

$\begin{array}[]{cccc}\vspace{2em}\includegraphics[scale={0.33},angle={90}]{226cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed226_2sp_2dustnew.ps}\\\
\vspace{1em}\includegraphics[scale={0.33},angle={90}]{228cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed228_2sp_dustnew.ps}\\\
\vspace{2em}\includegraphics[scale={0.33},angle={90}]{234cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed234_2sp_dustnew.ps}\\\
\vspace{1em}\includegraphics[scale={0.33},angle={90}]{236cspnew.ps}&\end{array}$

Figure 15.— Continued

$\begin{array}[]{cc}\vspace{2em}\includegraphics[scale={0.33},angle={90}]{258cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed258_2sp_dustnew.ps}\\\
\vspace{2em}\includegraphics[scale={0.33},angle={90}]{285cspnew.ps}&\includegraphics[scale={0.33},angle={90}]{sed285_2sp_2dustnew.ps}\end{array}$

Figure 15.— Left panels: SEDs of the sample (green line) fitting the
photometric points, as result of Hyperz. We plot the photometric point at 24
$\mu$m (Spitzer/MIPS), if available, but is not included in the modeling.
Right panels: SEDs of the objects fitted with 2SPD. The total model is the
green line, the YSP is the blue line, the OSP is the red line, and the dust
component(s) is the light blue line. For both the panels, the wavelengths on
top of the plots correspond to observed wavelengths, while those on bottom are
at rest frame. For object 236 (QSO) we only show the photometric points

Table 6COSMOS multiwavelength counterparts of the sample

ID | $u^{*}$ | $B_{J}$ | $g^{+}$ | $V_{J}$ | $r^{+}$ | $i^{*}$ | $i^{+}$ | $F814W$ | $z^{+}$ | $J$ | $K_{S}$ | $K$
---|---|---|---|---|---|---|---|---|---|---|---|---
(1) | (2) | (3) | (4) | (5) | (6) | (7) | (8) | (9) | (10) | (11) | (12) | (13)
1 | 26.56$\pm$0.15 | 25.89$\pm$0.13 | 25.61$\pm$0.10 | 24.42$\pm$0.05 | 23.55$\pm$0.03 | 22.24$\pm$0.05 | 22.29$\pm$0.11∗ | 21.87$\pm$0.06 | 21.32$\pm$0.01 | 20.46$\pm$0.10∗ | 19.59$\pm$0.09∗ | 19.57$\pm$0.03
2 | 26.79$\pm$0.20∗ | 26.43$\pm$0.11∗ | 26.29$\pm$0.11∗ | 25.93$\pm$0.10∗ | 25.26$\pm$0.08∗ | 24.35$\pm$0.45∗ | 24.43$\pm$0.10∗ | 24.49$\pm$0.34∗ | 23.44$\pm$0.07∗ | 22.18$\pm$0.15∗ | 20.85$\pm$0.02 | 20.90$\pm$0.15
3 | 26.69$\pm$0.17 | 25.58$\pm$0.10 | 25.81$\pm$0.12 | 25.60$\pm$0.11 | 25.53$\pm$0.10 | 25.31$\pm$0.78 | 25.10$\pm$0.08 | 24.96$\pm$0.79 | 24.48$\pm$0.12 | 23.39$\pm$0.25∗ | 21.57$\pm$0.03 | 22.84$\pm$0.96
4 | 26.64$\pm$0.17 | 25.93$\pm$0.13 | 25.88$\pm$0.13 | 25.46$\pm$0.10 | 25.12$\pm$0.08 | 24.27$\pm$0.36 | 24.20$\pm$0.05 | 23.82$\pm$0.29 | 23.27$\pm$0.04 | 21.93$\pm$0.10∗ | 20.53$\pm$0.01 | 20.74$\pm$0.16
5 |  $<$25.77∗ |  $<$25.68∗ |  $<$25.53∗ |  $<$25.17∗ |  $<$25.23∗ | 24.12$\pm$0.27 | 24.59$\pm$0.41∗ | 23.98$\pm$0.33 | 23.41$\pm$0.25 | 22.24$\pm$0.13∗ | 20.65$\pm$0.12 | 20.82$\pm$0.15
11 |  $<$29.40∗ | 28.75$\pm$1.14 | 28.00$\pm$0.76 | 27.11$\pm$0.35 | 26.27$\pm$0.17 | 25.16$\pm$0.84 | 24.96$\pm$0.07 | 24.64$\pm$0.60 | 24.09$\pm$0.08 | 22.28$\pm$0.07∗ | 21.18$\pm$0.16 | 21.75$\pm$0.49
13 | 26.14$\pm$0.15∗ | 25.75$\pm$0.13∗ | 25.74$\pm$0.13 | 25.02$\pm$0.09∗ | 24.38$\pm$0.06∗ | 23.20$\pm$0.17 | 23.29$\pm$0.04∗ | 22.90$\pm$0.13 | 22.21$\pm$0.02 | 21.44$\pm$0.09∗ | 20.46$\pm$0.10 | 20.54$\pm$0.16
16 | 26.07$\pm$0.12 | 25.71$\pm$0.12 | 25.32$\pm$0.11 | 24.84$\pm$0.08 | 24.23$\pm$0.05 | 23.10$\pm$0.14 | 23.15$\pm$0.03 | 22.85$\pm$0.10 | 22.16$\pm$0.02 | 21.47$\pm$0.08∗ | 20.40$\pm$0.08 | 20.50$\pm$0.11
18 | 26.16$\pm$0.13∗ | 25.76$\pm$0.10∗ | 25.51$\pm$0.10 | 24.61$\pm$0.08∗ | 23.99$\pm$0.06∗ | 22.91$\pm$0.11∗ | 22.80$\pm$0.04∗ | 22.51$\pm$0.10∗ | 22.04$\pm$0.03∗ | 21.54$\pm$0.07∗ | 20.71$\pm$0.04 | 20.92$\pm$0.18
20 | 26.11$\pm$0.15 | 25.74$\pm$0.15 | 25.51$\pm$0.10 | 24.57$\pm$0.06 | 23.67$\pm$0.04 | 22.35$\pm$0.09 | 22.40$\pm$0.02 | 22.19$\pm$0.08 | 21.53$\pm$0.02 | 20.79$\pm$0.06∗ | 19.91$\pm$0.07 | 20.10$\pm$0.07
22 |  $<$27.38∗ |  $<$27.44∗ |  $<$26.97∗ |  $<$26.85∗ |  $<$26.00∗ |  |  $<$25.13∗ |  | 23.93$\pm$0.06 | 22.14$\pm$0.17∗ | 20.79$\pm$0.10 | 21.24$\pm$0.15
25 | 26.66$\pm$0.22 | 25.36$\pm$0.09 | 25.44$\pm$0.10 | 24.97$\pm$0.07 | 24.55$\pm$0.05 | 23.64$\pm$0.32 | 23.52$\pm$0.03 | 23.32$\pm$0.19 | 22.57$\pm$0.03 | 21.43$\pm$0.09∗ | 20.57$\pm$0.05 | 20.59$\pm$0.17
26 | 26.58$\pm$0.20 | 25.44$\pm$0.10 | 25.56$\pm$0.10 | 24.92$\pm$0.07 | 23.90$\pm$0.04 | 22.68$\pm$0.10 | 22.60$\pm$0.02 | 22.32$\pm$0.09 | 21.57$\pm$0.01 | 20.42$\pm$0.04∗ | 19.42$\pm$0.03 | 19.48$\pm$0.04
28 |  $<$27.39∗ |  $<$27.50∗ |  $<$27.20∗ |  $<$26.90∗ |  $<$26.80∗ |  | 25.55$\pm$0.30∗ |  | 24.27$\pm$0.20∗ | 22.48$\pm$0.20∗ | 20.99$\pm$0.08∗ |
29 | 25.81$\pm$0.12 | 25.51$\pm$0.10 | 25.49$\pm$0.09 | 25.31$\pm$0.09 | 25.18$\pm$0.08 | 25.00$\pm$0.87 | 24.78$\pm$0.07 | 24.92$\pm$0.74 | 24.30$\pm$0.11 | 22.86$\pm$0.27∗ | 22.19$\pm$0.10 | 22.89$\pm$0.80
30 | 27.85$\pm$0.63 | 26.77$\pm$0.22 | 26.82$\pm$0.23 | 25.90$\pm$0.13 | 24.69$\pm$0.06 | 23.41$\pm$0.19 | 23.41$\pm$0.03 | 23.01$\pm$0.15 | 22.32$\pm$0.02 | 21.25$\pm$0.08∗ | 20.02$\pm$0.04 | 20.19$\pm$0.09
31 | 24.75$\pm$0.08 | 24.50$\pm$0.06 | 24.51$\pm$0.05 | 23.94$\pm$0.04 | 23.42$\pm$0.03 | 22.33$\pm$0.08 | 22.33$\pm$0.02 | 22.07$\pm$0.07 | 21.61$\pm$0.02 | 20.97$\pm$0.08∗ | 20.12$\pm$0.05 | 20.32$\pm$0.09
32 |  $<$27.54∗ | 27.85$\pm$0.48∗ |  $<$27.19∗ |  $<$27.00∗ | 26.44$\pm$0.30∗ |  | 26.07$\pm$0.25∗ |  | 25.15$\pm$0.28∗ | 24.13$\pm$0.50∗ | 22.43$\pm$0.17∗ | 23.24$\pm$1.44
34 | 26.09$\pm$0.40∗ | 26.19$\pm$0.20∗ | 26.00$\pm$0.19∗ | 25.58$\pm$0.12∗ | 25.44$\pm$0.10∗ | 25.09$\pm$0.66∗ | 24.88$\pm$0.07∗ | 24.95$\pm$0.25∗ | 24.12$\pm$0.11∗ | 23.10$\pm$0.51∗ | 21.10$\pm$0.29∗ | 20.96$\pm$0.14
36 |  $<$27.8∗ |  $<$26.80∗ |  $<$26.40∗ | 26.18$\pm$0.45∗ | 25.50$\pm$0.31∗ | 24.31$\pm$0.33∗ | 24.15$\pm$0.15∗ | 24.00$\pm$0.25∗ | 23.20$\pm$0.24∗ | 22.04$\pm$0.15∗ | 20.62$\pm$0.07 | 20.44$\pm$0.10
37 | 22.70$\pm$0.02 | 22.35$\pm$0.02 | 22.63$\pm$0.02 | 22.38$\pm$0.02 | 22.18$\pm$0.02 | 21.93$\pm$0.05 | 21.93$\pm$0.01 |  | 21.17$\pm$0.01 | 20.31$\pm$0.05∗ | 20.14$\pm$0.11∗ | 20.06$\pm$0.08
38 | 25.50$\pm$0.10 | 25.00$\pm$0.08 | 24.96$\pm$0.07 | 24.66$\pm$0.06 | 24.20$\pm$0.04 | 23.58$\pm$0.32 | 23.49$\pm$0.03 | 23.47$\pm$0.22 | 22.69$\pm$0.03 | 22.08$\pm$0.16∗ | 20.89$\pm$0.05 | 21.34$\pm$0.22
39 |  $<$27.07∗ | 26.05$\pm$0.15∗ |  $<$26.00∗ | 25.28$\pm$0.09 | 24.42$\pm$0.05 | 23.13$\pm$0.18 | 23.13$\pm$0.03 | 22.86$\pm$0.15 | 22.18$\pm$0.02 | 21.06$\pm$0.08∗ | 19.99$\pm$0.04 | 20.27$\pm$0.09
52 | 24.24$\pm$0.04 | 23.64$\pm$0.04 | 23.76$\pm$0.04 | 23.13$\pm$0.03 | 22.62$\pm$0.02 | 21.62$\pm$0.05 | 21.67$\pm$0.01 | 21.52$\pm$0.05 | 21.18$\pm$0.01 | 20.58$\pm$0.05∗ | 19.79$\pm$0.04 | 19.78$\pm$0.06
70 | 26.16$\pm$0.14∗ | 25.14$\pm$0.09∗ | 25.27$\pm$0.14∗ | 24.79$\pm$0.08∗ | 24.51$\pm$0.06 | 24.34$\pm$0.58 | 24.35$\pm$0.06 | 24.13$\pm$0.43 | 24.01$\pm$0.09 | 23.42$\pm$0.50∗ | 21.76$\pm$0.09 | 21.40$\pm$0.27
202 | 27.03$\pm$0.28 | 27.11$\pm$0.29 | 27.40$\pm$0.42 | 26.92$\pm$0.29 | 25.86$\pm$0.12 | 24.44$\pm$0.53 | 24.41$\pm$0.06 | 24.29$\pm$0.45 | 23.48$\pm$0.05 | 22.19$\pm$0.17∗ | 21.07$\pm$0.05 | 21.62$\pm$0.30
219 | 25.68$\pm$0.10 | 25.09$\pm$0.08 | 25.23$\pm$0.08 | 24.49$\pm$0.05 | 23.81$\pm$0.04 | 22.74$\pm$0.09 | 22.71$\pm$0.02 | 22.37$\pm$0.09 | 21.70$\pm$0.02 | 21.02$\pm$0.06∗ | 19.93$\pm$0.04 | 20.15$\pm$0.37
224 | 26.40$\pm$0.18 | 25.75$\pm$0.13 | 25.52$\pm$0.10 | 25.36$\pm$0.09 | 24.62$\pm$0.06 | 23.52$\pm$0.27 | 23.49$\pm$0.03 |  | 22.50$\pm$0.03 | 21.31$\pm$0.08∗ | 20.39$\pm$0.04 | 20.49$\pm$0.18
226 | 25.40$\pm$0.11 | 25.30$\pm$0.09 | 25.48$\pm$0.11 | 25.23$\pm$0.09 | 25.17$\pm$0.08 | 24.59$\pm$1.00 | 24.75$\pm$0.08 | 24.78$\pm$0.68 | 24.37$\pm$0.11 | 23.41$\pm$0.46∗ | 21.90$\pm$0.10 | 21.75$\pm$0.36
228 |  $<$28.50∗ |  $<$27.70∗ |  $<$27.70∗ |  $<$26.80∗ | 26.17$\pm$0.16 |  | 25.17$\pm$0.10 | 25.17$\pm$0.67 | 24.38$\pm$0.11 | 23.10$\pm$0.40∗ | 21.14$\pm$0.06 | 21.23$\pm$0.22
234 | 27.40$\pm$0.51∗ | 26.43$\pm$0.22∗ | 26.65$\pm$0.28∗ | 25.50$\pm$0.12∗ | 24.88$\pm$0.08∗ | 23.76$\pm$0.30 | 23.52$\pm$0.08∗ |  | 22.76$\pm$0.03 | 21.38$\pm$0.20∗ | 20.31$\pm$0.11∗ | 20.57$\pm$0.39
236 | 20.92$\pm$0.01 | 20.73$\pm$0.01 | 21.02$\pm$0.01 | 20.59$\pm$0.01 | 20.65$\pm$0.01 | 20.32$\pm$0.02 |  | 20.26$\pm$0.02 | 19.94$\pm$0.01 | 20.05$\pm$0.04∗ | 19.27$\pm$0.02 | 19.32$\pm$0.04
258 | 24.01$\pm$0.04 | 23.54$\pm$0.04 | 23.70$\pm$0.03 | 23.19$\pm$0.03 | 22.68$\pm$0.02 | 21.87$\pm$0.06 | 21.85$\pm$0.01 | 21.56$\pm$0.04 | 21.11$\pm$0.01 | 20.61$\pm$0.05∗ | 19.62$\pm$0.04 | 19.71$\pm$0.04
285 | 24.60$\pm$0.18∗ | 24.65$\pm$0.15∗ | 24.63$\pm$0.24∗ | 23.88$\pm$0.14∗ | 23.70$\pm$0.12∗ | 23.32$\pm$0.28 | 23.06$\pm$0.10∗ |  | 22.36$\pm$0.11∗ | 21.80$\pm$0.13∗ | 20.78$\pm$0.10∗ | 20.84$\pm$0.10

Note. — Column description: (1) ID number of the object; (2)-(3) right
ascension and declination of radio source; (4) photometric redshifts found by
Ilbert et al. (2009); (5) spectroscopic redshift: $a$ from zCOSMOS survey
(Lilly et al., 2007), $b$ form Magellan survey (Trump et al., 2007), and $c$
from Prescott et al. (2006). The spectroscopic redshift of object 25 (marked
with $*$) is considered incorrect (see Section 6.1); (6)-(7) flux at 1.4 GHz
from NVSS (from http://www.cv.nrao.edu/nvss/NVSSlist.shtml) and FIRST (taken
from Chiaberge et al. 2009) catalogs in mJy.

Note. — List of the filters used for our analysis, the associated telescopes,
the effective wavelengths $\lambda_{eff}$, the PSF full-width half maximum
(FWHM) of the images in each band, and their sensitivities at 5$\sigma$ (mag).

Note. — Results from the analysis of the SEDs of the sample with Hyperz.
Column description: (1) ID number of the object; (2) photometric redshift as
result of Hyperz; all errors are quoted at at 99% of confidence level; (3)
photometric redshifts obtained with 2SPD; (4) source library of the best
template which fits the observed SED: BC from Bruzual & Charlot (2009) and MA
from Maraston (2005); (5) star formation history of the best template: $ssp$
corresponds to a single stellar population and $\tau$ = ’N’ corresponds to the
$\mu$-models with exponentially decaying star formation rate with star
formation timescale $\tau$ = N Gyr; (6) age of the best template in Gyr; (7)
$A_{V}$ associated with the best template; (8) $\chi_{rid}^{2}$ of the best
template; (9) stellar mass and its error of the galaxy in $Log\,M_{\odot}$ as
result of Hyperzmass.

Note. — Results from the analysis of the SEDs with 2SPD. Column description:
(1) ID number of the object; (2) photometric redshift measured with 2SPD;
(3)-(4)-(5)-(6) age in Gyr, AV, flux fraction and mass fraction of the young
stellar population (YSP) at 4800 Årest frame; (7)-(8) age in Gyr and AV of the
old stellar population (OSP); (9) the total stellar mass of the galaxy in M⊙;
(10)-(11) the effective temperature (in K) of the one or two dust components
and their luminosities, Ldust (in units of 109 L⊙); (12)-(13) the infrared
excess luminosity (in erg s-1) defined in the text (Section 6.4) and the
spectral index measured on the infrared excess at 8 and 24$\mu$m; (14) UV
luminosity at 2000 Å in the rest frame in erg s-1. The marginal UV excesses
are marked with a m.

Note. — Column description: (1) ID number of the object; (2) redshift of the
object used throughout the work, with the spectroscopic redshifts marked with
a s; (3)-(4) K-corrected radio luminoisty at 1.4 GHz from NVSS (from
http://www.cv.nrao.edu/nvss/NVSSlist.shtml) and FIRST (taken from Chiaberge et
al. 2009) in erg s-1 Hz-1; (5) classification based on the radio power: low or
high power (LP or HP) radio sources.

Note. — Column description: (1) ID number of the object; (2) CFHT $u^{*}$
magnitude with its error; (3)-(4)-(5)-(6) Subaru $B_{J}$, $g^{+}$, $V_{J}$,
$r^{+}$ magnitudes with their errors; (7) CFHT $i^{*}$ magnitude with its
error; (8) Subaru $i^{+}$ magnitude with its error; (9) HST/ACS $F814W$
magnitude with its error; (10) Subaru $z^{+}$ magnitude with its error; (11)
UKIRT $J$ magnitude with its error; (12) CFHT $K$ magnitude with its error;
(13) NOAO$K_{S}$ with its error. The values marked by $*$ are measured by our
3″-aperture photometry on the images.

Table 7COSMOS multiwavelength counterparts of the sample

ID | $FUV$ | $NUV$ | $IRAC1$ | $IRAC2$ | $IRAC3$ | $IRAC4$ | $MIPS$
---|---|---|---|---|---|---|---
1 |  |  | 54.89$\pm$0.23∗ | 36.50$\pm$0.29∗ | 23.87$\pm$0.97∗ | 14.73$\pm$2.07∗ |  $<$0.15
2 |  |  | 41.44$\pm$0.19 | 45.75$\pm$0.30 | 28.07$\pm$0.99 | 21.8$\pm$2.49 |  $<$0.15
3 |  |  | 13.65$\pm$0.15 | 23.62$\pm$0.25 | 55.88$\pm$0.94 | 166.89$\pm$1.98 | 1.47$\pm$0.02
4 |  |  | 42.97$\pm$0.19 | 48.78$\pm$0.29 | 37.71$\pm$1.04 | 29.34$\pm$2.19 |  $<$0.08
5 |  |  | 66.44$\pm$0.19 | 89.03$\pm$0.29 | 91.57$\pm$0.99 | 63.65$\pm$1.97 | 0.87$\pm$0.02
11 |  |  | 25.34$\pm$0.17 | 29.34$\pm$0.26 | 25.07$\pm$1.00 | 9.64$\pm$1.98 |  $<$0.15
13 |  |  | 49.51$\pm$0.18 | 51.39$\pm$0.26 | 53.24$\pm$0.98 | 75.84$\pm$2.04 | 0.32$\pm$0.03
16 |  |  | 44.29$\pm$0.18 | 32.54$\pm$0.28 | 26.19$\pm$0.94 | 18.39$\pm$2.26 | 0.22$\pm$0.02
18 |  |  | 31.03$\pm$0.17 | 24.94$\pm$0.28 | 18.13$\pm$0.99 | 14.03$\pm$2.15 |  $<$0.61∗
20 |  |  | 63.11$\pm$0.18 | 44.73$\pm$0.27 | 25.12$\pm$0.90 | 17.46$\pm$2.06 |  $<$0.08
22 |  |  | 57.27$\pm$ 0.17 | 64.90$\pm$0.29 | 41.55$\pm$0.89 | 32.84$\pm$2.16 | 0.25$\pm$0.03∗
25 |  |  | 49.17$\pm$0.19 | 55.52$\pm$0.28 | 38.82$\pm$0.99 | 25.08$\pm$2.19 | 0.18$\pm$0.05∗
26 |  |  | 126.74$\pm$0.25 | 108.17$\pm$0.29 | 59.54$\pm$0.98 | 44.75$\pm$1.92 |  $<$0.15
28 |  |  | 32.46$\pm$0.16 | 39.98$\pm$0.26 | 46.73$\pm$0.91 | 40.34$\pm$1.98 | 0.15$\pm$0.02
29 |  |  | 9.81$\pm$0.16 | 10.17$\pm$0.25 | 6.50$\pm$1.00∗ |  $<$7.38∗ |  $<$0.15
30 |  |  | 74.26$\pm$0.19 | 58.67$\pm$0.28 | 33.32$\pm$0.91 | 22.76$\pm$2.05 |  $<$0.15
31 |  | 25.15$\pm$0.24 | 47.86$\pm$0.16 | 34.60$\pm$0.25 | 24.98$\pm$0.81 | 19.36$\pm$1.78 |  $<$0.15
32 |  |  | 16.97$\pm$0.15 | 32.41$\pm$0.24 | 58.17$\pm$0.95 | 98.81$\pm$1.81 | 0.26$\pm$0.02
34 |  |  | 22.80$\pm$0.44∗ | 26.73$\pm$0.48∗ | 25.44$\pm$1.10∗ | 12.07$\pm$2.40∗ |  $<$0.26∗
36 |  |  | 49.50$\pm$0.16 | 43.80$\pm$0.27 | 30.24$\pm$0.87 | 21.89$\pm$1.97 | 0.15$\pm$0.04∗
37 |  | 23.55$\pm$0.15 | 88.32$\pm$0.22 | 123.39$\pm$0.32 | 158.41$\pm$1.10 | 249.52$\pm$2.18 | 1.58$\pm$0.39
38 |  |  | 30.47$\pm$0.17 | 28.00$\pm$0.28 | 29.99$\pm$0.95 | 24.94$\pm$2.35 | 0.13$\pm$0.06∗
39 |  |  | 74.19$\pm$0.21 | 66.35$\pm$0.32 | 34.74$\pm$0.96 | 25.29$\pm$2.40 |  $<$0.08
52 |  | 24.97$\pm$0.21 | 59.84$\pm$0.21 | 41.39$\pm$0.28 | 36.37$\pm$1.09 | 28.67$\pm$2.04 | 0.16$\pm$0.02
70 |  |  | 17.58$\pm$0.16 | 20.57$\pm$0.26 | 29.83$\pm$1.04 | 21.32$\pm$2.17 | 0.13$\pm$0.03∗
202 |  |  | 24.81$\pm$0.14 | 22.02$\pm$0.24 | 15.90$\pm$0.84 | 6.46$\pm$1.87 |  $<$0.15
219 |  |  | 98.71$\pm$0.23 | 86.12$\pm$0.32 | 48.50$\pm$1.02 | 38.60$\pm$2.26 | 0.12$\pm$0.01
224 |  |  | 48.09$\pm$0.17 | 41.65$\pm$0.26 | 25.56$\pm$0.88 | 18.53$\pm$2.07 |  $<$0.15
226 |  |  | 16.94$\pm$0.15 | 30.93$\pm$0.28 | 53.89$\pm$0.89 | 104.92$\pm$2.16 | 0.54$\pm$0.02
228 |  |  | 39.72$\pm$0.18 | 47.82$\pm$0.27 | 36.06$\pm$0.99 | 23.46$\pm$2.03 |  $<$0.15
234 |  |  | 50.19$\pm$0.19 | 45.78$\pm$0.26 | 33.95$\pm$1.03 | 14.69$\pm$2.00 |  $<$0.15
236 |  |  | 99.23$\pm$0.20 | 146.61$\pm$0.36 | 224.02$\pm$1.04 | 324.18$\pm$2.50 | 0.98$\pm$0.04
258 | 24.97$\pm$0.15 | 24.32$\pm$0.09 | 84.11$\pm$0.23 | 59.75$\pm$0.30 | 37.90$\pm$1.06 | 31.78$\pm$2.12 |  $<$0.15
285 |  | 24.85$\pm$0.18 | 34.06$\pm$0.17 | 33.78$\pm$0.26 | 33.68$\pm$0.95 | 32.70$\pm$2.03 | 0.13$\pm$0.03∗

Note. — Column description: (1) ID number of the object; (2)-(3) GALEX FUV and
NUV magnitudes with their errors; (4)-(5)-(6)-(7) Spitzer/IRAC 4-channel (3.6,
4.5, 5.8, and 8.0 $\mu$m) fluxes with their errors; (8) Spitzer/MIPS flux at
24$\mu$m with its error. The values marked by $*$ are measured by our
3″-aperture photometry on the images.