EDP Sciences
Free Access
Issue
A&A
Volume 607, November 2017
Article Number A89
Number of page(s) 22
Section Extragalactic astronomy
DOI https://doi.org/10.1051/0004-6361/201730866
Published online 20 November 2017

© ESO, 2017

1. Introduction

The star formation history (SFH) in the Universe is one of the key constraints to understand the evolution of galaxies. The combination of various tracers (Hα, far UV, far infrared and millimeter) was successfully used in the last 20 yr to measure the star formation rate density (SFRD) up to very high redshift (z ~ 8, see Madau & Dickinson 2014, for a review). At z ≥ 2−3, building complete spectroscopic samples becomes very challenging and continuum emission is mainly used to derive star formation rates (SFR). Consequently, the prime tracer of recent star formation is the redshifted far-UV emission from young stars. However, even at early epochs, massive galaxies have already formed a large amount of dust and UV light is thus absorbed (e.g., Takeuchi et al. 2005; Heinis et al. 2014). Two main approaches can then be used to derive the intrinsic SFR: correct the UV absorption using the UV spectral slope as a proxy of attenuation (e.g., Calzetti et al. 2000) or directly detect the reprocessed UV light emitted by dust in the far-infrared and millimeter (Kennicutt 1998).

Far-infrared and submillimeter observations are challenging because of the limited angular resolution of the instruments. The deepest observations of the most modern single-dish instruments are limited by the confusion, that is, the blending of sources in the same beam of the instrument (e.g., Dole et al. 2003). Only the brightest galaxies emerge from the confusion and can be extracted individually from far-infrared and submillimeter maps. However, because of the large beam of the single-dish instruments, their measured flux density can be contaminated by their fainter neighbors. Indeed, follow-up observations of the brightest 850 μm sources at high-resolution with ALMA revealed that a large fraction are multiple sources (e.g., Karim et al. 2013; Hodge et al. 2013). Because of this, the flux density distributions measured with single-dish instruments and interferometers such as ALMA (Karim et al. 2013; Simpson et al. 2015) strongly disagree.

Also, the Herschel space observatory (Pilbratt et al. 2010) has a limited angular resolution and could be affected by similar effects. However, it is very difficult to verify because interferometric follow-up observations are not possible at the high frequencies of the Herschel observations. Consequently, other approaches such as modeling must be used to explore possible biases induced by the angular resolution on Herschel number counts. In particular, we have to understand why the number density of red Herschel sources found in the extragalactic surveys is almost an order of magnitude higher than that predicted by the models (Asboth et al. 2016; see also Dowell et al. 2014; Ivison et al. 2016).

In addition to studies of bright sources above the confusion limit, various advanced techniques were developed to probe galaxy populations in the confusion such as the stacking method (e.g., Dole et al. 2006; Marsden et al. 2009; Béthermin et al. 2010a; Viero et al. 2013a), P(D) measurements (e.g., Condon 1974; Patanchon et al. 2009; Glenn et al. 2010), or source extraction using position priors coming from shorter wavelengths (e.g., Magnelli et al. 2009; Béthermin et al. 2010b; Roseboom et al. 2010; Hurley et al. 2017). These methods can also be biased by the contamination of the measured flux by faint clustered sources.

Simulations were developed to test these possible biases (e.g., Fernandez-Conde et al. 2008), but the clustering of infrared galaxies at high redshift was poorly constrained at that time. Important progress has been made recently. In particular, Planck and Herschel measured cosmic infrared background (CIB) anisotropies with an unprecedented precision (Planck Collaboration XVIII 2011; Amblard et al. 2011; Planck Collaboration XXX 2014; Viero et al. 2013b). Their modeling showed that the typical mass of the dark-matter halos hosting the bulk of the obscured star formation is almost constant and around 1012M up to z ~ 3 (e.g., Béthermin et al. 2012b, 2013; Viero et al. 2013b; Planck Collaboration XXX 2014; Wu et al. 2016). In addition, clustering studies of bright high-redshift far-infrared and millimeter galaxies showed they are hosted by massive halos (~ 1013M, e.g., Farrah et al. 2006; Weiß et al. 2009; Magliocchetti et al. 2014; Béthermin et al. 2014; Wilkinson et al. 2017). These massive halos are strongly clustered. The impact of clustering on the extraction of sources from confusion-limited surveys might be stronger than predicted by pre-Planck and Herschel simulations, which assumed a weaker clustering.

It is thus timely to develop new simulations that are able to reproduce simultaneously the far-infrared and millimeter observations at various angular resolutions. These simulations must include clustering and take into account all the lessons learnt from Herschel and ALMA. On the one hand, Hayward et al. (2013b) built a simulation based on abundance matching, but this analyzes only the galaxy populations selected at 850 μm (see also Muñoz Arancibia et al. 2015; Cowley et al. 2015). On the other hand, Schreiber et al. (2017) built a simulation of the panchromatic properties of galaxies, but did not include a physical clustering model. Our new simulation combines the strengths of these two approaches and accurately reproduces spectral and spatial properties of galaxies and CIB anisotropies. In this paper, we focus on the continuum properties of galaxies and the effect of angular resolution from 70 μm to 1.2 mm. In a future paper, we will introduce the (sub)millimeter line ([CII], [NII], [CI], CO...) properties of galaxies, discuss the perspectives for (sub)millimeter intensity mapping and test methods of line deblending.

Our simulation is based on the Bolshoi-Planck simulation (Klypin et al. 2016; Rodríguez-Puebla et al. 2016), from which a lightcone covering 2 deg2 was produced. We populate the dark-matter halos using an abundance-matching technique (e.g., Vale & Ostriker 2004). The luminous properties of the galaxies are derived using an updated version of the 2SFM (2 star-formation modes) model (Sargent et al. 2012; Béthermin et al. 2012a, 2013). This model is based on the observed evolution of the main sequence of star forming galaxies (e.g., Noeske et al. 2007; Elbaz et al. 2007; Daddi et al. 2007), that is, a SFR-M correlation evolving with redshift, and the observed evolution of the spectral energy distributions (SEDs) with redshift. In the new version of the model, we take into account the increase of dust temperature in main sequence galaxies recently measured from z = 2 to z = 4 (Béthermin et al. 2015a), extending the increase found from z = 0 to z = 2 by Magdis et al. (2012). We also include the latest calibration of the evolution of the main sequence (Schreiber et al. 2015).

In Sect. 2, we present the ingredients of our simulation and discuss its limitations. We compare our results with observed number counts and discuss the effects of resolution in Sect. 3. We then discuss the redshift-dependent observables and the consequences on the obscured star formation history (Sect. 4). We then show the significant impact of clustering on the pixel histogram of the Herschel maps, also known as P(D), and check that our model correctly reproduces the CIB anisotropies measured by Herschel and Planck (Sect. 5). Finally, we discuss the existence of the red sources found by Herschel surveys (Sect. 6).

We assume a Planck Collaboration XIII (2016) cosmology and a Chabrier (2003) initial mass function (IMF). The products of our simulation, called SIDES (Simulated Infrared Dusty Extragalactic Sky)1.

2. Ingredients of the simulation

This section describes the ingredients used to build our simulated sky, namely,

  • the dark-matter lightcone, which is the starting point(Sect. 2.1);

  • the stellar mass function (Sect. 2.2);

  • the abundance-matching procedure used to populate the dark-matter halos with galaxies (Sect. 2.3);

  • our recipe to split galaxies into a star forming and a passive population (Sect. 2.4);

  • our method to derive a SFR for each galaxy (Sect. 2.5);

  • the assignment of SEDs to our simulated galaxies (Sect. 2.6);

  • the implementation of strong and weak lensing (Sect. 2.7).

Finally, in Sect. 2.8, we discuss the limitations of our simulation. We homogenized the cosmology used in the dark matter simulation and in the observed stellar mass functions. Our method is described in Appendix A.

2.1. Dark matter simulation and lightcone catalog

We use the publicly available halo catalogs from the Bolshoi-Planck simulation (Rodríguez-Puebla et al. 2016)2. The simulation has a volume of (250 h-1 Mpc)3, with a dark-matter particle mass of 1.5 × 108h-1M. The cosmological parameters are compatible with Planck Collaboration XIII (2016): h = 0.678, σ8 = 0.823, ΩΛ = 0.693, ΩM = 0.307, Ωb = 0.048, ns = 0.96.

Dark matter halos are identified by the phase-space halo finder Rockstar (Behroozi et al. 2013b). We use the halo mass M200, which is defined by the radius within which the spherical overdensity is 200 times the critical density of the Universe. We only use halos with mass above 1010M, which have more than 50 dark matter particles. We have explicitly verified that above 1010M, the halo mass function from our simulation agrees with the analytic halo mass function.

Using the simulation snapshots at different redshifts, we construct a lightcone catalog of 1.4 deg × 1.4 deg, 0 <z< 10, corresponding to a comoving volume of 0.17 Gpc3, approximately three times the volume of the Bolshoi-Planck simulation. In a lightcone catalog, each object is at a cosmic distance that corresponds to the cosmic time that it emits light. The simulation outputs are saved at discrete time steps, and we use snapshots approximately spaced by Δz = 0.25. We have explicitly checked that the structure in such a narrow redshift bin has negligible evolution and can be represented by the same snapshot.

To construct the lightcone, we replicate the box in all three dimensions, using the periodic boundary condition inherent in the simulations. Since our lightcone catalog has a pencil-beam geometry, we use a “slanted” line of sight to reduce the repeated structure; that is, the line of sight is not parallel to any of the axes or diagonals of the box. Specifically, we first rotate the box by 10° along the y-axis and another 10° along the z-axis. We then transform the Cartesian coordinates into the equatorial coordinates, following the convention of astropy. The distance of an object from the observer is converted into the cosmological redshift, and we add to the redshift the peculiar velocity along the line-of-sight. For more details about the constructions of the lightcone, we refer to Merson et al. (2013).

thumbnail Fig. 1

Left: evolution of the stellar mass function used in our simulation with redshift. The construction of this stellar mass function is explained in Sect. 2.2. The solid line is the SMF computed from the catalog in a bin Δz = 0.4. The dashed line is the analytic SMF at the mean redshift of the bin. The cut at low mass (M< 107M) is due to the halo mass limit of the dark-matter simulation. Right: relation between the halo mass and the stellar mass from our abundance-matching procedure (see Sect. 2.3). The error bar indicates the 0.2 dex scatter on this relation.

Open with DEXTER

2.2. Stellar mass function

In our simulation, the stellar mass function (SMF) is the starting point from which we generate all the properties of the galaxies. Similarly to the approach presented in Bernhard et al. (2014), we assume that it can be described by a double Schechter function (e.g., Baldry et al. 2012): (1)where is the characteristic mass of the knee of the SMF, and are the normalization of the two components, and α1 and α2 the power-law slopes at low mass. We use the same functional representation of the SMF at all redshifts to avoid discontinuities of its evolution with redshift.

The evolution with redshift of the parameters described above is based on the observations. We use the data points of Kelvin et al. (2014) in the GAMA field in the local Universe, Moutard et al. (2016) from the VIPERS survey up to z = 1.5, Davidzon et al. (2017) in the COSMOS field from z = 1.5 to z = 4, and Grazian et al. (2015) at z> 4. Grazian et al. (2015) uses a simple Schechter function. At z> 4, to ensure a smooth transition with smaller redshifts at which a double Schechter function is used, we fix to 0 and use the Φ and α of the single Schechter function for the second component. We connect the data points (taken at the center of the redshift bins of the authors) using a linear interpolation of each parameter (log(), , log(), α1, α2) as a function of (1 + z). We chose to use and log() to avoid problems with the log where is fixed to zero and to avoid negative values at z> 7, respectively. The stellar mass function of the galaxies used to generate our simulation is shown in Fig. 1. Below 108M, the number density at fixed mass no longer evolves monotonically with redshift. This unphysical behavior is caused by the uncertainties on the low-mass slope of the observed data we used. This is a limitation of our empirical approach. However, these low-mass sources have a small impact on our simulation, since they emit only 4% of the infrared luminosity. Thus, we have chosen to keep these sources in the simulation, since they contribute to confusion noise.

2.3. Abundance matching

To assign stellar mass to dark matter halos and subhalos, we perform subhalo abundance matching between the halo catalogs and the stellar mass functions described above. The basic idea of abundance matching is to assign higher stellar mass to more massive halos or subhalos, either monotonically or with some scatter, according to the number densities of the objects in the Universe (e.g., Vale & Ostriker 2004; Shankar et al. 2006; Behroozi et al. 2013a; Moster et al. 2013). In this work, instead of mass, we use the peak circular velocity vpk of dark matter halos and subhalos to perform the abundance matching, since vpk is known to be more tightly correlated with stellar mass (e.g., Reddick et al. 2013).

We assume that the stellar mass has an intrinsic scatter of 0.2 dex at a given vpk, which is required for the resulting galaxy catalog to reproduce the observed galaxy clustering (Reddick et al. 2013). The input SMF (Eq. (1)) is deconvolved into a stellar mass function without the intrinsic scatter on the stellar versus halo mass relation; this deconvolved stellar mass function is then used to match the number density of halos monotonically. We use the implementation by Y.-Y. Mao3, and we refer the readers to Behroozi et al. (2010, 2013a) and Wu & Doré (2017) for the detailed implementation. The left panel of Fig. 1 demonstrates that the stellar mass functions resulted from this abundance-matching calculation (solid curves) recover the input stellar mass functions (dashed curves). There is a slight tension at low mass in some redshift bins caused by the evolution of the SMF inside a redshift bin. We also observe a sharp cut below 107M for 0 <z< 0.4, which is caused by the halo mass limit of the simulation. Since halo and stellar masses are correlated, this also implies a low-mass cut in the stellar mass function.

The right panel of Fig. 1 shows the stellar mass–halo mass relation resulting from the abundance-matching calculation.

2.4. Fraction of star forming galaxies

thumbnail Fig. 2

Fraction of star-forming galaxies versus stellar mass at various redshifts. The data points are from Davidzon et al. (2017). The solid line is a fit by the parametric form described in Sect. 2.4.

Open with DEXTER

We will draw randomly galaxy properties from their stellar mass and redshift using the prescriptions of the 2SFM formalism (Sargent et al. 2012; Béthermin et al. 2012a), which applies only to star forming galaxies. First, we have to estimate the probability of a galaxy at a given M and redshift to be star forming. Accordingly, we split the galaxies in our simulation in two populations: passive galaxies, which have a negligible star formation, and star forming galaxies. We used the observed evolution of the star forming fraction by Davidzon et al. (2017) to derive this fraction. In this work, the authors classified the galaxies as star forming or not using their position in the (NUV-r) versus (rK) color diagram (Arnouts et al. 2013; Ilbert et al. 2013). We fit their results with the following parametric form (see also Fig. 2): (2)where fQ,0(z) is the fraction of passive galaxies at low mass (MMt(z)). The fraction is higher at low redshift, where a significant fraction of low-mass galaxies in dense environments are passive. In Eq. (2), Mt(z) is the stellar mass of the transition between passive and star forming galaxies, and σSF(z) the width of this transition. These three quantities evolve with redshift. Their evolution is parametrized in the following way: This parametric form provides an excellent fit of the measurements (reduced χ2 of 0.82). The best fit parameters are fQ,0,z = 0 = 0.1017, log 10(Mt,z = 0) = 10.53, σSF,z = 0 = 0.8488, α1 = 0.2232, α2 = 0.0913, β1 = 0.0418, β2 = −0.0159, and γ = −1.039.

This approach neglects environmental effects, since it depends only on the stellar mass and redshift. Our simulation is optimized for field galaxies, cosmic infrared background, and intensity mapping studies. Cosmic infrared background and intensity mapping are dominated by central galaxies and are thus not severely affected by these effects (e.g., Béthermin et al. 2013). The limitations implied by this simplification are discussed in Sect. 2.8.

2.5. Star-forming properties

We assume that only galaxies classified as star-forming have far-infrared and millimeter outputs. In passive galaxies, some residual emission of cirrus heated by the old stellar populations has been observed. However, at a given stellar mass, these galaxies usually have infrared luminosities lower by at least one order of magnitude than galaxies on the main sequence (e.g. Viero et al. 2013a; Amblard et al. 2014; Man et al. 2016; Gobat et al. 2017b,a). Neglecting their infrared outputs is thus a fair assumption.

The SFR of star forming galaxies is derived using the 2SFM formalism (Sargent et al. 2012; Béthermin et al. 2012a). The first step is to compute the mean SFR of sources from the measured evolution of the main sequence of star forming galaxies, written afterwards SFRMS. Schreiber et al. (2015) measured the evolution of this main sequence up to z = 4 and proposed the following parametric description: (6)with m0 = 0.5, a0 = 1.5, a1 = 0.3, m1 = 0.36, a2 = 2.5. In Béthermin et al. (2012a), we assumed a simple power law for the main sequence at a given redshift. In addition, at z> 2.5, there was no evolution of sSFR, that is, SFR /M, with z at fixed M. In this updated version (Eq. (6)), the SFR decreases sharply at high M and sSFR continues to evolve at higher redshift. This rising sSFR was already discussed in Béthermin et al. (2013), since it reproduces the CIB anisotropies better. The Schreiber et al. (2015) formula is fitted on observations at z> 0.5 and sSFR is too high at lower redshift. To correct for this offset, we applied a  dex offset to the Schreiber et al. (2015) formula at z< 0.5. The detailed explanations are provided in Appendix B.

Star forming galaxies are not all on the main sequence. In this paper, a starburst is defined as a positive outlier of the main sequence (e.g., Elbaz et al. 2011; Sargent et al. 2012). Following Béthermin et al. (2012a), the fraction of starburst does not vary with stellar mass; it grows linearly with redshift from 1.5% at z = 0 to 3% at z = 1 and stays flat at higher redshift. We randomly drew a main sequence or a starburst galaxy using this probability.

The main sequence is of course not a perfect correlation and it has a non-negligible scatter. We followed a procedure similar to Béthermin et al. (2012a) to distribute the galaxies around the main sequence. We randomly drew the SFR of each source using a log-normal distribution in agreement with the observational results (e.g., Rodighiero et al. 2011). Following Sargent et al. (2012), the Béthermin et al. (2012a) model assumed a width of 0.15 and 0.2 dex for main sequence galaxies and starbursts, respectively. However, more recent measurements by Schreiber et al. (2015) and Ilbert et al. (2015) found a slightly higher width of 0.3 dex (see also Sun et al. 2016). We thus use this updated value in our simulation. The distribution of main sequence galaxies is centered on 0.87 SFRMS and 5.3 SFRMS for the starbursts (Schreiber et al. 2015). Since a log-normal distribution centered on 1 has a mean value above unity, the center of main sequence is set to 0.87 SFRMS in order to have the correct mean SFR (see Schreiber et al. 2015; and Ilbert et al. 2015, for more explanations).

Using follow-up observations of submillimeter sources detected by single-dish telescopes with interferometers, Karim et al. (2013) showed that the brightest of these sources have multiple components (see also Simpson et al. 2015). They found that the bright end of the number counts at 850 μm were significantly overestimated and that none of the single components have a SFR significantly above 1000 M/yr. The SFR distribution of 870 μm-selected galaxies measured by da Cunha et al. (2015) drops strongly above 1000 M/yr. A rapid drop of the number density at was also found by high-resolution radio observations (Barger et al. 2014, 2017). For simplicity, we implemented a sharp SFR limit at 1000 M/yr. The SFR of each galaxy is redrawn until it is lower than this limit. Consequently, the sSFR distribution of the most massive galaxy populations is truncated at high sSFR. Wide surveys found some rare sources with a higher SFR suggesting this is not a sharp limit (e.g., Riechers et al. 2013; Tan et al. 2014; Ma et al. 2016). However, using a sharp limit rather than an exponential cut of the sSFR distribution is a reasonable assumption considering the small size of our field. The impact of this SFR limit on the number counts is discussed in Sect. 3.5. So far, the physical origin of this SFR cut is not totally clear. These objects could be Eddington-limited starbursts limited by the radiative pressure (e.g., Thompson et al. 2005). This could also be explained by the weaker boost of star formation induced by mergers in gas-rich systems (e.g., Fensch et al. 2017).

2.6. Spectral energy distributions and continuum fluxes

We then assigned a SED to each of our sources to derive their flux densities in a large set of instrument filters from their total infrared luminosity (LIR). LIR is directly derived from SFR using the Kennicutt (1998) conversion factor (1.0 × 10-10M/yr/L after converting to Chabrier 2003 IMF). We use Magdis et al. (2012) SED library. The shape of the SEDs depends on the galaxy type (main sequence or starburst) and on the U parameter, that is, the mean intensity of the radiation field. This parameter is strongly correlated with the dust temperature (e.g., Dale & Helou 2002). It evolves with redshift for main sequence galaxies (see Fig. 3). In Béthermin et al. (2012a), we had no data above z = 2 and we assumed a flattening at z> 2, since it provides a better agreement with the observed submillimeter number counts. Two new observational inputs motivated us to update the evolution of U in our simulation. In Béthermin et al. (2015a), we measured that U continues to rise at z> 2 using a stacking analysis. In addition, Karim et al. (2013) also showed that submillimeter number counts were overestimated because of blending effects (see the discussion in Sect. 3.4). The previous measurements favored a scenario with no evolution of U at z> 2, because it was producing colder SEDs and consequently higher submillimeter counts. This is no longer true with the new number counts that are observed to be lower.

For main sequence galaxies, we used UMS ⟩ (z = 0) = 5 as in Béthermin et al. (2012a). An evolution in (1 + z)α does not fit very well the observational data from Béthermin et al. (2015a) with an overly sharp decrease with decreasing redshift at z< 0.5. This artificially low U at low redshift is responsible of an excess of the bright number counts at 160 μm. We thus used another parametric form, which fits better the observational data: (7)with αU = 0.25. Following Béthermin et al. (2015a), we use a constant USB ⟩ = 31. However, at z> 3, this would lead to starbursts colder than main sequence galaxies. This behavior could be considered as unphysical and we thus assumed UMS ⟩ = ⟨ USB at higher redshift. At z> 4, we have no constraints on an evolution of UMS. Extrapolating this behavior up to z ~ 10 would imply unphysically high values. We thus assume a plateau in the high redshift regime (z> 4). In addition to this mean evolution, we also included a 0.2 dex scatter on U following Magdis et al. (2012).

thumbnail Fig. 3

Evolution of the mean intensity of the radiation field U (see Magdis et al. 2012) that is used in the current updated model. The evolution of main sequence galaxies and starbursts are plotted in blue (solid line) and red (dashed line), respectively. Also shown are the measurement of Béthermin et al. (2015a, triangles for main sequence galaxies and squares for starbursts), Magdis et al. (2012, diamonds), da Cunha et al. (2010, asterisk), and Ciesla et al. (2014, circle). This figure is adapted from Béthermin et al. (2015b).

Open with DEXTER

2.7. Magnification by lensing

Gravitational lensing can have a non-negligible impact on the bright submillimeter number counts, because of their steepness (Negrello et al. 2007, 2010, 2017; Béthermin et al. 2011, 2012a; Lapi et al. 2011, 2012; Vieira et al. 2013; Wardlow et al. 2013). At 350 and 500 μm, this effect is maximal around 100 mJy, where ~20% of the sources are lensed. Our simulation of a 2 deg2 field contains only six sources brighter than this threshold. The lensing has thus a relatively weak effect on the total number counts; however, it has a non-negligible impact on the number of bright red sources (see Sect. 6), since the fraction of lensed sources is higher at high redshift. It is therefore important to consider lensing.

For each source of our simulation, we randomly drew the magnification μ. The determination of the magnification does not include any spatial information (see Sect. 2.8 for a discussion about this approximation). For the strong lensing (μ> 2), we used the probability distribution of Hezaveh & Holder (2011) used also in Béthermin et al. (2012c), which depends only on the redshift. We also included a simplified weak lensing model for the other sources. We randomly drew their magnification from a Gaussian, whose width and mean value are derived from Hilbert et al. (2007, their Figs. 1 and 2).

2.8. Limits of our simulation

Our simulation is based on the observed evolution of star forming galaxies and aims to accurately reproduce current observations of the far-infrared and (sub)millimeter Universe. However, the current version of this simulation has several limitations, which should be kept in mind while comparing it with observations. Our simulation is based on a single 2 deg2 field. Since it is based on a dark-matter simulation, it is thus affected by cosmic variance beyond simple Poisson fluctuations and can contain under- or overdensities at some specific redshift.

Our abundance-matching procedure assumes that the stellar mass of a galaxy is associated with vpk (proxy of the potential well of dark matter halos or subhalos), with some scatter. During our abundance-matching procedure, we implicitly assume that main halos and subhalos follow the same relation. In addition, the probability of a galaxy to be passive at a given z depends only on its stellar mass. Our simulation thus neglects the environmental quenching observed in the most massive halos (Mhalo> 1014M) at z< 1 (e.g., Peng et al. 2010). There are only 26 such halos in our simulation. Moreover the contribution of these massive structures to the star formation density is small (Popesso et al. 2015). This approximation should thus only be a problem if the simulation is used to study low-redshift overdensities.

Our description of the lensing depends only on the redshift and ignores the position of foreground sources. This treatment is thus inconsistent with the large-scale structures of our simulation. Overdensities of low-z galaxy populations are associated with massive halos, which can strongly magnify high-redshift sources (e.g., Wang et al. 2011; Welikala et al. 2016). A full consistent treatment of lensing is beyond the scope of this paper and we thus decided to have a purely probabilistic treatment of lensing magnification. The impact of this simplification should be small on most of the statistics, but spatial correlations between bright-lensed sources and their neighbors could be significantly affected.

Finally, at z> 4, our simulation relies on extrapolations of relations calibrated at lower redshift. The SEDs used in our simulation evolve only up to z = 4. At higher redshift, we assume no evolution due to the lack of constraints. Potentially, SEDs could become even warmer because of the effect of CMB (da Cunha et al. 2013; Zhang et al. 2016). However, the temperature of our SEDs is higher than the dust temperature assumed in these studies (~40 K versus ~15 K). The CMB effect should thus be much smaller than estimated in these studies, but it might be non-negligible for the highest-redshift objects of our simulation. The evolution of sSFR in our simulation is based on the Schreiber et al. (2015) relation, which is derived from z< 4 data. The scatter on the main sequence is also assumed to be constant with mass and redshift, since there is currently no evidence of the contrary. Finally, the evolution of the parameters of the stellar mass function are extrapolated at z> 6. The predictions of our simulation at z> 4 should thus be taken with caution.

Our simulation currently contains only the far-infrared and millimeter observables and we thus assumed in Sect. 2.6 that LIR traces the total star formation. However, in low-mass and high-redshift galaxies, the fraction of UV photons escaping the galaxies can be non-negligible. The impact of neglecting unobscured star formation on the infrared observables was discussed extensively in Bernhard et al. (2014). They showed that the scatter of the infrared excess (IRX = LIR/LUV) has a negligible impact on infrared observables. In contrast, the increasing IRX with increasing M implies that low-mass objects have a smaller fraction of their UV reprocessed by dust and the faint-end slope of the number counts should be slightly steeper if we include the UV (Béthermin et al. 2012a). The impact of these lower number counts on the confusion noise is small (<5%), since the shot noise is proportional to dN/ dSS2dS, where S is the flux density and dN/ dS are the number counts.

3. Number counts and multiplicity of sources detected by single-dish instruments

In this section, we demonstrate that our simulation is able to reproduce the observed number counts, when the effects of angular resolution on source extraction are properly taken into account. We also discuss in particular the multiplicity of Herschel/SPIRE and NIKA2 sources and the bias caused by clustering on stacking measurements.

3.1. Simulating the observational process

Karim et al. (2013) showed that the 850 μm sources found by single-dish telescopes are often blends of several sources. The same phenomenon could also impact other single-dish observations and especially Herschel. We thus compare the measurements with both the intrinsic number counts from our simulated catalog and number counts extracted from simulated maps.

We built Herschel simulated maps from our simulated catalog. We used Gaussian beams with full widths at half maximum (FWHM) of 5.5, 6.5, 11, 18.2, 24.9, and 36.3 arcsec at 70, 100, 160, 250, 350, and 500 μm, respectively, corresponding to the measured size of the Herschel beams. We did not include any instrumental noise, since we are only interested in the effect of angular resolution. The faint sources in the simulated map are responsible for the confusion noise. We measured a confusion noise of 6.0, 6.5, and 6.0 mJy at 250, 350, and 500 μm, respectively. This is compatible at 2σ with the measurements of Nguyen et al. (2010), who found 5.8 ± 0.3, 6.3 ± 0.4, and 6.8 ± 0.4, respectively.

We extracted the sources from Herschel maps using FASTPHOT (Béthermin et al. 2012c). This routine uses source positions from another wavelength as a prior to deblend their flux. A large fraction of Herschel catalogs were produced using the position of 24 μm sources as a prior (e.g., Roseboom et al. 2010; Berta et al. 2011; Béthermin et al. 2012a; Magnelli et al. 2013).

Photometry routines using positional priors are not converging when too many sources are located in the same beam because of degeneracies. We thus kept only the brightest 24 μm sources in a 0.5 FWHM of radius in our list of prior position. Finally, even if catalogs extracted using position priors are not affected by flux boosting, they are still affected by the Eddington bias. This bias appears when a steep distribution is convolved by measurement uncertainties (see Béthermin et al. 2012c). We estimated the correction factor following Béthermin et al. (2012a). We start from the flux distribution measured in the map. We then add a random Gaussian noise to each flux and compare the flux distribution before and after adding this noise. The photometric noise is estimated using the standard deviation of the residual map.

Contrary to λ ≤ 500μm, we cannot use 24 μm priors to extract the sources at 850 μm and 1.2 mm, since the 24 μm is no longer probing the dust emission (>8 μm rest-frame) at the typical redshift of the sources detected at these wavelengths (see Sect. 4). Thus, we extracted blindly the >5σ peaks in our simulated map. This task is relatively easy, since we have no instrumental noise in our simulations. We then measured the flux density of the detected sources using FASTPHOT and deboosted the fluxes following Geach et al. (2017).

Single-dish observations were performed with various angular resolutions. We chose to use the resolution of the James Clerk Maxwell Telescope (JCMT, 15-m diameter). This choice was guided by the fact that the most recent single-dish surveys at 850 μm and 1.1 mm/1.2 mm were performed with this telescope. Usually, ground-based (sub)millimeter maps are convolved by the Gaussian of the size of the beam before extracting the sources. This technique is optimal to extract point sources in noise-limited maps, but it increases the confusion and blending problems. The convolved map is usually called beam-smoothed map. We thus produced beam-smoothed simulated maps, with an effective FWHM after convolution of 21 and 26 arcsec at 850 μm and 1.1 mm, respectively. We used the beam at 1.1 mm instead of 1.2 mm, since the most accurate number counts were measured at this wavelength with the AzTEC camera (Scott et al. 2012).

thumbnail Fig. 4

Differential number counts from 70 μm to 500 μm. The counts are multiplied by to reduce the dynamic range of the plot and to highlight the plateau at high flux densities where the Euclidian approximation is valid (e.g., Planck Collaboration Int. VII 2013). The solid black line is the prediction from our simulated catalog. The red line is derived from the extraction of sources in simulated Herschel maps using a method similar to that used by Béthermin et al. (2012a). The error bars on the prediction of the simulation are derived assuming Poisson statistics. The source extraction procedure is limited by confusion and only bright sources can be extracted reliably. From 70 μm to 160 μm, the light blue circles, the dark blue diamonds, and the green triangles are the Spizter counts of Béthermin et al. (2010a), Herschel/PEP counts of Berta et al. (2011), and Herschel/GOODS counts of Magnelli et al. (2013), respectively. From 250 μm to 500 μm, the light blue circles, the dark blue diamonds, the green triangles, turquoise squares, and purple asterisks are the Herschel/SPIRE measurements of Béthermin et al. (2012c), Oliver et al. (2010), Glenn et al. (2010), Clements et al. (2010), and Valiante et al. (2016), respectively. The SCUBA2 measurements at 450 μm of Chen et al. (2013), Casey et al. (2013), Geach et al. (2013), and Zavala et al. (2017) are shown using gold down-facing triangles, orange right-facing triangles, red open circles, and brown stars, respectively.

Open with DEXTER

thumbnail Fig. 5

Integral number counts at 850 μm and 1.2 mm. The black solid line is the intrinsic counts from our simulated catalog. The red solid line is the result of source extraction from the simulated map of single-dish instrument (see Sect. 3.3). The number counts extracted from high-resolution interferometric data (int.) are colored in blue and purple (Oteo et al. 2016; Karim et al. 2013; Simpson et al. 2015; Fujimoto et al. 2016; Hatsukade et al. 2013; Aravena et al. 2016), while we use orange for single-dish (SD) results (Coppin et al. 2006; Geach et al. 2017; Lindner et al. 2011; Scott et al. 2012).

Open with DEXTER

3.2. Spitzer and Herschel number counts

The comparison between our simulation and the observed number counts is presented in Fig. 4. The intrinsic number counts in our simulated catalog (black solid lines) agree well with the data overall. However, there are some tensions at some specific wavelengths and flux regimes. The number counts at 70 μm in our simulation (both intrinsic and extracted from the simulated maps) are 2σ high at the bright end compared with Béthermin et al. (2010a)Spitzer measurements, but agree at 1σ with the Hercshel/PACS measurements of Berta et al. (2011). The intrinsic faint-end slope (<2 mJy) of the PACS number counts (70, 100, and 160 μm) is less steep in our simulation than in the observations, but the number counts recovered after a source extraction in our simulated map (red solid lines, Sect. 3.1) agree with the observations. Jin et al. (in prep.) also found that the published PACS number counts are underestimated using advanced source extraction techniques.

Below 5 mJy, the intrinsic Herschel/SPIRE number counts (250, 350, and 500 μm) are 2σ higher than the constraints derived by stacking by Béthermin et al. (2012c) and by P(D) analysis by Glenn et al. (2010). These constraints come essentially from the GOODS fields, which are deep but small and thus strongly affected by the cosmic variance. For instance, only the S< 5 mJy data points of Béthermin et al. (2012c) come from GOODS-N. The S> 5 mJy data points are dominated by COSMOS, which probes a much larger volume than the GOODS fields, and agree well with our simulation at 250 μm. In addition, the pixel histograms of the COSMOS maps, that is, P(D), which is very sensitive to the number of faint sources (see Sect. 5), agree well with our simulation.

The main disagreement between intrinsic and measured number counts is located between 5 mJy and 50 mJy at 350 μm and 500 μm, where the simulation is a factor of 2 below the measurements. In contrast, the number counts extracted from the simulated maps (red solid line) agree well with the observations. The resolution has thus a strong impact on the bright Herschel/SPIRE number counts and models should thus be compared with observations only after having simulated these resolution effects. Consequently, models adjusted directly on the observed number counts potentially overestimate the number of bright dusty star forming galaxies.

The SCUBA2 camera observed deep fields at 450 μm with a 8 arcsec angular resolution (Chen et al. 2013; Casey et al. 2013; Geach et al. 2013; Zavala et al. 2017). In Fig. 4, these data points are shown using yellow, orange, and brown colors. We did not attempt to correct for the slightly different wavelength, since the 450 μm/500 μm color varies strongly with redshift. The latest data points of Zavala et al. (2017) agree very well with the intrinsic number counts in our simulation. This is not surprising, because the much better resolution of SCUBA2 compared with SPIRE limits the effect of resolution on the number counts. Our simulation also well agree with Chen et al. (2013) and Geach et al. (2013). Casey et al. (2013) measurements have a 3σ excess between 10 and 20 mJy and disagree with both the previously quoted measurements and our simulation.

3.3. Ground-based (sub)millimeter number counts

Contrary to number counts at λ ≤ 500μm, number counts at 850 and 1.2 mm were measured with both interferometers and single-dish telescopes. Karim et al. (2013; see also Simpson et al. 2015) showed that number counts derived using low- and high-angular-resolution data are inconsistent. These wavelengths are thus essential to test the ability of our simulation to consistently describe these resolution effects. The comparison between our simulation and the observed number counts at 850 μm and 1.2 mm is presented in Fig. 5. In order to homogenize these data taken at heterogeneous wavelengths, we applied a multiplicative factor of 0.8 to the 1.1 mm data to convert them at 1.2 mm and a factor of 1.07 to 870 μm data to convert them at 850 μm. These factors are derived using our main sequence SED template at z = 2 and are only weakly redshift dependent.

At 850 μm, our model agrees well with sub-mJy number counts extracted by Oteo et al. (2016) using ALMA calibration observations. Above 1 mJy, we have access to two types of constraints: single-dish measurements (in orange, Coppin et al. 2006; Geach et al. 2017) and interferometric follow-up of these bright single-dish sources (blue and purple, Karim et al. 2013; Simpson et al. 2015). As explained in Karim et al. (2013), the number counts derived from the interferometric follow-up of bright sources are lower than the number counts extracted directly from single-dish data, because the flux density of some single-dish sources is coming from several galaxies. Our intrinsic number counts agree perfectly with the interferometric data. The number counts extracted from the simulated map agree well with Geach et al. (2017), but are slightly lower than Coppin et al. (2006).

At 1.2 μm, our intrinsic number counts are in good agreement with the deep blank ALMA fields (Hatsukade et al. 2013; Fujimoto et al. 2016; Aravena et al. 2016; Oteo et al. 2016). The number counts extracted from the simulated single-dish maps agree perfectly with Scott et al. (2012) and are 1σ higher than Lindner et al. (2011). Scott et al. (2012) used a mix of ASTE and JCMT data. The angular resolution is thus similar to that used in our simulated map. Lindner et al. (2011) data were taken with the IRAM 30-m telescope and have thus an angular resolution two times higher, which explains why these measurements are significantly below the number counts extracted from the simulated map. In contrast, they agree well with the intrinsic number counts of the simulated catalog.

We thus managed to reproduce simultaneously the interferometric and single-dish number counts at 850 μm and 1.2 mm together with those from Herschel. This reconciles the observations at low- and high-angular-resolution and highlights the importance of taking into account both the clustering and resolution effects in the modeling of the evolution of dusty galaxies.

thumbnail Fig. 6

Average fraction of the flux density emitted by the brightest galaxy in the beam as function of the flux density measured in our simulated map with a limited angular resolution (red diamonds). We present our results for Herschel/SPIRE at 250 μm (left) and 500 μm (middle), and for NIKA2 at 1.2 mm (right). The gold squares show the average contribution of the galaxies physically related (| Δz | < 0.01) to the brightest galaxy in the beam. The orange open circles are the sum of the flux density fraction from the brightest galaxy and the other galaxies at the same redshift. The black downward-facing triangles is the flux density fraction from the brightest galaxy measured by Scudder et al. (2016). They used a different definition from ours and divided the flux density of the brightest galaxy by the total of the flux density of all the galaxies in a 1 FWHM radius. The blue upward-facing triangles are the results from our simulation assuming their definition (see Sect. 3.4).

Open with DEXTER

3.4. Multiplicity of single-dish sources

As we shown in Sects. 3.2 and 3.3, number counts at λ ≥ 350μm derived from single-dish observations are severely affected by the limited angular resolution of the instruments. We thus expect that the flux density of bright single-dish sources is emitted by several galaxies. This phenomenon has been well studied at 850 μm from both an observational and theoretical point of view (Karim et al. 2013; Hayward et al. 2013a; Hodge et al. 2013; Cowley et al. 2015). In contrast, it is much less explored for Herschel sources, because of the difficulty in observing with interferometers from the ground below 850 μm. In this paper, we present the results of our simulation of the Herschel sources and predictions for the new NIKA2 camera at IRAM (Monfardini et al. 2011).

For each single-dish source extracted from the simulated map with the method described in Sect. 3.1, we searched in our simulated catalog for the brightest galaxy in the beam. We used a search radius of 0.5 FWHM, since the brightest galaxy is usually close to the center of the single-dish source (<0.15 FWHM on average for Herschel and NIKA2 data) and we want to avoid selecting a galaxy contributing to another close single-dish source. We then computed the ratio between the flux density of this brightest galaxy in our simulated catalog and the flux density of the single-dish source measured in our simulated map. In Fig. 6, we show the average ratio as a function of the measured single-dish flux density.

We also estimated the fraction of the flux density emitted by other galaxies at a similar redshift as the brightest galaxy. We chose to define a redshift as similar if | Δz | < 0.01. This value was determined using the histogram of the difference between the redshift of the brightest source and the other sources in the beam. This histogram has a very sharp peak around Δz = 0 with a FWHM of 0.0072, 0.0064, 0.0047 for SPIRE 250 μm, SPIRE 500 μm, and NIKA2 1.2 mm, respectively. Our | Δz | < 0.01 criterion thus corresponds to at least 3σ. We then computed the contribution of these physically related sources to the single-dish flux density. The easiest way to proceed would be to sum the flux density of all the galaxies at the same redshift and closer than a given distance. Unfortunately, this definition is problematic, since the result will depend significantly on the chosen search radius. We thus chose the following alternative method. For every galaxy at the same redshift, we computed their contribution at the center of the single-dish source by multiplying their flux density in the simulated catalog by , where d is the distance between the galaxy and the center of the single-dish source and σbeam is the size of the Gaussian beam. We finally divided the sum of the contribution of all these sources at the same redshift as the brightest source by the measured single-dish flux density measured in the simulated map. The results are presented in Fig. 6.

At 250 μm, 80  to 90% of the flux density is emitted by the brightest galaxy. Scudder et al. (2016) found ~50% based on a Bayesian source-extraction method using shorter wavelength priors (black downward-facing triangles). These results could seem to contradict our analysis. However, they used a very different definition of the flux density fraction. They divided the flux density of brightest galaxy by the sum of the flux density of all the galaxies in a 1 FWHM radius. In our simulated catalog, this sum is larger than the flux density measured in the simulated map. Indeed, the numerous faint sources are responsible for a background (Dole et al. 2003), which is removed by photometric tools, and thus do not contribute to the flux densities measured in our simulated maps. In addition, the galaxies at 1 FWHM from the Herschel sources can contribute to another close single-dish source. Using the same method as Scudder et al. (2016), we find a similar value of 50%. However, the trend with the flux density is different. We find a rising trend, while they have a decreasing one. Their observational method is based on several important assumptions and only high-resolution far-infrared observations will allow to identify which are the most reliable. Finally, we estimated the average contribution of the other sources at the same redshift and found 5%. The sum of the flux density of the brightest galaxy and other galaxies at the same redshift remains smaller than unity. There is thus a significant contribution of galaxies at different redshifts than the brightest galaxy to SPIRE 250 μm sources.

At 500 μm, resolution effects are much stronger and on average only 58% of the flux density is coming from the brightest galaxy. At 70 mJy, this fraction is compatible with unity. The >60 mJy SPIRE sources are essentially local star-forming objects and lensed galaxies, which are sufficiently bright to be detected by themselves. The clustering of nearby objects is weaker than at high redshift. The contrast between a magnified source and its unlensed environment is also high. This explains why these galaxies have a smaller contamination from other galaxies when their flux density is measured with a single dish. The contribution to the measured flux density from the physically related neighbors is ~10% between 20 mJy and 40 mJy and decreases to 5% for these bright sources, in agreement with our understanding.

At 1.2 mm, we produced predictions for the NIKA2 camera (e.g., Calvo et al. 2016). The data will be less affected by resolution effects. The contribution of the brightest galaxy to the NIKA2 sources is ~95% at all flux densities, except the faintest ones that are close to the confusion limit. The resolution effects will be smaller with NIKA2, essentially because of the smaller beam (~12 arcsec). The contribution from galaxies at the same redshift is ~5% and this fraction does not evolve significantly with the flux density. At all flux densities, the sum of the brightest galaxy and other galaxies at the same redshift is responsible for at least 97% of the single-dish flux density measured in our simulated map. The contamination by low-redshift galaxies will thus be much smaller than with Herschel, because they are observed far from their peak of emission.

3.5. The importance of a SFR limit

thumbnail Fig. 7

Effect of the SFR cut on the number counts at 850 μm. The data points are similar to those of Fig. 5. The black solid line is our standard model, that is, with a SFR limit of 1000 M/yr, and the brown dashed line is the model without SFR limit.

Open with DEXTER

In Sect. 2.5, we introduced a SFR limit at 1000 M/yr. The impact on number counts below 500 μm is moderate: the model without SFR limit slightly overproduces the number of sources above 100 mJy. On the contrary, the impact is much stronger at 850 μm, as shown in Fig. 7. This is not surprising because longer wavelengths are dominated by higher redshifts, where the sSFR is on average higher and more sources are thus affected by this limit. The models with and without SFR limit start to diverge at 4 mJy. Above 10 mJy, the model without SFR limit is 5σ above the counts of Geach et al. (2017), which should already be taken as upper limits since they are extracted from single-dish observations. The version of the model without SFR limit is thus clearly ruled out, proving a posteriori the necessity to introduce this threshold.

The SFR limit used in our simulation is an effective way to obtain number counts at the bright end in agreement with observations in the submillimeter. Other modifications could have produced similar number counts. Without SFR limit, the S850> 10 mJy galaxies in our simulation are massive (M ⟩ = 8.6 × 1010M) and at relatively high redshift (z ⟩ = 2.9). A smaller number density of massive star-forming galaxies could thus have a similar impact on the number counts. Because of the steepness of the SMF at the high-mass end, the uncertainties on the stellar mass measurements could produce an artificial excess of massive objects (an effect similar to the Eddington bias). However, it was taken into account by Davidzon et al. (2017) in their fit of the SMF. Some massive passive galaxies could also have been wrongly classified as star forming. However, the main sequence measured by stacking of star forming galaxies would also be lower. Finally, the boost of star formation in starbursts (5.3 in our simulation following Schreiber et al. 2015) could be lower in massive galaxies at z> 2. A lower boost or a SFR limit causing a truncated sSFR distribution are very hard to disentangle with the current data. We thus chose the SFR-limit solution for its simplicity.

The infrared luminosity functions at z> 2 measured with Herschel contain objects above 1013L (SFR> 1000 M) even if their density drops quickly above this luminosity (e.g., Gruppioni et al. 2015; Mancuso et al. 2016). We showed in Sect. 3.4 that the SPIRE fluxes could be overestimated because of resolution effects. This could propagate to the luminosity function as discussed in Sect. 4.2. Most of the interferometric follow-up observations of these abundantly star-forming objects were performed at λ> 850μm. Future ALMA band-9 observations (450 μm) would thus be valuable for confirming the measurements of their obscured SFR.

3.6. Impact of clustering on stacking analysis

Since confusion limits the detection of faint individual galaxies with single-dish instruments, a large fraction of the far-infrared and millimeter observables were measured by stacking analysis. Stacking analysis can also be biased by clustering effects. Since galaxies are clustered, there is a higher probability of finding a source in the beam of a stacked source than at a random position (e.g., Marsden et al. 2009; Béthermin et al. 2010b). Consequently the average flux density of a galaxy population measured by stacking tends to be biased toward higher values. This bias was extensively discussed in the literature and various methods were proposed to correct for this effect (e.g., Marsden et al. 2009; Kurczynski & Gawiser 2010; Béthermin et al. 2010b; Viero et al. 2013a; Heinis et al. 2013; Welikala et al. 2016).

Our simulation is built using two observational studies based on stacking: the evolution of the main sequence measured by Schreiber et al. (2015) and the evolution of the SEDs presented in Béthermin et al. (2015a). These results were corrected for the clustering bias using empirical approaches. Since they are key elements in the calibration of our simulation, we checked that these empirical corrections are consistent with the biases we measure in our simulation. We discuss only Herschel/SPIRE data, since shorter wavelengths have a negligible bias (<10%, Béthermin et al. 2015b; Schreiber et al. 2015). As detailed in Appendix C, our values of the excess of flux density caused by clustered neighbors agree well with the estimate of Schreiber et al. (2015): 13 ± 1% in our simulation versus % in theirs at 250 μm, 21 ± 1% versus % at 350 μm, and 34 ± 1% versus %. In Béthermin et al. (2015a), we used a redshift-dependent correction estimated using two different techniques, which also agrees with our simulation as explained in the Appendix C. The observables derived from a stacking analysis corrected from clustering were thus paradoxically more reliable than the statistical properties derived from catalogs of individually-detected sources.

4. Redshift-dependent observables and consequences on the star formation history

In this section, we compare the results of our simulation with redshift-dependent observables (redshift distributions, number counts per redshift slice) and discuss the impact of these results on the determination of the obscured star formation history.

4.1. Comparison with observed redshift distributions

thumbnail Fig. 8

Comparison between the measured redshift distributions and the predictions of our simulation. The orange histograms are the intrinsic redshift distributions from our simulation. The data points are extracted from Berta et al. (2011) at 100 and 160 μm, Béthermin et al. (2012c) at 250 μm, Geach et al. (2013) and Casey et al. (2013) at 450 μm, Wardlow et al. (2011), Simpson et al. (2014), and Chapman et al. (2005) at 850 μm, and Michałowski et al. (2012), Smolčić et al. (2012), and Yun et al. (2012) at 1.1 mm. Figure adapted from Béthermin et al. (2015b).

Open with DEXTER

thumbnail Fig. 9

Illustration of the impact of the 24 μm-prior extraction on the redshift distribution at 250, 350, and 500 μm. A flux density cut of 20 mJy was used to select the SPIRE sources. The red line is the intrinsic redshift distribution, while the blue line is obtained after extracting the sources in our simulated Herschel/SPIRE maps using 24 μm positions as a prior. The histograms are normalized in order to have dN/ dz dz = 1. We compare these model predictions with the observational redshift distribution of Béthermin et al. (2012c, black diamonds). The data points were extracted in real Herschel data using position, 24-μm flux density, and redshift as a prior. This more complex method was chosen to avoid the potential biases associated with a source extraction using only 24-μm positions as a prior (see discussion in Sect. 4.1).

Open with DEXTER

In our simulation, we implemented significant modifications compared to the Béthermin et al. (2012a) version of the model as the updated evolutions of the SEDs and of the SFR-M relation. We thus checked if this updated model reproduces correctly the observed redshift distributions in Fig. 8 (see Béthermin et al. 2015b for a detailed discussion about the modeling of the redshift distributions). There is an overall good agreement between the intrinsic redshift distributions in our simulation and the measured ones from 100 μm to 1.1 mm.

However, the measurement of the redshift distributions is a complicated task, which requires identification of the galaxy responsible for the main fraction of the far-infrared or submillimeter flux and measurement of its redshift. Various methods can be used. The procedures based on high-resolution follow-up are difficult to reproduce with our simulation. In contrast, our simulation is perfectly suited to testing the prior-based source extraction, which was used to derive Herschel redshift distributions (e.g., Berta et al. 2011; Béthermin et al. 2012c).

In Fig. 9, we compared the intrinsic redshift distribution in our simulated catalog with the redshift distribution of the sources extracted in our simulated map using 24 μm positions as a prior as described in Sect. 3.1. At 250 μm, the two distributions are very similar, showing this extraction technique does not bias the results. We obtained similar results at shorter wavelength with Herschel/PACS. At 350 μm, we found that the redshift distribution derived from the source extracted in the simulated map is slightly biased toward lower redshifts compared to the intrinsic distribution. At 500 μm, this bias becomes stronger. As discussed in Sect. 3.4, the flux density of 500 μm sources is emitted by several galaxies. In addition, the 24 μm/500 μm color varies more with redshift than colors between 24 μm and shorter wavelengths. The brightest 24 μm galaxy in a 500 μm beam is thus not systematically the main contributor to the 500 μm flux density. In conclusion, the Herschel redshift distribution extracted using 24 μm positional priors are thus accurate only below 250 μm. At longer wavelength, other methods must be used.

In Béthermin et al. (2012c), we used a prior-based extraction based on both the 24 μm flux density and the redshift. Instead of directly selecting the brightest 24 μm source in a 0.5 FWHM radius as an input for FASTPHOT, we predicted the 500 μm flux from the 24 μm flux density and the redshift, and kept in the prior list the galaxy with the highest predicted flux density at 500 μm in a 0.5 FWHM radius. For the prediction, we used the average colors measured by stacking. We thus kept more high-redshift sources in the prior list. This approach agrees with the intrinsic redshift distribution in our simulation (Fig. 9), but not with the extracted one. This highlights that more advanced prior-based source extraction techniques could be sufficient to derive accurate redshift distributions from confusion-limited maps. Our simulation will be particularly useful for validating future studies of the redshift distributions.

thumbnail Fig. 10

Number counts per redshift slice. Purple, blue, orange, and red are used for the z< 0.5, 0.5 <z< 1, 1 <z< 2, and z> 2 slices, respectively. The data points are from Berta et al. (2011) at 70, 100, and 160 μm and Béthermin et al. (2012c) at 250, 350, and 500 μm. The dotted lines are the intrinsic distributions in the simulation and the solid lines are the counts extracted using 24 μm priors (see Sect 3). Figure adapted from Béthermin et al. (2012c).

Open with DEXTER

4.2. Number counts per redshift slice

In Fig. 10, we compare the results of our simulation with the measured number counts per redshift slice measured with PACS (Berta et al. 2011) and SPIRE (Béthermin et al. 2012c). This observable is very close to monochromatic luminosity functions (Gruppioni et al. 2013; Magnelli et al. 2013), but is not affected by the assumptions made on the K-corrections, which are necessary to determine the luminosity functions. The full observational process used to measure the number counts per redshift slice is thus easier to simulate. There is an overall good agreement between 70 and 160 μm at z< 2. At z> 2, our simulation under-predicts the source counts at 70 μm and 100 μm by a factor of 5 and 2.5, respectively. The number counts in our simulated catalog and after simulating the full source-extraction procedure are similar. This is thus not a problem caused by the resolution. The most likely explanation is contamination by active galactic nuclei (AGNs), since these Herschel bands at z> 2 correspond to <23 and <33 μm rest-frame. Indeed, we did not implement the contribution of AGNs to the mid-infrared emission in our simulation, which focuses on the far-infrared and millimeter domain. However, at these wavelengths, 99% of the sources lie at z< 2 and the AGN contribution to the SEDs has thus a negligible impact on the global statistical properties of the galaxies.

At 250 μm, the number counts in our simulation at z< 2 agree well with observations and there is no significant difference between the intrinsic number counts and those extracted using 24 μm priors. At z> 2, the intrinsic counts under-predict the observations at the bright end, but the number counts extracted from the simulated map agree well with the data. At 350 μm and 500 μm, the intrinsic number counts are systematically below the observations at z> 0.5, but the number counts extracted from the simulated maps agree better with the data. However, the number counts extracted from the simulated map tend to be lower than the observations at z> 2 and higher at z< 0.5. We note however, that the source extraction from the simulated map was done using only the 24 μm position as a prior and is thus slightly biased toward low redshift, as shown in Sect. 4.1.

thumbnail Fig. 11

Evolution of the obscured star formation density as a function of redshift. The black line is the result of our simulation. We show for comparison the infrared measurements of Sanders et al. (2003, diamond), Takeuchi et al. (2003, triangle), Magnelli et al. (2011, squares), Magnelli et al. (2013, circles), Gruppioni et al. (2013, stars), and Rowan-Robinson et al. (2016, circles with cross). We also plot the UV estimates corrected for dust attenuation from Wyder et al. (2005, asterisk), Schiminovich et al. (2005, downward-facing triangles), Robotham & Driver (2011, right-facing triangle), Cucciati et al. (2012, left-facing triangle), Dahlen et al. (2007, hourglass), Reddy & Steidel (2009, hat), and Bouwens et al. (2012a,b, bowtie). All the data have been converted to Chabrier (2003) IMF and Planck Collaboration XIII (2016) cosmology.

Open with DEXTER

4.3. Consequences on the obscured star formation history

In the previous sections, we have shown that the flux densities of individually-detected sources (Sect. 4.2) are biased toward higher values because of angular resolution effects, while stacking-derived observables were already corrected from the clustering effects. Since the peak of the far-infrared emission of galaxies is around 100 μm rest-frame, Herschel/SPIRE data are thus essential to derive accurate obscured SFR, but unfortunately they are affected by these resolution effects. They are also limited by the confusion and only sources brighter than ~20 mJy can be extracted reliably from the maps. These bright, individually detected sources have an important role in understanding the evolution of the massive systems, but they contribute only marginally to the global star formation budget.

At 250 μm, the resolution has an impact only at z> 2 (Fig. 10). In our simulated catalog, at z> 2, the S250> 20 mJy galaxies contribute to only 2.5% of the obscured star formation density. At 350 and 500 μm, the galaxies brighter than 20 mJy host only 2.9 and 1.4% of the SFRD, respectively, at z > 2. At those fluxes, the excess of flux density caused by the resolution effects (see Sect. 3.4 and Fig. 6) is 21, 46, and 96% at 250, 350, and 500 μm, respectively. It is hard to propagate this effect to the estimate of the total infrared luminosity density and star formation density (SFRD), since it requires combining several wavelengths. However, even in the worst case scenario of using only 500 μm as a SFR estimator, the excess of SFRD caused by sources brighter than 20 mJy will remain below 10%. This effect remains thus below the systematic uncertainties associated with the extrapolation of the contribution of the faint sources.

We checked if the SFRD in our simulation agrees or not with other estimates from the literature. In Fig 11, we compare the obscured SFRD from our simulation with the latest observations compiled by Madau & Dickinson (2014). Our simulation agrees well with both the IR- and UV-derived measurements up to z ~ 3. This confirms that the impact of resolution effects are minors on the global star formation budget.

At z > 3, our simulation is 2σ higher than the measurements of Gruppioni et al. (2013) derived from Herschel observations. However, they have only three data points at LIR > 1012.5L and have to make strong assumptions about the faint-end slope of the luminosity function. In contrast, our simulation is 0.5σ and 2σ lower than the estimate of Rowan-Robinson et al. (2016) at z = 4 and z = 6, respectively. There is thus significant tension between the various estimates of the obscured SFRD at z > 3. Our simulation agrees with the measured redshift distributions and deep millimeter counts and is thus compatible with the current non-extrapolated data. These differences between studies highlight how uncertain the obscured star formation history at z > 3 remains. Concerning the observation derived from dust-corrected UV, our simulation agrees with Bouwens et al. (2012a,b) measurement at z > 5, but is 50% higher at z ~ 4. This suggests that a fraction of the star formation might have been missed at this redshift by optical surveys. Future wide and deep millimeter surveys with NIKA2 at IRAM and with the large millimeter telescope (LMT) will be essential to confirm or not this result.

5. One- and two-point map statistics

In addition to the statistical properties of the sources, we checked the agreement of our model with map statistics. This is particularly important for SPIRE data, which has a limited resolution.

thumbnail Fig. 12

Pixel histograms of the Herschel/SPIRE maps in COSMOS (black) and comparison with our simulation using the same instrumental noise map. The red histograms are the result of our simulation and the blue ones are the histogram obtained after randomizing the position of the sources to illustrate the impact of clustering.

Open with DEXTER

thumbnail Fig. 13

Power spectrum of cosmic infrared background anisotropies at 250, 350, and 500 μm. The blue solid lines are CIB anisotropies measured in our simulation and 1σ confidence regions are represented by the dashed lines. The black diamonds and the red triangles are the measurements of Viero et al. (2013b) and Planck Collaboration XXX (2014), respectively.

Open with DEXTER

5.1. Pixel histograms: P(D)

The distribution of the surface brightness in the pixels (P(D)) of a map is directly connected to the number counts of the objects in this map (Scheuer 1957; Condon 1974). This method was used to measure the faint source counts with, for example, Bolocam (Maloney et al. 2005), LABOCA (Weiß et al. 2009), BLAST (Patanchon et al. 2009), and Herschel (Glenn et al. 2010). As discussed in Takeuchi & Ishii (2004) and Patanchon et al. (2009), clustering could impact P(D) analysis. However, when these analyses were performed, simulations did not include clustering that we know is critical. We can now investigate this effect using our new simulations.

In Fig. 12, we compare the pixel histograms of the simulated maps and of real Herschel maps. We used the COSMOS maps4 from the HerMES survey (Oliver et al. 2012), which match the size of our simulation. We used the real noise maps released by the HerMES team to generate a similar Gaussian instrument noise in our simulated maps. In order to evaluate the impact of clustering, we produce another simulated map without clustering by randomly reshuffling the positions of the galaxies.

At 250 μm, the clustering has an impact of less than 5% and both clustered (red) and unclustered (blue) simulated maps agree at 5% with the observed histogram. This is a very good agreement considering the 4% calibration uncertainty of SPIRE (Bendo et al. 2013). At 350 and 500 μm, the effect of clustering is much larger because of the larger beam and can reach 15%. The clustered maps agree well with the observed one, but the randomized maps have a large excess at the peak. This shows that clustering has a non-negligible effect on the P(D) analysis and must be taken into account. This could explain why the P(D) analysis of Glenn et al. (2010) agrees with individual source counts even if they are biased high compared to the intrinsic counts (see Fig. 4).

5.2. Anisotropies of the cosmic infrared background: P(k)

Measuring the clustering of individually detected population in confusion-limited data is difficult. The sample sizes remain limited to obtain good statistics (e.g., Béthermin et al. 2014). The contamination of the fluxes by the neighbors tends to introduce artificial correlation between redshift slices and to bias the measurements (Cowley et al. 2016). The power spectrum of the CIB anisotropies, which is not affected by this problem, is currently the best way to constrain how the star formation is distributed in dark-matter halos (e.g., Lagache et al. 2007; Béthermin et al. 2013; Planck Collaboration XXX 2014; Viero et al. 2013b). CIB anisotropies are a powerful observation to test that the model simultaneously reproduces the infrared emission and the spatial distribution of galaxies.

In Fig. 13, we compare the power spectrum measured with Herschel (Viero et al. 2013a; black diamonds) and in our simulation (blue solid lines). In order to reduce the Poisson noise, the brightest sources are usually masked or subtracted from the maps. We chose to use a Sν,cut = 50 mJy flux density cut, which is the deepest cut used by Viero et al. (2013a). We also included the Planck data at 857 GHz (350 μm, red triangles). We shifted these data by the difference of the Poisson noise in our simulations between a flux density cut of 50 mJy and 710 mJy (used by Planck). Finally, as shown by Bertincourt et al. (2016), the Herschel/SPIRE 500 μm absolute flux calibration is 4.7% too high compared to Planck. We thus corrected Viero et al. (2013a) data points accordingly. To accurately measure the power spectrum, we generate a map without convolving it by the PSF and without including the sources above the flux cut. To be fully consistent with the observational process, we produced a map using the SPIRE spectral response to extended emission to measure the power spectrum, but we used the flux densities from the point-source spectral response to select the sources to put in the map (see Lagache et al., in prep., for a detailed discussion). We measured the power spectrum from these simulated maps using the POKER software (Ponthieu et al. 2011; Planck Collaboration XVIII 2011; Planck Collaboration XXX 2014). This software accounts for non-periodic boundary conditions of the map that otherwise bias large-scale measurements. The error bars are estimated via Monte-Carlo simulations of the estimated power spectrum.

At small scale (k> 0.3 arcmin-1), the power-spectrum is dominated by the shot noise from galaxies. These Poisson fluctuations of the number of galaxies in a patch of sky produce a plateau in the power spectrum, which can be derived directly from the number counts (Lagache et al. 2000): (8)where Sν is the flux density and are the differential number counts. At 350 and 500 μm, our simulation agrees at 1σ with the measurements of Viero et al. (2013a). At 250 μm, our simulation is systematically 1.5σ above the measurements. Since the measurements are dominated by systematic effects (e.g., deconvolution of the beam), their error bars are strongly correlated. This offset is thus not statistically significant. Overall, the Poisson level in our simulation and in real data agrees.

Mak et al. (2017) found a discrepancy between the Poisson level measured in Herschel and Planck data and that derived from the measured Herschel number counts using Eq. (8). The Poisson level derived directly from the number counts are higher than the measurements. This problem can be solved naturally considering the discrepancy between the measured and the intrinsic number counts that we identified in Sect. 3.2. Indeed, the Poisson level depends on the number counts. Since the observed number counts are overestimated at 350 and 500 μm because of resolution effects, the Poisson levels derived from them are thus overestimated.

At 250 μm, our simulation agrees at better than 1σ with the data at large scale (k < 0.3 arcmin-1). In this regime, the power spectrum is dominated by the large-scale clustering of galaxies. At 350 and 500 μm, at k < 0.1 arcmin-1, our simulation underestimates the power spectrum by 1.5σ. Future larger simulations will allow us to determine if this deficit is real or just a statistical fluctuation.

6. The nature of Herschel red sources

Dowell et al. (2014) and Asboth et al. (2016) found a large population of red Herschel sources in the HerMES survey (Oliver et al. 2012). They claimed that the number of sources they found is one order of magnitude higher than predicted by the models. If confirmed, these results would suggest that the models strongly underpredict the number of bright z> 4 dusty star forming objects. Ivison et al. (2016) also found a large number of red high-redshift candidates in the H-ATLAS survey using a slightly different selection.

In this section, we verify that our simulation accurately reproduces the statistics of red sources. First, we check the statistics of red sources in our 2 deg2 simulation, which is a small area but includes clustering (Sect. 6.1). We then investigate the statistics of red sources in a large simulated catalog, without clustering (Sect. 6.2). The criteria of Ivison et al. (2016) are hard to reproduce, since they involve some visual inspection. We thus focus our analysis on the results of Asboth et al. (2016), who used the following criteria: S250 < S350 < S500, S500 > 52 mJy, and D = 0.92 M500−0.392 M250 > 34 mJy, where M500 and M250 are the values of the maps at the position of a source at 250 and 500 μm, respectively, after matching all the maps at the resolution of 500 μm data.

6.1. Simulation in map space in 2 deg2

There is no galaxy in our 2 deg2 simulated catalog that follows the Asboth et al. (2016) criteria. However, as we will show, some of these sources could be explained by noise fluctuations and resolution effects.

Asboth et al. (2016) homogenized the beams to the size of the 500 μm one. For simplicity, we directly generated three SPIRE maps using a Gaussian beam with a FWHM of 36.3 arcsec. We then added instrumental noise using the same values as in Asboth et al. (2016). The D map is generated from the 250 μm and 500 μm maps. The noise in our simulated D map is very close to the observations: 8.8 versus 8.5 mJy. This shows that our approximation on the beam is sufficiently good to perform our analysis of red sources.

We extracted the peaks higher than 34 mJy in the D map and measured the photometry on the maps at the native SPIRE resolution using the flux in the central pixel of the source after subtracting the mean of the map. The number of detected red sources varies depending on the realization of the noise. We thus used 1000 realizations to estimate the mean number of detected sources. This estimate does not take into account the cosmic variance. We found 1.7 in our 2 deg2 field, which corresponds to 229 sources in 274 deg2. This agrees at 1σ with the 477 detections reported by Asboth et al. (2016). Even if the statistics are very limited, our results indicate that there might be no real tension between models and observations of red sources.

thumbnail Fig. 14

Differential number counts of red Herschel sources. The black diamonds are the measurements from Asboth et al. (2016). The red point is the result of an end-to-end analysis of our 2 deg2 simulated maps including clustering and instrumental noise. The turquoise, dark blue, and gold lines are the number counts in a simulated catalog of 274  deg2 without noise, with instrumental noise only, and with both instrumental and unclustered confusion noise, respectively. Blue and purple arrows illustrate the impact of instrumental and confusion noise, respectively.

Open with DEXTER

6.2. Red sources in a 274 deg2 catalog

Our simulation covers only 2 deg2 and thus contains small statistics compared to Asboth et al. (2016) who used 274 deg2. We thus generated another simulated catalog based on the same prescriptions and covering the same area of 274 deg2 as Asboth et al. (2016). Since we do not have a sufficiently large dark-matter simulation, we have to ignore the clustering and draw directly the sources from the stellar mass function. With this simplified method, we could potentially underestimate the number of red sources, since we neglect the boosting of the flux of massive high-redshift dusty galaxies by their neighbors. This effect can be potentially important because of the strong clustering of the most star-forming galaxies at z > 2 (e.g., Farrah et al. 2006; Béthermin et al. 2014). We generated only sources with M > 1010M and z > 1 to save memory, because no source with a lower mass can be sufficiently bright and no source below z = 1 can be sufficiently red to pass the “red source” selection. We do not produce maps so we compute D from the flux densities (0.92 S500−0.392 S250).

thumbnail Fig. 15

Left panel: redshift distribution of the red sources. The blue and black histograms are the distributions of the sources matching the red criteria before and after including noise (both instrument and confusion noises), respectively. Middle panel: distribution of the red excess (D = 0.92 M500−0.392 M250). Right panel: distribution of the 500 μm flux density. The blue and black histograms show the intrinsic (without including noise) 500 μm flux densities (middle panel) and D values (right panel) of the intrinsically red (red criteria applied before adding noise) and observed red (red criteria applied after adding noise) sources, respectively. The red histogram shows the distributions of the observed, that is, after including noise, 500 μm flux densities and D values of observed red sources. The difference between the black and red histogram illustrates that sources observed above the S500 and D thresholds of Asboth et al. (2016, vertical black dotted line) usually have lower intrinsic values.

Open with DEXTER

We have only 18 objects following the criteria of Asboth et al. (2016) in our simulated catalog compared to the 477 objects in Herschel data. All these 18 sources are strongly lensed (μ> 2). Our simulation contains 439 256 non-lensed sources with S250<S350<S500, but none of them are sufficiently bright to satisfy S500> 52 mJy. Number counts of red sources at 500 μm are shown in Fig. 14. While the counts from our simulation above 100 mJy are close to the data, they are well below for fainter flux densities. The problem is thus not coming from the lensing, but from a lack of intrinsically bright sources. If we remove the SFR limit in our simulation, we find 205 non-lensed sources matching the Asboth et al. criterion (676 if we add noise to the simulated catalog, see next paragraph). However, this SFR cut is necessary to be consistent with the 850 μm number counts (see Sect. 3.5). Another explanation must thus be found to understand this discrepancy.

The results described in the previous paragraph ignore the effect of noise on the number counts of red sources. We thus simulate the effect of both confusion and instrumental noise by adding a random Gaussian noise to the flux densities of the simulated sources. We used the noise values provided in Table 1 of Asboth et al. (2016). This method produces a higher noise on D than in the real D map, since the confusion noise from 250 and 500 μm tend to partially cancel each other out in the real maps. Because of the very steep color distribution and number counts, the noise strongly increases the number of red sources (see Fig. 14). We found 168 sources matching the Asboth et al. criteria (74 with only instrumental noise), of which only 29% are strongly lensed (60% with only instrumental noise). The noise has thus an important role in producing red sources without strong lensing. Concerning the weakly-lensed sources, the weakly-lensed red sources are on average 6% more magnified than the average magnification at z> 2. Red-source selections are thus biased toward higher magnifications. Weak lensing acts as an additional noise on the flux density and sources on a positive fluctuation of the magnification tend to pass the 500 μm flux density threshold more often. This is similar to the Eddington bias. Without weak lensing, the number of red sources decreases to 132.

In our simulation, the noise thus strongly increases the number density of detected red sources. These results could seem to be in contradiction with Asboth et al. (2016), who found that the number of injected and recovered red sources was similar with the number of recovered ones using an end-to-end simulation. However, we identified one potentially incorrect assumption in their simulation. Their simulation used the Béthermin et al. (2012a) model with the intrinsically red source removed and a power-law distribution of red sources with a fixed color based on the median observed one. The flux distribution of the red sources in their simulation is based on the number of detected red sources directly extracted in the real map, which is one order of magnitude higher than what is intrinsically in the Béthermin et al. (2012a) model. The relative contribution to the extracted number counts of intrinsically non-red sources, which matches the red criteria because of the noise and resolution effects, might thus be significantly underestimated because of the very high number of intrinsically red sources in their simulation. This highlights the difficulty to correct the biases in statistical measurements of red sources using only inputs from observations.

In Fig. 15, we illustrate the impact of the noise on the selection of red sources. If we apply the Asboth et al. (2016) criteria to the intrinsic fluxes of the sources, red sources are selected at z> 3. If we select red sources with the same criteria after adding noise, the number of detections at z< 4 increases dramatically, while the number of z> 4 detections remains almost constant (see left panel). The noise has thus a strong impact on the redshift distribution of red sources. The red sources selected in the noisy catalog usually have a measured D value and a 500 μm flux density just above the limit (vertical black dotted line in middle and right panels of Fig. 15). However, their intrinsic D value and 500 μm flux density are in general below the detection limit (D ⟩ = 23 mJy and S500 ⟩ = 42 mJy). The intrinsic non-red sources selected because of the noise are thus not purely spurious objects. They are mostly strongly star-forming objects with a D value and an intrinsic 500 μm flux density slightly below the cut. The first red sources that have been followed up were selected in deep fields and were usually the reddest sources in the sample (e.g., Riechers et al. 2013; Dowell et al. 2014). These objects were confirmed spectroscopically to be at high redshift. However, their selection was less affected by the noise because of their particularly red observed color and the lowest instrument noise in such fields.

6.3. The challenge of using red sources to constrain models

Overall, the combination of the noise and the weak lensing can dramatically boost the number density of detected red sources compared to their intrinsic number density. Our 2 deg2 end-to-end simulation including clustering is compatible with the observations, but the statistics are limited. We also built a larger simulation based on a catalog of galaxies with random positions and a random Gaussian noise. However, the number of red sources remains lower by a factor of 2.8 in this simplified simulation. This could be explained by several effects. As shown in the pixel histogram of Herschel maps in Fig. 12, the combination of instrument and confusion noise has an asymmetrical distribution with a large positive tail, which could be responsible for more bright 500 μm outliers than in the Gaussian case. In addition, the faint foreground sources are clustered and a local underdensity of faint blue foreground galaxies in a beam could create an artificial red color. Finally, the model is sensitive to the value of the SFR cut and better observational constraints should allow us in the future to determine its value or favor one of the alternative scenarios described in Sect. 3.5.

Comparing observations of red sources and models remains complicated, because it would require using the exact same algorithm on simulated maps with clustering of ~100 deg2. For wide shallow fields as the one used by Asboth et al. (2016), the noise also plays a crucial role. Applying a new method of deblending to the deeper 55 deg2 HeVICS field, Donevski et al. (2017) found that number counts of red sources are an order of magnitude lower than Asboth et al. (2016). They simulated the effect of instrument noise and confusion on our simulated catalogs and found an excellent agreement with their new observations. They also found that the noise has only a mild impact on the number counts of red sources in this deeper field. This highlights the need to perform end-to-end simulations with the same extraction algorithm and realistic noise properties to compare observations and model. However, these simulations remain extremely difficult to perform on large areas because they require both a large-volume dark-matter simulation and a sufficiently-high-mass resolution to have the faint galaxies responsible for confusion.

In addition to these difficulties, the lensing was included in a non-consistent way in our simulation, since it was drawn randomly in a distribution, which does not vary with the position of the source (see Sect. 2.8). Non-trivial biases can occur in color selections for lensed sources. For instance, the sources clustered with the lens can change the color of the source, since they are at lower redshift and thus bluer than the background source (Welikala et al. 2016). In addition, we used only a simplified Gaussian weak lensing. The number of sources with a magnification between 1.5 and 2 is thus underestimated compared to the numerical simulation of Hilbert et al. (2007). For instance, the extreme starbursts reported by Riechers et al. (2013) were later proved to be lensed (Cooray et al. 2014) by a small factor.

Interpreting the statistics of red sources with models is thus a very challenging task, because of the complexity of the various artifacts affecting the selection of these objects. Direct millimeter selections as SPT (Vieira et al. 2013) provides more straightforward constraints for models. However, we should mention that, despite the difficulty in interpreting their number counts, red-source selections in Herschel fields are a powerful tool to build large samples to study the physics of high-redshift, dusty star forming galaxies.

7. Conclusion

We presented a new simulation of the far-infrared and (sub)millimeter sky called SIDES. This simulation is based on an updated version of the Béthermin et al. (2012a) phenomenological galaxy evolution model using the latest observational constraints on the stellar mass function, the main sequence of star forming galaxies, and the evolution of the SEDs. To obtain realistic clustering, we used an abundance matching procedure to populate the dark-matter halos of a light cone constructed from the Bolshoi-Planck simulation (Rodríguez-Puebla et al. 2016) with the galaxies produced by our model. The intrinsic galaxy number counts in this new simulation are significantly lower than the measurements from single-dish instruments, while they agree with interferometric data. To understand this tension between our simulation and the observations, we simulated the full source extraction process and showed that the number counts extracted from our simulated maps agree with the observed ones. When we take into account the observational effects, our simulation is able to simultaneously reproduce the single-dish and interferometric number counts from the far-infrared to the millimeter domain together with redshift distributions, CIB anisotropies, and the pixel histograms of SPIRE maps.

Our simulation also allowed us to evaluate the impact of clustering and angular resolution on some statistical properties derived from far-infrared and (sub)millimeter surveys. We identified the following effects:

  • The flux density of Herschel sources is affected by resolutioneffects. The brightest galaxy in the Herschel beam is responsiblefor ~85% of the flux density at 250 μm, but only ~60% at 500 μm. Other galaxies contributing to the Herschel flux density are both galaxies at the same redshift as the brightest one and randomly aligned galaxies. Our simulation predicts that the fraction of the flux density coming from the brightest galaxy will rise to 95% in future millimeter single-dish surveys performed by 30 m-class telescopes (e.g., NIKA2 at IRAM).

  • Measurements using stacking are also biased by the clustering. However, this bias is already known. The corrections made by the observational studies are compatible with the corrections derived from our simulation. Paradoxically, the stacking studies, which took into account the clustering, are more accurate than the observations of individually detected sources.

  • The redshift distributions of Herschel sources extracted using 24 μm positional priors tend to be biased toward lower redshifts at 350 and 500 μm, but are reliable at shorter wavelengths.

  • Even if the flux density of the brightest Herschel sources tends to be overestimated, the impact of these sources on the global star formation history is small. Our simulation is compatible with the UV- and IR-derived measurements of the SFRD up to z< 3. At z ~ 3.5, our simulation predicts a higher SFRD than the measurements of Gruppioni et al. (2013, IR) and Bouwens et al. (2012b,a, UV). At higher redshift, our results are compatible with the constraints from the UV.

  • The clustering has a significant impact on the pixel histogram of Herschel maps used to perform P(D) analysis. This explains why the number counts derived by Glenn et al. (2010) using this statistical technique are compatible with the measurements derived using standard source-extraction methods, but not with the intrinsic number counts.

  • The resolution effects allow us to solve the tension between the measured level of the Poisson fluctuation of the CIB power spectrum and what is expected from the observed number counts. Indeed, the number counts are biased high and the shot-noise derived from the number counts is thus overestimated.

  • Recently, Asboth et al. (2016) identified a population of red Herschel sources. Their number density is one order of magnitude higher than in models, including our new simulation. However, after taking into account the noise and the resolution effects, our 2 deg2 simulation produces the correct number of objects.

These results highlight the difficulty to interpret the long-wavelength single-dish observations. Correcting the observations for all the observational effects is a complex task. The corrections usually assume an underlying model. In this paper, we started from our model and reproduced the full observational process. This approach is more direct and allowed us to test the validity of our model without having to rely on complex corrections of the data. This approach is probably the best way to deal with the complexity and the precision of the modern data sets.

Our simulation (SIDES), released publicly5, has many potential applications:

  • It can be used to prepare future surveys and in particular to predictthe number of detected galaxies and their properties (redshift,SFR, stellar mass). The realistic clustering included in oursimulation can also be used to accurately estimate the confusionlimit.

  • It is a powerful tool to test source-extraction techniques and characterize their biases. We can also use it to test and optimize methods to identify the counterparts at shorter wavelengths of single-dish sources. Finally, it can be used to validate stacking softwares and determine the most efficient and less biased ones.

  • The various biases affecting the extraction of single-dish sources can have strong impacts on clustering measurements and bias the estimates of the host halo mass (e.g., Cowley et al. 2016). Since it accurately reproduces a large set of observables, our simulation is well suited to characterize and correct for these effects.

  • Finally, in a future paper, we will include the emission of far-infrared and (sub)millimeter lines in our simulation and perform predictions for spectroscopic surveys and for (sub)millimeter intensity mapping experiments.


1

Are publicly available at http://cesam.lam.fr/sides

2

The catalogs are available at http://hipacc.ucsc.edu/Bolshoi/MergerTrees.html

3

The code is publicly available at https://bitbucket.org/yymao/abundancematching

Acknowledgments

We thank the anonymous referee for his/her very useful and constructive comments. We thank Y.-Y. Mao for providing the abundance matching code and for assistance. We thank Eric Jullo, Stéphane Arnouts, Olivier Ilbert, and Corentin Schreiber for their explanations. The Bolshoi-Planck simulation was performed by Anatoly Klypin within the Bolshoi project of the University of California High-Performance AstroComputing Center (UC-HiPACC) and was run on the Pleiades supercomputer at the NASA Ames Research Center. We acknowledge financial support from “Programme National de Cosmologie and Galaxies” (PNCG) funded by CNRS/INSU-IN2P3-INP, CEA and CNES, France. This work has been partially funded by the ANR under the contract ANR-15-CE31-0017. This work has been carried out thanks to the support of the OCEVU Labex (ANR-11-LABX-0060) and the A*MIDEX project (ANR-11-IDEX-0001-02) funded by the “Investissements d’Avenir” French government program managed by the ANR. H.W. acknowledges the support by the US National Science Foundation (NSF) grant AST1313037. Part of the research described in this paper was carried out at the Jet Propulsion Laboratory, California Institute of Technology, under a contract with the National Aeronautics and Space Administration.

References

Appendix A: Homogenization of cosmology

thumbnail Fig. A.1

Left: differential number counts at 100 μm. We used the same data points as in Fig. 4. The red dashed line is before the correction at low redshift of the position of the main sequence (see Appendix B) and the black solid line is the final version of our model. Since this effect is only significant at bright flux density, where the number of sources is small, we used an unclustered 10 deg2 to reduce the statistical uncertainties in our model predictions. Right: star formation density as a function of redshift. The data are similar to Fig. 11, but we added the prediction of our model before the low-redshift correction of the main sequence (red dashed line).

Open with DEXTER

Our simulation is based on observational quantities (e.g., stellar masses, SFR). They were derived assuming a cosmology that was different from the one used in the dark-matter simulation. To be consistent, we convert the observational constraints to the Planck cosmology. At fixed SED, the SFR and stellar masses are proportional to the intrinsic luminosity of an object. In the Planck cosmology, the luminosity distance is ~3% (small redshift dependance) larger than in the 773 cosmology used in most of the observational papers (h = 0.7, ΩΛ = 0.7, ΩM = 0.3). The intrinsic luminosity of the objects, and consequently the stellar masses and SFRs, are thus slightly higher than in the observational papers in 773 cosmology. We thus applied a (DL,Planck/DL,773)2 correction to the observed stellar masses and SFR. Similarly, the volume corresponding to a redshift slice and estimated with 773 cosmology is smaller than with Planck cosmology. The number density of observations estimated in 773 cosmology is thus overestimated. We thus apply (dVcomoving, 773/ dz)/(dVcomoving, Planck/ dz) corrections to the characteristic densities Φ. The correction is computed at redshift corresponding to the center of the redshift bins used to derive the mass and luminosity functions. We checked that this correction does not vary by more than 1% inside a bin.

Appendix B: Correction of the main sequence at low-z

When we use the Schreiber et al. (2015) analytic description of the evolution of the main sequence, the simulation exhibits an excess of SFRD at z< 0.5 and overpredicts the bright end (S100> 30 mJy) of the number counts at 100 μm by ~30% (see Fig. A.1, red dashed line). We found a similar excess at the very bright end at other wavelengths. The simulation of Schreiber et al. (2017), which uses the same description of the main sequence but different SEDs, has a similar excess at S100> 30 mJy (see their Fig. 10). This is thus not a problem of SEDs. These bright sources have a median redshift of 0.22 and a median stellar mass of 2 × 1010M. At this stellar mass and redshift, the SFRMS provided by Schreiber et al. (2015) is 0.1 dex higher than Sargent et al. (2014) estimate based on a large compilation of data. We thus offset the sSFRMS by  dex to correct for this offset. After implementing this correction, the simulation reproduces the SFRD and the bright end of the number counts (Fig. A.1, black solid line).

Appendix C: The impact of clustering on stacking results as a function of redshift

In Sect. 3.6, we discussed the impact of clustering on the measurements by stacking. In this Appendix, we explain the technical details of our comparison between the corrections used by Schreiber et al. (2015) and Béthermin et al. (2015a) and the ones derived from our simulation.

Schreiber et al. (2015) used the real position of their stacked galaxies. They used their stellar mass and redshift to predict their mean expected flux densities at SPIRE wavelengths. They then built a map based on these predicted flux densities. They finally compared the mean flux density in their simulated catalog with the measurements by stacking in their simulated map. They used several methods to perform the photometry (small apertures, PSF-fitting photometry, signal in the central pixel). We chose to use the signal in the central pixel, since it is the easiest to simulate. After we had subtracted the mean value of the map to remove the background, we computed the mean flux density in all the SPIRE pixels hosting a galaxy in the input stacked catalog. The relative excess of flux density caused by clustered neighbors is computed using: (C.1)where Sstack is the mean flux density measured by stacking in the simulated map and Scat is the mean flux of the stacked sources in the simulated catalog. We found 13 ± 1%, 21 ± 1%, 34 ± 1% at 250, 350, and 500 μm, respectively, which agrees well with the values of (Schreiber et al. 2015) of %, %, and %, respectively.

We also compared our results with Béthermin et al. (2015b), who used a redshift-dependent correction. To allow an easier comparison, we used the same stellar mass selection (3 × 1010M) and redshift bins and derived the stacking excess using the method described in the previous paragraph. Our results are presented in Fig. C.1 (black squares). The relative excess caused by clustered neighbors is compatible with zero at z = 0, rises up to a maximum at z ~ 1, and slightly decreases with increasing redshift at z> 2. This decrease of the relative excess at higher z was already discussed in Béthermin et al. (2015b) and was interpreted as the result of the rising of both the rarity and the infrared brightness of star forming galaxies with increasing redshift, causing a higher contrast between the massive galaxies and their environment. The first approach used in (Béthermin et al. 2015b) uses an initial stacking to determine a mass-to-flux-density ratio evolving on the redshift (blue triangles). A simulated map is then produced from the real COSMOS catalog from the real position of the galaxies and their stellar mass. The relative excess caused by clustering is then estimated comparing the injected flux density with the stacked flux density measured in this simulated map. This method neglects the diversity of the SEDs, the non-linearity of the M-SFR relation, and the scatter around it. The trend between this method and our simulation agrees overall. However, our simulation found slightly lower values at z > 2. This disagreement of ~15% is hard to explain and could come from the cosmic variance, a systematic effect at small scale in the real catalogs (incompleteness, problem of deblending), or a poor description of the small-scale clustering at z > 2 in our simulation. The other method (red diamonds) is based on a stacking in map space and a fit of the resulting radial profile by both a point-like and an extended clustered component (Heinis et al. 2013; Béthermin et al. 2015a; Welikala et al. 2016). The results are similar to those obtained with the method described previously and our simulation at 250 μm and 350 μm. At 500 μm, this method predicts higher values than the other ones at z> 2. However, at 500 μm, the decomposition is hard to perform, since the typical scale of the intra-halo clustering is close to the size of the beam.

Larger biases can be found if we stack fainter or more clustered galaxy populations. It can reach 90% if we stack, for example, all sources with M > 109M at 500 μm. The impact of clustering should thus be carefully checked while stacking Herschel data. Our simulation will be particularly well suited to checking the accuracy of the stacking approaches in future studies.

thumbnail Fig. C.1

Relative excess of flux density in stacking measurements caused by clustered neighbors as a function of redshift. These values were derived for a stellar mass selection (3 × 1010M). The black squares are the results from our simulation. The blue triangles were derived using a simulated map built from the real positions and the stellar masses of the galaxies in the COSMOS field (Béthermin et al. 2015a). The red diamonds were derived using a decomposition of the real stacked images into a point-like and an extended clustered component (Béthermin et al. 2015a).

Open with DEXTER

All Figures

thumbnail Fig. 1

Left: evolution of the stellar mass function used in our simulation with redshift. The construction of this stellar mass function is explained in Sect. 2.2. The solid line is the SMF computed from the catalog in a bin Δz = 0.4. The dashed line is the analytic SMF at the mean redshift of the bin. The cut at low mass (M< 107M) is due to the halo mass limit of the dark-matter simulation. Right: relation between the halo mass and the stellar mass from our abundance-matching procedure (see Sect. 2.3). The error bar indicates the 0.2 dex scatter on this relation.

Open with DEXTER
In the text
thumbnail Fig. 2

Fraction of star-forming galaxies versus stellar mass at various redshifts. The data points are from Davidzon et al. (2017). The solid line is a fit by the parametric form described in Sect. 2.4.

Open with DEXTER
In the text
thumbnail Fig. 3

Evolution of the mean intensity of the radiation field U (see Magdis et al. 2012) that is used in the current updated model. The evolution of main sequence galaxies and starbursts are plotted in blue (solid line) and red (dashed line), respectively. Also shown are the measurement of Béthermin et al. (2015a, triangles for main sequence galaxies and squares for starbursts), Magdis et al. (2012, diamonds), da Cunha et al. (2010, asterisk), and Ciesla et al. (2014, circle). This figure is adapted from Béthermin et al. (2015b).

Open with DEXTER
In the text
thumbnail Fig. 4

Differential number counts from 70 μm to 500 μm. The counts are multiplied by to reduce the dynamic range of the plot and to highlight the plateau at high flux densities where the Euclidian approximation is valid (e.g., Planck Collaboration Int. VII 2013). The solid black line is the prediction from our simulated catalog. The red line is derived from the extraction of sources in simulated Herschel maps using a method similar to that used by Béthermin et al. (2012a). The error bars on the prediction of the simulation are derived assuming Poisson statistics. The source extraction procedure is limited by confusion and only bright sources can be extracted reliably. From 70 μm to 160 μm, the light blue circles, the dark blue diamonds, and the green triangles are the Spizter counts of Béthermin et al. (2010a), Herschel/PEP counts of Berta et al. (2011), and Herschel/GOODS counts of Magnelli et al. (2013), respectively. From 250 μm to 500 μm, the light blue circles, the dark blue diamonds, the green triangles, turquoise squares, and purple asterisks are the Herschel/SPIRE measurements of Béthermin et al. (2012c), Oliver et al. (2010), Glenn et al. (2010), Clements et al. (2010), and Valiante et al. (2016), respectively. The SCUBA2 measurements at 450 μm of Chen et al. (2013), Casey et al. (2013), Geach et al. (2013), and Zavala et al. (2017) are shown using gold down-facing triangles, orange right-facing triangles, red open circles, and brown stars, respectively.

Open with DEXTER
In the text
thumbnail Fig. 5

Integral number counts at 850 μm and 1.2 mm. The black solid line is the intrinsic counts from our simulated catalog. The red solid line is the result of source extraction from the simulated map of single-dish instrument (see Sect. 3.3). The number counts extracted from high-resolution interferometric data (int.) are colored in blue and purple (Oteo et al. 2016; Karim et al. 2013; Simpson et al. 2015; Fujimoto et al. 2016; Hatsukade et al. 2013; Aravena et al. 2016), while we use orange for single-dish (SD) results (Coppin et al. 2006; Geach et al. 2017; Lindner et al. 2011; Scott et al. 2012).

Open with DEXTER
In the text
thumbnail Fig. 6

Average fraction of the flux density emitted by the brightest galaxy in the beam as function of the flux density measured in our simulated map with a limited angular resolution (red diamonds). We present our results for Herschel/SPIRE at 250 μm (left) and 500 μm (middle), and for NIKA2 at 1.2 mm (right). The gold squares show the average contribution of the galaxies physically related (| Δz | < 0.01) to the brightest galaxy in the beam. The orange open circles are the sum of the flux density fraction from the brightest galaxy and the other galaxies at the same redshift. The black downward-facing triangles is the flux density fraction from the brightest galaxy measured by Scudder et al. (2016). They used a different definition from ours and divided the flux density of the brightest galaxy by the total of the flux density of all the galaxies in a 1 FWHM radius. The blue upward-facing triangles are the results from our simulation assuming their definition (see Sect. 3.4).

Open with DEXTER
In the text
thumbnail Fig. 7

Effect of the SFR cut on the number counts at 850 μm. The data points are similar to those of Fig. 5. The black solid line is our standard model, that is, with a SFR limit of 1000 M/yr, and the brown dashed line is the model without SFR limit.

Open with DEXTER
In the text
thumbnail Fig. 8

Comparison between the measured redshift distributions and the predictions of our simulation. The orange histograms are the intrinsic redshift distributions from our simulation. The data points are extracted from Berta et al. (2011) at 100 and 160 μm, Béthermin et al. (2012c) at 250 μm, Geach et al. (2013) and Casey et al. (2013) at 450 μm, Wardlow et al. (2011), Simpson et al. (2014), and Chapman et al. (2005) at 850 μm, and Michałowski et al. (2012), Smolčić et al. (2012), and Yun et al. (2012) at 1.1 mm. Figure adapted from Béthermin et al. (2015b).

Open with DEXTER
In the text
thumbnail Fig. 9

Illustration of the impact of the 24 μm-prior extraction on the redshift distribution at 250, 350, and 500 μm. A flux density cut of 20 mJy was used to select the SPIRE sources. The red line is the intrinsic redshift distribution, while the blue line is obtained after extracting the sources in our simulated Herschel/SPIRE maps using 24 μm positions as a prior. The histograms are normalized in order to have dN/ dz dz = 1. We compare these model predictions with the observational redshift distribution of Béthermin et al. (2012c, black diamonds). The data points were extracted in real Herschel data using position, 24-μm flux density, and redshift as a prior. This more complex method was chosen to avoid the potential biases associated with a source extraction using only 24-μm positions as a prior (see discussion in Sect. 4.1).

Open with DEXTER
In the text
thumbnail Fig. 10

Number counts per redshift slice. Purple, blue, orange, and red are used for the z< 0.5, 0.5 <z< 1, 1 <z< 2, and z> 2 slices, respectively. The data points are from Berta et al. (2011) at 70, 100, and 160 μm and Béthermin et al. (2012c) at 250, 350, and 500 μm. The dotted lines are the intrinsic distributions in the simulation and the solid lines are the counts extracted using 24 μm priors (see Sect 3). Figure adapted from Béthermin et al. (2012c).

Open with DEXTER
In the text
thumbnail Fig. 11

Evolution of the obscured star formation density as a function of redshift. The black line is the result of our simulation. We show for comparison the infrared measurements of Sanders et al. (2003, diamond), Takeuchi et al. (2003, triangle), Magnelli et al. (2011, squares), Magnelli et al. (2013, circles), Gruppioni et al. (2013, stars), and Rowan-Robinson et al. (2016, circles with cross). We also plot the UV estimates corrected for dust attenuation from Wyder et al. (2005, asterisk), Schiminovich et al. (2005, downward-facing triangles), Robotham & Driver (2011, right-facing triangle), Cucciati et al. (2012, left-facing triangle), Dahlen et al. (2007, hourglass), Reddy & Steidel (2009, hat), and Bouwens et al. (2012a,b, bowtie). All the data have been converted to Chabrier (2003) IMF and Planck Collaboration XIII (2016) cosmology.

Open with DEXTER
In the text
thumbnail Fig. 12

Pixel histograms of the Herschel/SPIRE maps in COSMOS (black) and comparison with our simulation using the same instrumental noise map. The red histograms are the result of our simulation and the blue ones are the histogram obtained after randomizing the position of the sources to illustrate the impact of clustering.

Open with DEXTER
In the text
thumbnail Fig. 13

Power spectrum of cosmic infrared background anisotropies at 250, 350, and 500 μm. The blue solid lines are CIB anisotropies measured in our simulation and 1σ confidence regions are represented by the dashed lines. The black diamonds and the red triangles are the measurements of Viero et al. (2013b) and Planck Collaboration XXX (2014), respectively.

Open with DEXTER
In the text
thumbnail Fig. 14

Differential number counts of red Herschel sources. The black diamonds are the measurements from Asboth et al. (2016). The red point is the result of an end-to-end analysis of our 2 deg2 simulated maps including clustering and instrumental noise. The turquoise, dark blue, and gold lines are the number counts in a simulated catalog of 274  deg2 without noise, with instrumental noise only, and with both instrumental and unclustered confusion noise, respectively. Blue and purple arrows illustrate the impact of instrumental and confusion noise, respectively.

Open with DEXTER
In the text
thumbnail Fig. 15

Left panel: redshift distribution of the red sources. The blue and black histograms are the distributions of the sources matching the red criteria before and after including noise (both instrument and confusion noises), respectively. Middle panel: distribution of the red excess (D = 0.92 M500−0.392 M250). Right panel: distribution of the 500 μm flux density. The blue and black histograms show the intrinsic (without including noise) 500 μm flux densities (middle panel) and D values (right panel) of the intrinsically red (red criteria applied before adding noise) and observed red (red criteria applied after adding noise) sources, respectively. The red histogram shows the distributions of the observed, that is, after including noise, 500 μm flux densities and D values of observed red sources. The difference between the black and red histogram illustrates that sources observed above the S500 and D thresholds of Asboth et al. (2016, vertical black dotted line) usually have lower intrinsic values.

Open with DEXTER
In the text
thumbnail Fig. A.1

Left: differential number counts at 100 μm. We used the same data points as in Fig. 4. The red dashed line is before the correction at low redshift of the position of the main sequence (see Appendix B) and the black solid line is the final version of our model. Since this effect is only significant at bright flux density, where the number of sources is small, we used an unclustered 10 deg2 to reduce the statistical uncertainties in our model predictions. Right: star formation density as a function of redshift. The data are similar to Fig. 11, but we added the prediction of our model before the low-redshift correction of the main sequence (red dashed line).

Open with DEXTER
In the text
thumbnail Fig. C.1

Relative excess of flux density in stacking measurements caused by clustered neighbors as a function of redshift. These values were derived for a stellar mass selection (3 × 1010M). The black squares are the results from our simulation. The blue triangles were derived using a simulated map built from the real positions and the stellar masses of the galaxies in the COSMOS field (Béthermin et al. 2015a). The red diamonds were derived using a decomposition of the real stacked images into a point-like and an extended clustered component (Béthermin et al. 2015a).

Open with DEXTER
In the text

Current usage metrics show cumulative count of Article Views (full-text article views including HTML views, PDF and ePub downloads, according to the available data) and Abstracts Views on Vision4Press platform.

Data correspond to usage on the plateform after 2015. The current usage metrics is available 48-96 hours after online publication and is updated daily on week days.

Initial download of the metrics may take a while.