Asantha Cooray, UC Irvine Far-Infrared 2016 AAS 250Μm

Extragalactic Sciences with the Far-Infrared Surveyor Architecture Consensus

70% 30%

Architecture Consensus CALISTO SPIRIT 70% Some (old) new ideas appeared,30% though: • Polarimetry • Very high spectral resolution

15 December 2016 3rd FISICA Workshop: London, UK 7

CALISTO SPIRIT Some (old) new ideas appeared, though: Asantha• PolarimetryCooray • Very high spectral resolution

15 December 2016 3rd FISICA Workshop: London, UK 7 Outline Quick summary of some key extragalactic astrophysics results from the Herschel Space Observatory

Science opportunities with Far-IR Surveyor (single aperture and an interferometer)

New things: (I) 3D spectral line intensity mapping

(II) Far-infrared probes of reionization (especially molecular Hydrogen 10 < z < ~15)

review of dusty star-forming galaxies Casey, Narayanan, Cooray 2014 Physics Reports

Asantha Cooray, UC Irvine Far-infrared 2016 AAS 250µm

350µm

500µm

10 arcmin GOODS-N • FIRAS total background Source Counts 250, 350, 500 µm 15%, 10%, 6%

• P(D) 250, 350, 500 µm 65%, 60%, 45% confusion confusion σ σ 5 5 • Stacking: 250, 350, 500 µm 80%, 80%, 85%

Of course: The remainder are the most interesting sources! E.g. z > 3 galaxy populations

Resolving the extragalactic background spectrum

Asantha Cooray, UC Irvine Far-Infrared 2016 AAS 8 D. Brisbin et al.

µ Figure 3. GOODS-N galaxy SEDs plotted as in Figure 2. The 500 m Figure 5. Spectral energy distributions of the most luminous sources in measurements for GOODS sources A, C, E, F, H, K, and L are compatible GOODS North (a) to (d) Lockman North (e) to (k). As in Figures 2 and 3, with zero, so here we plot only the upper error bar, barely visible within the we solely plot the upper error bars at 500 µm for sources (a) to (e), whose diamond symbol. As in Figure 2, we weighted the 24 µm data a quarter as lower error bars are compatible with zero. We similarly plot an upper error heavily as the longer wavelengths. Observational data shortward of 24 µm bar at 170 µm for LN (k). Luminosities and positions for these sources are is plotted for reference but not used in fitting. presented in Table 4. Seven of these eleven sources are triply secure; two of these are also included in Figures 2 and 3. A brief description of sources (c), (d), (f) and (i) is provided in Section 6. Although some of the ultraluminous sources exhibit significant AGN activity, we have applied S&K model fits, as discussed in Section 7. As in Figures 2 and 3 we have weighted the 24 µm HLIRGS data only one quarter as heavily as the longer wavelengths. Observational 2000data shortward of 24 µm is plotted for reference but not used in fitting.

ULIRGS Star-Formation 200 Rate in solar masses per year LIRGS 20

2 (~Milky-way SFR) Far-IR Surveyor

10 Figure(i) ULIRGS/HyLIRGS 4. Infrared source typically luminosities have (integrated about over ~10 8-1000 solarµm) masses in in stars GOODS-N and LN plotted as a function of redshift for all sources detected at all three SPIRE wavelengths. Diamonds indicate GOODS-N luminosities (ii) So the time scale for star-formation is [M*/(dM*/dt)] ~ 5 to 100 Million years obtained from SEDs ﬁtted by S&K models, and triangles indicate similar luminosities for LN sources. The solid line shows the growth (star-bursting of luminosity galaxies!) distance squared with z. It serves as a rough lower bound to the luminosities in our selectionWhat of observed kind sources; of the scattergalaxies of data points did about we the detect with Herschel? curve provides a visual impression of the uncertainties in those luminosities.

Asantha Cooray, UC Irvine Far-infrared 2016 AAS

c 2010 RAS, MNRAS 000, 2–?? More Generally: When are galaxies at high-z mergers?

In the local Universe ~100% of SMGs, 24 sources starbursts are driven by gas- rich galaxy mergers.

BzKs But at z ~ 1 to 2, observations show that some starburst galaxies are simple disks.

Is there a different mechanism to trigger a Hayward, Narayanan et al. in prep. starburst at high redshifts? (theorists: cold accretion mode) Tacconi, L. J. et al. 2008, ApJ, 680, 246 Dekel, A. et al. 2009, ApJ, 703, 785 (a) What fraction of starbursts are mergers vs. cold ﬂows? Hopkins, Younger, Hayward,(b) Do the DN, mergers Hernquist evolve 2010 differently from cold ﬂows? what stops the starburst? What are Dusty Star Forming Galaxies (DSFGs)? LOCAL ULIRGS REVIEW: SANDERS, D. AND MIRABEL, I. 1996, ARAA, 34, 749 Asantha Cooray, UC Irvine Far-infrared 2016 AAS

Desika Narayanan Thursday, May 12, 2011 Figure 3: (a): density-ionization dia- gram (Spinoglio et al. 2009). SPIRE- AGNs FTS/PACS analysis will result in data re- −2 10 lated to many of the spectral lines that NGC 1068

F10214 probe a wide range of conditions. (b): 10−3

m) / FIR [OIV]/FIR limit as a test of AGN contri- µ bution. From Sturm et al. (2010). The −4 Mrk 231 10 points with labeled names are PACSspec [O IV] (25.91 SFGs Arp 220 detections or upper limits during the Sci- 10−5 (b) Why is infrared spectroscopy the best tool to isolate star (a) M82 ence Demonstration Phase observations. formation and accretion? (c): The [CII]158/[OI]63 line ratio vs. Spinoglio et al. 108 109 1010 1011 1012 10 13 Infrared(fine(structure(lines( LFI R [ L ] the [OIII]88/[OI]63 ratio: PACSspec ob- IR fine structure lines: - separate different servations of low-redshift galaxies are physical mechanisms, shown as different types of points (Star- - fully cover the ionization-density burst galaxies: filled pentagons,“non- parameter space Seyfert” active galaxies: filled circles; - do not suffer from “pure” Seyfert 2’s: filled squares, high extinction (c) (d) Balmer line ratio galaxies: open trian- Spinoglio & Malkan ‘92 gles, Seyfert 1’s: filled triangles). Here predicted line intensities of IR lines in active and photoionisation models of AGNs and # Stars( b.2h.(accre4on( starburst galaxies, before starburst galaxies are shown as blue and ISO launch. black grids, respectively, with logarith- density ioniza/on# mic values of the density n and ioni- sation potential U of indicated on the

Luigi#Spinoglio#B#ESO,#May#26th,#2015;## 18# plot. From Spinoglio et al. (2014). (d): Karin#also#showed#this# Far-IR rich in spectral lines The [CII]158/[OI]63 line ratio vs. the [OIII]88/[OIV]26 ratio. From Spinoglio Asantha Cooray, UC Irvine far-infrared 2016 AAS et al. (2014). The plot follows left panel.

tentative detections of [OI]145.5 and [NII]122.1. The subsequent analysis, also in Danielson et al. 2011, demonstrated that the ISM in this z =2.3 DSFG is multi-phase, with bulk of the gas mass in a low-metallicity, cool (T 25 k) and extended component. A higher metallicity, warmer phase ⇠ (T 45 k) traces the regions of current star formation. ⇠ The [CII] emission from SMM J2135 is 10 times more luminous than expected from local ⇠ ULIRGs (e.g. Luhman et al. 2003; Brauher et al. 2008). Indeed, the data currently available suggest that at a ﬁxed far-IR luminosity, high-redshift DSFGs generally have enhanced [CII] compared to AGN-dominated systems that are no longer actively star-forming (Figure 3 bottom left). It has been speculated that this supports the presence of a low-metallicity gas reservoir. Differ- ences on [CII] to CO ratios at low vs. high redshifts at the same ﬁxed far-IR luminosity suggest that this reservoir may be more extended than those seen around the compact starbursts at z 0 ⇠ (Hailey-Dunsheath et al. 2010; Ivison et al. 2010). Further support for such a suggestion has also come from sub-mm, radio and CO sizes and near- and mid-IR colours and spectra (e.g. Tacconi et al. 2008; Menendez-Delmestre´ et al. 2009). The lower dust content leads to a lower attenuation to UV photons, potentially making the [CII]-emitting region larger and lowering the far-IR luminosity (Maiolino et al. 2009). However, [CII]/LFIR is also high in local LIRGs and this is likely not due to an extended low-metallicity gas reservoir. Thus, the mechanism that suppresses [CII] in ULIRGs cannot also operate in LIRGs. The most recent interpretation of the drop in [CII]/LFIR

6 SDSS

SPICA ~ 103 galaxies (but limited z range)

Far-IR Surveyor ~ 106 galaxies?

Luigi#Spinoglio#B#ESO,#May#26th,#2015# 29# [slide adapted from Matt Malkan] Far-IR rich in spectral lines

Asantha Cooray, UC Irvine far-infrared 2016 AAS N (>S) N

N (>S) N ⨂ = Probability

Flux mu Flux

Intrinsic Source Count Magnification Cross Section Lensed Source Count

Negrello et al. 2010 Science Julie Wardlow et al. 2013, ApJ (UCI ex-postdoc)

Lensing galaxy selection at sub-mm wavelengths > 95% efﬁcient

The Nature of Brightest high-z Herschel Galaxies

Asantha Cooray, UC Irvine far-infrared 2016 AAS Jae Calanog UCI PhD 2014

We now have 60 images like these in total with Keck/LGS AO

Figure 5: Herschel -selected bright, gravitationally Keck lensed LGS-AO SMGs conﬁrmed Imaging by Keck adaptive op- tics 2.1 µm or HST/WFC3 1.1 µm images (Calanog et al. 2013 in prep). Tick marks are separated Fu et al.in 2012; intervals Bussmann of 100 in all et panels.al. 2012; See Fu Fig. et 6al.for 2013; the lensCalanog model et of al. G12v2.30. 2014; Timmons et al. 2015 Asantha Cooray, UC Irvine far-infrared 2016 AAS

Figure 6: Lens modeling reveals a highly structured galaxy in the source plane of G12v2.30 at z = 3.26 (Fu et al. [48]). (a) Observed K-band image with the foreground lenses subtracted. (b) Lens model. Critical curves are in red and caustics are in blue. The box delineates the on-sky area covered by the source-plane image in panel (d). (c) Residual image. (d) Parameterized source mor- phology (grey scale) and a direct inversion of the observed image (red contours). For comparison, we also show the distributions of the molecular gas (green) and starburst-heated dust (purple) from JVLA CO J =1 0 and SMA 880-µm observations. !

10 H-ATLAS: 650 sq. degrees. ~2 lensed Planck CSC Starbursting knot in a sources.Far-infrared One spectroscopyin HerMES over of a 370 lensed sq. starburstdegrees. 3 spiral galaxy. Disk is mostly Table 1. Continuum ﬂux densities. an old stellar population.

Wavelength Sν Notes

100 µm41± 4a mJy HerschelLensedPACS 160 µm180± 14a mJy HerschelopticalPACS arc 250 µm347± 25b mJy Herschel SPIRE 350 µm378± 28b mJy Herschel SPIRE 500 µm268± 21b mJy Herschel SPIRE 870 µm30.2± 5.2 mJy SMA 2mm1.2± 0.1 mJy IRAM PdBI 3.5 mm 200 ± 170 µJy CARMA 4.3 cm 350 ± 30 µJy VLA 21 cm 1.95 ± 0Lensed.24 mJy Herschel-detected VLA (FIRST) a Errors include 3- and 5-per-centsource (starbursting calibration uncertaintie s at 100- and 160-µm, respectively.clump) b Errors include the contribution due to confusion and a 7- per-cent1” calibration uncertainty has been added in quadra- ture (Valiante et al., in preparation). Figure 3. The continuum-subtracted Herschel SPIRE FTS spectrum. No z=1.68,other features z determined correspond to expected from transitionthe Herschel-SPIRE/FTS lines. Inset:Zoomedinon the region indicated by the red line in the parent plot. Velocity scale corre- spectrumsponds to [C withII]atz the=1. 676158.Thebest-ﬁtsincproﬁleisoverlaidinred. micron CII line George et al. 2014; Timmons et al. 2015

Nick Timmons UCI PhD 2017 Herschel Lensed Sources

Asantha Cooray, UC Irvine Far-infrared 2016 AAS

Figure 4. CO J =2→1 spectrum from CARMA, binned to 20.8 MHz, with the CO J =3→ 2 IRAM PdBI spectrum shown in red, binned to 40 MHz and put on the same brightness temperature (Tb)scale. Figure 2. Far-IR-through-radio SED of HATLASJ132427, with the best-fit synchrotron and power-law temperature distribution dust model shown in black. The FTS spectrum is shown in blue. Inset:15-arcmin× 15-arcmin accurate to within ∼ 15 per cent. We obtained an 8 σ detection of false-colour image of HATLASJ132427 in the three Herschel SPIRE filters. line emission, close to the expected frequency – see Fig. 4 – thus confirming the redshift, with z =1.676 ± 0.001.The34.8GHz GBT line is therefore presumably spurious. galaxies, this line is likely to be the most significant; indeed, it is We also imaged the source in CO J =3→ 2 with the IRAM often the only transition detected (e.g. Ivison et al. 2010b). PdBI, also during 2012 November. We obtained 1.1 hr of integra- The SPIRE FTS spectrum of HATLAS J132427, shown in tion time, using all six of the 15-m antennas, this time in D con- Fig. 3, displays a 7.5 σ marginally resolved (1.2 GHz spectral reso- figuration – the most compact. The observing frequency was set lution) emission line at (709.9±0.4) GHz. Attributing this to [C II] to 129.028 GHz, corresponding to the redshifted frequency ofCO indicates a redshift, z =1.677 ± 0.001.Nootherlineswerede- − J =3→ 2 (νrest =345.796 GHz) for z =1.68.Again,wefounda tected, with 3σ < 200,230,150,600Jykms 1 for [O III]88µm, bright 3 σ emission line at the expected frequency, (Fig. 4). [N II]122µm, [O I]145µm, and [N II]205µmrespectively,mean- The two line profiles are consistent with one another; neither ing that this redshift remained tentative at this point. can be described well with a single Gaussian, suggesting thattheir shape is due to either a merger or a rotating, gas-rich disk (e.g. Engel et al. 2010; Ivison et al. 2013). The line width (deducedfrom 4.2 Redshift confirmation via CO − fits using a single Gaussian), (500± 140) km s 1 FWHM,istypical To verify the SPIRE FTS redshift determination, we used the of those seen for DSFGs (Greve et al. 2005). Combined Array for Research in Millimeter-wave Astronomy (CARMA) to search for the CO J =2→ 1 line (νrest = 230.538 GHz), which should be redshifted to approximately ν = obs 5DISCUSSION 86.0 GHz for z =1.68.Observationswerecarriedoutusing the 3-mm receivers during 2012 November 23 in D configuration Characterisation of the dust emission of HATLAS J132427 was (11–146 m baselines), with 2.3 hr spent on source. The blazars, performed by fitting the power-law dust temperature model of 1310+323 and 0927+390, were used for complex gain calibration Kovács et al. (2010) to the continuum flux densities detailedinTa- and to derive the bandpass shape. Absolute flux densities should be ble 1. Derived quantities are detailed in Table 2.

⃝c 0000 RAS, MNRAS 000,000–000 H-ATLAS: 650 sq. degrees. ~2 lensed Planck CSC Source reconstructed: SMA/JVLA (red) sources. One in HerMES over 370 sq. degrees. and Keck/Hubble (green/blue) n = 103 Galactic SF Regions L[C II ] / LCO(1 0) = 4100 Normal Galaxies n = 102 10− 2 Starburst Nuclei n = 104 Local ULIRGs HATLAS J132427 LESS J033229.4 BR1202-0725 SMM J2135 5

FIR n = 10

L HDF 850.1 / ] HFLS 3 II − 3 C

[ 10 SDP.81 L ID 141 NBv1.43

n = 106

− 4 4 3 2 10 G0 = 10 G0 = 10 G0 = 10

10− 7 10− 6 10− 5 LCO(1 0) / LFIR 1” z=1.68, z determined from the Herschel-SPIRE/FTS spectrum with the 158 micron CII line George et al. 2014; Timmons et al. 2015

Herschel Lensed Sources

Asantha Cooray, UC Irvine Far-infrared 2016 AAS PACS spectroscopy of z > 1 galaxies - mainly lensed galaxies - about 50 targets - Mostly undetected - detections are at best 3 to 5 sigma

70 to 500 micron spectroscopy was not easy with Herschel - tons of upper limits over close to 500 hours unpublished.

Wardlow et al. in prep

Asantha Cooray, UC Irvine Far-infrared 2016 AAS 500 um peaked sources S250 < S350 < S500: z > 4? 3 *Confusion reduced S(500) – fS(250) 250 um 350 um 500 um * HST WFC3+ACS 250 um 350 um 500 um * FLS3 z=6.34

Dowell et al. 2014 ApJ technique IRAM 1.1mm

z Figure= 6.34 1. Left: DustyThe three-color Starburst image using Galaxy HST/ACS in combined HerMES F625W and F814W (blue), HST/WFC3 combined F160W, F125W and F105, and Keck/NIRC- 2 Ks-band LGS-AO (red)Riechers, images. D. et Note al. Nature the clear 2013; detection Cooray ofet twoal. 2014 galaxies close to HFLS3 shown here in terms of the IRAM/PdBI 1.1 mm (rest-frame 158 1 Asantha Cooray, UC Irvine µ m) emission. The r.m.s. uncertainty in the PdBI A-array Far-infrared con 2016figuration AAS data is 180 µJy beam− and the contours are shown in steps of 3σ starting at 5σ. The instrumental beam is shown to the bottom right with FWHM of 0.35′′ 0.23′′. Right: The three-color GALFIT residual map where we remove models for the HST/ACS-detected galaxies in HST/WFC3. Here we show the combination× of ACS/F625W+F814W (blue), WFC3/F105W (green) and WFC3/F160W (red). Both G1 and G2 are detected in the combined ACS/F625W and F814W stack, consistent with the scenario that both G1 and G2 are at z<6 and G2 is not the optical counterpart of HFLS3, as was previously assumed. We find a marginal detection of rest-UV emission at the location of HFLS3 (labeled R2) and a higher significance diffuse emission 0.′′5 to the South-West of HFLS3 (labeled R1). We use WFC3 fluxes and ACS upper limits of R2 for combined SED modeling of HFLS3 with far-IR/sub-mm flux densities. We detemine a photometric redshift for R1 and find it to be consistent with emission from either a galaxy at z 6 or a dusty galaxy at z 2. ∼ ∼ relative to ACS would suggest the presence of detectable rest- UV emission from HFLS3. As shown in Fig. 1 (right panel) we find excess emission primarily 0.′′5 to the South-West of HFLS3 (labeled R1). We also find some marginal evidence for excess emission near the 1.1-mm peak (labeled R2), with detection levels between 2.5 to 3.2σ. In Table 1 we summarize GALFIT and other intrinsic properties of the two foreground galaxies G1 and G2 as well as the residual emission R1 and R2, with R2 emission assumed to be from HFLS3. We also use the latter for a combined UV to far-IR SED modeling with MAGPHYS (da Cunha et al. 2008). We also model the SED of R2 to determine its photometric redshift and address its asso- ciation with HFLS3.

4. LENS MODELING We use the publicly available software GRAVLENS (Keeton 2001) to generate the lens model. As the background source is not multiply-imaged, and remains undetected in the rest-frame optical, we simply model the highest resolution Figure 2. Spitzer/IRAC 3.6 µm image from Riechers et al. (2013) showing a detected source. The contours on the intensity scale show the region of G1 IRAM/PdBI 1.1-mm continuum emission map. This map and G2 (blue) and R1 and R2 (red). IRAC 2′′ PSF is marked with a white is currently our highest resolution view of HFLS3, and the circle. ∼ source is resolved in these data. The magnification factor we determine here with lens modeling is the value for the mm- detected galaxies near HFLS3 and to see if there is any ex- wave continuum emission. It could be that HFLS3 will be cess emission in WFC3 data relative to the ACS images. Us- subject to differential magnification, with different emission ing GALFIT on the individual Hubble/ACS and WFC3 frames components within the galaxy subject to different magnifica- proved to be difficult because the output models tend to over- tion factors (e.g, Serjeant 2012; Hezaveh et al. 2012). This fit regions of low signal in which HFLS3 is expected to re- is especially true if the different components associated with side. To remedy this, we stacked the HST/ACS in two bands HFLS, such as dust, gas, and stellar mass, have peak intensiti- to increase the signal-to-noise ratio and to model the fore- ties that are offset from each other, as in the case of a complex ground galaxies in the combined F625W and F814W images. merging galaxy system. Under the assumption that the stacked model best represents To simplify the lens modeling, we use singular isothermal the two foreground galaxies, we then subtracted the stacked ellipsoidal (SIE) models to fit for the Einstein radius and po- model from individual HST/ACS and WFC3 frames, with the sitions of the two lens galaxies. The position angles and ellip- flux density at each wavelength allowed to vary as an over- ticities for G1 and G2 are fixed to the values derived from pro- all normalization in GALFIT models. Any excess in WFC3 file fitting using GALFIT, but their masses are allowed to vary 250 um 350 um 500 um *

! ! ! Asantha Cooray, UC Irvine Far-infrared 2016 AAS ! ! ! !

! 3 3

z=2.1

z=6.3

Weakly lensed by two z=2.1

galaxies with magnification 12 Figure 1. Left:FigureThe 1.three-colorLeft: The image three-color using image HST/ACS using combined HST/ACS combinedF625W and F625W F814 andW(blue),HST/WFC3-IRcombinedF160W,F125WandF105W,andL F814WFIR = (blue), 6X10 HST/WFC3 L⊙ combinedM F160W,DUST F125W> 109 and M F105,⊙ and Keck/NIRC- Keck/NIRC-2 K2 Ks-bands-band LGS-AO LGS-AO (red) (red) images. images. Note Note the the clear clear detection detection of of two two ga galaxieslaxies close close to to HFLS3 HFLS3 shown shown here here in in terms terms of of the the IRAM/PdBI IRAM/PdBI 1.1 1.1mm mm (rest- (rest-frame 158 1.6 +/- 0.3 − 1 10 frame 158 µm)µ emission.m) emission. The The r.m.s. r.m.s. uncertainty uncertainty in the in PdBIthe PdBI A-array A-array confi congurationfiguration data data is 180SFR is 180µJyµ beamJy~ beam13001−andand theM the contours⊙/yr contours areM are shownSTARS shown in steps in~ steps5X10 of 3 ofσ 3startingσ startingM⊙ at at 5σ. The 5σ.TheinstrumentalbeamisshowntothebottomrightwithFWHMinstrumental beam is shown to the bottom right with FWHMof 0.35 of′′ 0.35× 0.23′′ ′′0.23. Right:′′. Right:The three-colorThe three-color GALFITGALFIT residualresidual map wheremap where we remove we remove models models for the × MGAS ~ 1011 M⊙ for the HST/ACS-detectedHST/ACS-detected galaxies galaxies in HST/WFC3. in HST/WFC3. Here we Here show wethe show combination the combination of ACS/F625WT ofDUST ACS/F625W +=F814W 55+ ±F814W (blue),10 K (blue), WFC3/F105W WFC3/F105W (green) (green) and WFC3/F160W and WFC3/F160W (red). (red). Both G1Both and G2 G1 are and detected G2 are[G2 detected in the identification combined in the combined ACS/F625W ACS/F625W in aR13nd F814W as and stack, F814W consistent stack, consistent with the with scenario the scenario that both that G1 both and G2 G1 are and at G2z< are6 atandz< G26 isand not G2 is not the the least obscuredoptical region, counterpart or the rest-frame of HFLS3, optical as was counte previouslyrpart, assumed. of HFLS3, We asfi wasnd a previously marginal detection assumed. of We rest-UV find a emission marginal atdetection the location of rest-UV of HFLS3 emission (labeled at R2) the and a higher location of HFLS3signi (labeledficance diffuse R2) andK-band emission a higher 0. significance′′ID5 to of the South-WestFLS3 diffuse incorrect] emission of HFLS30′′. (labeled5 to the South-West R1).No We use evidence of WFC3 HFLS3fluxes (labeled for and ACS R1).a quasar/massive upper We use limits WFC3 of R2 fluxes for combinedand AGN!ACS upper SED modeling of limits of R2 forHFLS3 combined with SED far-IR/sub-mm modeling offlux HFLS3 densities. with We far-IR/s detemineub-mm a photometric flux densities. redshift We fordetemine R1 and afi photometricnd it to be consistent redshift for withR1 emission and find from it to either be consistent a galaxy at z 6 or a dusty galaxy at z 2. ∼ with emission from either a galaxy∼ at z ∼ 6 or a dusty galaxy at z ∼z2. = 6.34 Dusty Starburst Galaxy in HerMES IRAF.STSDAS pipeline for flat-fielding and cosmic-ray re- G1 and G2).relative This to is ACS similar would to what suggest was the previously presence ofreported detectable rest- jection. Individual exposures in each of the filters were com- with Keck/NIRC2UVRiechers, emission LGS-AO D. fromet al. imaging HFLS3.Nature 2013; data, As shown Cooray with thein et Fig. al. southern 2014 1 (right panel) bined with ASTRODRIZZLEAsantha(Fruchter Cooray, UC Irvine & et al. 2010) and we component we (G2) fi nd taken excess to emission be the rest-frame primarily Far-infrared 0. optical′′5 2016 to the counter-AAS South-West of produced images at a pixel scale of 0′′. 06 from the native scale part to HFLS3HFLS3 (Riechers (labeled et R1). al. 2013). We also Iffi thisnd some assumption marginal is evidence of 0′′. 13 per pixel. For flux calibration we made use of the correct wefor expect excess the emission southern near component the 1.1-mm to be peak invisible (labeled in R2), with latest zero-points from STScI, with values of 26.27, 26.26 the shortestdetection wavelength levels images, between as 2.5 it tois 3.2 a Lymanσ. In Table drop-out 1 we atsummarize and 25.96 in F105W, F125W and F160W, respectively. Simi- GALFIT and other intrinsic˚ properties of the two foreground wavelengthsgalaxies shorter G1 than and 8900 G2 asA. well Here, as the however, residual we emission have R1 and larly, Hubble/ACS imaging data were flat-fielded, cosmic ray- detected both galaxies in Hubble/ACS images, establishing rejected and charge transfer efficiency (CTE)-corrected with R2, with R2 emission assumed to be from HFLS3. We also that G2 isuse a galaxy the latter at z< for a combined5.Sincethese UV toHubble far-IR SEDobserva- modeling with the pipeline CALACS (version 2012.2). Exposures were tions, we have reanalyzed the Keck/LRIS spectrum shown in ′′ MAGPHYS (da Cunha et al. 2008). We also model the SED of remapped with ASTRODRIZZLE to a pixel scale of 0. 03.The Riechers etR2 al. to (2013) determine with itsz photometric=2.1 CIV (1549 redshiftA)˚ and and address OIII] its asso- ACS zero points used from an online tool are 25.94 and 25.89 ′′ (1661, 1666ciationA)˚ emission with HFLS3. lines within 1 of HFLS3. We now for F814W and F625W, respectively. find some marginal evidence that this emission is extended, The final Hubble mosaics were astrometrically calibrated consistent with the scenario that4. moreLENS MODELING than one galaxy may to the wider SDSS frame with an overall rms uncertainty, rel- be contributing to the emission lines. A further confirmation ′′ We use the publicly available software GRAVLENS (Keeton ative to SDSS, of less than 0.05 .Thisastrometriccalibra- of the redshift of G2 will require additional spectroscopic ob- 2001) to generate the lens model. As the background tion involved more than 60 galaxies and stars. The previous servations or UV imaging data where z ∼ 2 galaxies would source is not multiply-imaged, and remains undetected in the Keck/NIRC2 imaging data, due to the limited field of view be Lyman dropouts. For simplicity, hereafter, we assume that of 40′′ in the highest resolution NIRC2 imaging data used for rest-frame optical, we simply model the highest resolution Figure 2. Spitzer/IRAC 3.6 µm image from Riechers et al. (2013) showingboth G1 andIRAM/PdBI G2 are at 1.1-mm the same continuum redshift of emission 2.1. The map. SED This map LGS/AO observations,a detected source. had The large contours astrometric on the intensity errors scale as astro showm- the regionmodeling of G1 we discuss later is consistent with this assumption. and G2 (blue) and R1 and R2 (red). IRAC 2 PSF is marked with a white is currently our highest resolution view of HFLS3, and the etry was determined based on two bright sources∼ ′′ that were also detectedcircle. in 2MASS. Once the HST frames are calibrated, source3. REST-FRAME is resolved UV in these FLUXES data. OF The HFLS3 magnification factor we determine here with lens modeling is the value for the mm- we fixed the astrometry of Keck/NIRC2 image with close to We use the publicly available software GALFIT (Peng et al. 10 fainter sourcesdetected detected galaxies in bothnear WFC3HFLS3 and and NIRC2 to see images.if there is any ex- wave continuum emission. It could be that HFLS3 will be cess emission in WFC3 data relative to the′′ ACS images.2002) Us- to modelsubject the to differential surface brightness magnification, profiles with of differentHubble- emission This astrometric recalibration resulted in a small (0. 1)shift detected galaxies near HFLS3 and to see if there is any ex- to the opticaling sourcesGALFIT relativeon the to individual the peakHubble PdBI/1.1/ACS mm and emis- WFC3 frames components within the galaxy subject to different magnifica- proved to be difficult because the output models tend to over-cess emissiontion in factors WFC3 (e.g, data Serjeant relative 2012; to the Hezaveh ACS images. et al. Us- 2012). This sion from HFLS3, as can be seen by comparing Figure 1 here ing GALFIT on the individual Hubble/ACS and WFC3 frames with Figurefi 3t of regions Riechers of low et al. signal (2013). in which There HFLS3 is still an is over expected- to re- is especially true if the different components associated with side. To remedy this, we stacked the HST/ACS in two bandsproved to beHFLS, difficult such because as dust, the gas, output and stellar models mass, tend have to over- peak intensiti- all systematic uncertainty in the relative astrometry between fit regions of low signal in which HFLS3 is expected to re- IRAM/PdBIto image increase and theHubble signal-to-noise/Keck images ratio of and about to model0′′. 1, the fore- ties that are offset from each other, as in the case of a complex ground galaxies in the combined′′ F625W and F814W images.side. To remedymerging this, galaxy we stacked system. the HST/ACS in two bands with this value possibly as high as 0. 3 in an extremely un- to increase the signal-to-noise ratio and to model the fore- likely scenario.Under We the account assumption for such that a systematicthe stacked offset model in best the represents To simplify the lens modeling, we use singular isothermal the two foreground galaxies, we then subtracted the stackedground galaxiesellipsoidal in the (SIE) combined models F625W to fit for and the F814W Einstein images. radius and po- lens model by allowing the peak 1.1-mm flux to have an offset model from individual HST/ACS and WFC3′′ frames, withUnder the thesitions assumption of the that two the lens stacked galaxies. model The position best represents angles and ellip- from the two lens galaxies with a value as high as 0. 3. flux density at each wavelength allowed to vary as an over-the two foregroundticities for galaxies, G1 and G2 we are thenfixed subtracted to the values the derived stacked from pro- As shown in Figure 1 (left panel), we detect optical emis- all normalization in GALFIT models. Any excess in WFC3model fromfile individualfitting using HST/ACSGALFIT and, but WFC3 their masses frames, are with allowed the to vary sion from more than one galaxy near HFLS3 (galaxies labeled flux density at each wavelength allowed to vary as an overall z > 6 galaxies can be discovered with just 100 to 600 micron coverage.

Need a survey area of around 1000 deg2 for statistically interesting number of targets.

Dowell et al. 2014 ApJ

Figure 35: Spire color ratios for candidate z > 4 DSFGs in Dowell et al. (2013) (black dots) with confirmed sources shown as red diamonds, including the highest redshift Herschel-selected SMG HFLS3 at z = 6.34. SED tracks based on Arp 220 and the Cosmic Eyelash (Swinbank et al., 2010) are shown for comparison. They gray dashed lines are the selection from Dowell et al. (2013). The figure is a modified version of the same figure published in Dowell[How et al. (2013); angular its reproduction hereresolution is done with permission improvement of the authors and AAS. with CALISTO increases or enhances identification of z > 5 galaxies with far-IR alone?]

“red” galaxies in Herschel

Asantha Cooray, UC Irvine Far-infrared 2016 AAS

79 Galaxy proto-clusters at z >2 (before clusters “virialized” and bright in X-rays and SZ) Galaxy proto-clusters at z >2 Casey et al. 2015: Herschel/SCUBA-2 + redshifts from Keck/ MOSFIRE z=2.47, 8 dusty, starbursting galaxies and 40+ Lyman-break galaxies + radio + AGNs

Far-IR Surveyor over 1000 deg2 will find many 100s of these things - no follow-up as automatic redshifts Intensity Mapping

1. Individual sources are diﬃcult to detect (sources are intrinsically faint, large instrument beam, etc),

2. We are interested in the total power from all sources, or

3. There is truly diﬀuse emission,

Science Applications: • Galaxy Evolution • Dark Matter and Galaxy Formation • Epoch of Reionization • Baryon Acoustic Oscillations.

CMB is the canonical example of IM (Planck Collaboration 2013). CIB Power Spectrum l = 216 l = 21,600 Viero et al. 2012

2-halo

Combined 5 ﬁelds -1] over 70 deg2 sr

2 2-halo 1-halo ) [Jy θ Poisson P(k 1-halo

Halo Model: see e.g., Cooray & Sheth (2000), Zehavi et al. (2005, 2008)

-1 kθ (arcmin ) Smaller Scales Cosmic Infrared Background Fluctuations with SPIRE Viero et al. 2012; Thacker et al. 2013 Asantha Cooray, UC Irvine Far-infrared 2016 AAS 11 Cosmic Infrared Background Fluctuations = Dust Content

Cameron Thacker UCI PhD 2015

Asantha Cooray, UC Irvine Far-infrared 2016 AAS

Fig. 12.— The cosmic density of dust Ωdust vs redshift as determined from the CFIRB power spectra from H-ATLAS GAMA-15 field (shaded region). The thickness of the region corresponds to the 1 − σ ranges of the halo model parameter uncertainties as determined by MCMC fits to the data (Table 1). We also compare our estimate to previous measurements in the literature. The measurements labeled H-ATLAS dust mass function are from the low-redshift dust mass function measurements in Dunne et al. (2011). The other estimates are based on extinction measurements from the SDSS (e.g., Ménard et al. 2010, 2012; Fukugita 2011; Fukugita & Peebles 2004) and 2dF (e.g., Driver et al. 2007).

9 values. At z = 1, as show in Fig. 7, for LIR > 10 Xia et al. (2012). L⊙ galaxies, the HOD drops quickly for masses smaller The MCMC fits to the CFIRB power spectrum data than log(Mmin/M⊙) ≃ 10.7 and the high-mass end has show that the charicteristic mass scale M0l evolves with a power-law behavior with a slope ∼ 1. By design, this redshift as (1 + z)−2.9±0.4. In order to compare this halo model based on CLFs has the advantage that it does with existing models, we convert this evolution in the not lead to unphysical situations with power-law slopes charachteristic mass scale to an evolution of the L(M,z) −αl for the HOD greater than one as found by Amblard et relation. As L(M) ∝ (M/M0l) ,wefindL(M,z) ∝ α −pM αl al. (2011). Ml (1 + z) . Using the best-fit values, we find Both the HOD and the underlying luminosity-mass L(M,z) ∝ M 0.70±0.05(1+z)2.0±0.4. In Lapi et al. (2011), relations are consistent with De Bernardis & Cooray their equation 9 with the SFR as a measure of the IR lu- (2012), where a similar model was used to reinterpret minosity, this relation is expected to be M(1 + z)2.1.In Amblard et al. (2011) anisotropy measurement. The key Dekel et al. (2009), the expectation is M 1.15(1 + z)2.25. difference between the work of De Bernardis & Cooray While we find a lower value for the power-law depen- (2012) and the work here is that we introduce a dust dence on the halo mass with IR luminosity, the redshift SED to model-fit power spectra measurements in the evolution is consistent with both these models. three wavebands of SPIRE, while in earlier work only To test the overall consistency of our model relative to 250 µmmeasurementswereusedforthemodelfit.For existing observations at the bright-end, in Fig. 8, we com- comparison with recent model descriptions of the CFIRB pare the predicted luminosity functions 250 µm-selected power spectrum, we also calculate the effective halo mass galaxies in several redshift bins with existing measure- scale given by ments in the literature from Eales et al. (2010) and Lapi N (M) et al. (2011). The former relies on the spectroscopic M = dMn (M)M T . (21) redshifts in GOODS fields while the latter makes use of eff h n ! g photometric redshifts. We find the overall agreement to be adequate given the uncertainties in the angular power With this definition we find M =3.2×1012 M at z = eff sun spectrum and the resulting parameter uncertainties of 2, consistent with the effective mass scales of Shang et al. the halo model. In future, the overall modeling could (2011) and De Bernardis & Cooray (2012) of M ∼ 4 × eff be improved with a joint-fit to both the angular power 1012 and slightly lower than the value of ∼ 5 × 1012 from 3-D Intensity Mapping

Sky map at z Intensity map at z

• No need to resolve individual source • Measure the collective emission from many sources • Map large volume and faint sources at high z economically • Astrophysical and cosmological applications from structure formation to measurement of SFRD of the universe at z > 2 4CALISTO"SensiGvity" Uzgil, Aguirre, Bradford, & Lidz

Table 1 Li LIR relation variables from Spinoglio et al. (2012) Uzgil et al. 2014

Line i A SAOFIA B B [CII]158µm0.890.032.440.07 [NII]122µm1.010.043.540.11SPIRE [OI]63µm0.980.032.700.10 [OIII]88µm0.980.102.860.30PACS [OIII]52µm0.880.102.540.31 [SiII]35µm1.040.053.150.16 [SIII]33µm0.990.053.210.14IRS [SIII]19µm0.970.063.470.20 [NeII]13µm0.990.063.260.20 [NeIII]16µm1.100.073.720.23 2.5m,"4%"at"6"K,"NEP"2e)19,"4"beams" MIRI SPICA higher redshift, high luminosity systems of our model. While there are undeniably a number of uncertainties Figure 2. Intensity of fine-structure(CCAT,"N=100"MOS)" line emission as a function with the combined Béthermin-Spinoglio model, a simple of observed wavelength for the empirical model based on the B11 luminosity function. Intensities of CO lines, which are not in- extrapolation from the Hopkins & BeacomFar-IR (2006) Surveyor star For a concept Far-IR Surveyor cluded in the3%"at"4.5"K,"NEP"2e)20,"100" IR luminosity relations frombms Spinoglio" et al. (2012), formation history clearly brackets our predicted/CALISTO [CII] in- havebetween been estimated 60 using to 650 a luminosity microns: scaling provided by Carilli tensity at the relevant redshifts, and so we adopt it as (2011) for CO(1-0) and the relative intensities of the higher-J lines our fiducial model throughout this paper. In Section 4, in Bothwell [CII] et al. at (2013). z = 0 and 3 however, we exploreintensity• variations mapping in the shape of the IR lu- [OI] at z = 1 and 7 - minosity function andCALISTO"reaches"0.1"L consider an alternative line-to-IR*"at"z=2.5,""ULIRG"at"z=6."extend to reionization luminosity ratiosignal-to-noise (depicted as the dotted ratios curves in in Fig- 2 • Discovery"potenGal"or"discovery"speed:"" [OIII] …N detetc"x"(A"/"NEP) .""" ure 1). excessJune"3,"2015" of 100 in redshift intervals Next, it becomes possible to write the cosmic meanCALISTO:""Bradford"+"Goldsmith" 9" of 0.3 around z of 2-3 1 intensity and shot noise of the line, in units of Jy sr , as a function of redshift based on the B11 luminosity function and Spinoglio et al. (2012) L L relation as i IR 3D intensity mapping with Far-IR Surveyor f L S¯ (z)= dlogL (L ,z) i IR yD2 (7) i Asantha Cooray,IR UC IrvineIR 4 ⇡ D 2 A,co Far-infrared 2016 AAS Z L f L 2 P shot(z)= dlogL (L ,z) i IR yD2 (8) i,i IR IR 4⇡D2 A,co Z ✓ L ◆ where the limits of integration are over the full range of 8 13 expected IR luminosities, i.e. 10 to 10 L , and fi,i.e. Li(LIR) , is the fraction of IR luminosity emitted in line Figure 3. The fraction of total [CII] mean intensity as a function LIR of lower limit in the luminosity function. Di↵erent color curves i, as computed from equation (3). In other words, we represent di↵erent redshifts, as labeled on the plot. have written S¯ and P shot(z) as the first and the second i i,i largely on [CII] emission. moments of the luminosity function. The resulting mean intensities for a variety of far-IR lines are plotted in Fig- 2.3. [CII] Luminosity Functions and Expected Power ure 2 as a function of redshift and observed wavelength. ¯ Spectra Si vs obs can be interpreted as identifying the dominant clust source of fluctuations, according to our model, of a given As laid out in Equations 1 and 2, P[CII],[CII] is sensi- wavelength. As a specific example, if the target line of tive to intensity fluctuations from the full range of normal 11 12 an observation is [OI]63µm at z = 1, it is necessary to (LIR < 10 L ) to ULIRG-class (LIR > 10 L )sys- distinguish between the target line and interlopers like tems because its amplitude is proportional to the mean [OIII]88µm from z =0.4 and [OIII]52µm from z =1.4, line intensity, squared. The information contained in a which contribute power at the observed wavelength. Vis- power spectrum of individually detected galaxies is, in bal & Loeb (2010) showed how the cross spectra can be contrast to the line intensity mapping approach, neces- used to di↵erentiate between a target line and a contam- sarily limited to galaxies which are above a certain detec- inating line (or “bad line”, in their words), since emitters tion threshold, or LIR,min. Figure 3 shows the integrated at di↵erent redshifts will be spatially uncorrelated. For luminosity functions for [CII] in our model, which gives the observed wavelengths of [CII], however, it is apparent a sense of the depth that a galaxy survey must reach in from Figure 2 that, with the exception of contributions order to completely probe the full integrated [CII] emis- from [OIII]88µm and CO(8-7) near [CII] at z 0.01 and sion, i.e. all of S¯i. In this section, we examine the role z>2, respectively, the [CII] line is relatively⇠ una↵ected of the various luminosity ranges on the amplitude of the by interloper lines—a result of its luminosity and spec- observed [CII] power. tral isolation. It is for this practical reason, and for the Power spectra at four representative redshifts (z = astrophysical significance of [CII] mentioned in the In- 0.63, 0.88, 1.16, and 1.48) comprised of the sources above troduction, that we focus the remainder of this paper afewdi↵erent survey depths, or LIR,min, are repre- Molecular Hydrogen tracing primordial cooling sites/halos

molecular hydrogen cooling in “mini halos” critical at onset of reionization z>15

Outstanding problems at z > 6: billion to ten billion solar mass black-holes in SDSS quasars, Universe at < 600 Myr. One solution is massive PopIII clusters collapsing - seed blackholes. Need formation in minihalos at z > 15. Molecular Hydrogen tracing primordial cooling sites/halos

Gong et al. 2012, ApJ arXiv:1212.2964

Fig. 2.— Left:ThegastemperatureT and H2 fraction fH2 as functions of the gas density n,whicharederivedfromthesimulation results in Omukai (2001) and Yoshida et al. (2006). The uncertainties of the gas temperature and H2 fraction are shown in blue regions. Right: The density profile of the gas (blue solid line) and molecular hydrogen (blue dashed line) for the halo with M =106 M⊙/hat z =15.TheblueregionshowstheuncertaintyoftheH2 density profile estimated by the uncertainty of fH2 in the left panel. luminosity vs. halo mass relation for molecular hydro- the results on T (n)andfH2 (n)fromOmukai(2001)and gen cooling within primordial dark matter halos, we first Yoshida et al. (2006), which are available for n ≃ 10−2 to need to know the radial profile of the gas temperature 1023 cm−3 as shown in the left panel of Fig. 2. The uncer- and density within dark matter halos. tainties of the gas temperature and H2 fraction are shown Following the results from numerical simulations in- in blue regions. These uncertainties are evaluated based volving the formation of primordial molecular clouds on the differences in the far-ultraviolet radiation back- (e.g. Omukai & Nishi 1998; Abel et al. 2000; Omukai ground from the first stars and quasars (Omukai 2001). 2001; Yoshida et al. 2006; McGreer & Bryan 2008), we As can be seen, the gas temperature does not assume the gas density profile as monotonously increase with the gas density. For in- −2.2 stance, it drops from T ∼ 2000 to 200 K between n ≃ 1 r 4 −3 ρ(r)=ρ0 , (5) and 10 cm where molecular hydrogen density is ris- r ∼ −3 ! 0 " ing to fH2 10 . This indicates that H2 cooling is starting to become important in this gas density range. where we set r0 =1pcandρ0 is the normalization factor ≃ 4 −3 which is obtained by At n 10 cm ,H2 cooling saturates and turns into the cooling at the LTE. For n =1010 ∼ 1011 cm−3, al- rvir 2 most all of gas particles become molecular hydrogen due Mgas =4π r ρ(r)dr . (6) to the efficient H three-body reaction (Yoshida et al. 0 2 # ≃ ≃ 20 2006), and we find fH2 0.5 by definition. At n 10 Here Mgas =(Ωb/ΩM )M is the gas mass in the virial −3 4 cm with T ≃ 10 K, H2 begins to dissociate and the radius of the halo with dark matter mass M,andthe −5 23 fraction drops quickly to f 2 < 10 when n ≃ 10 r is the virial radius which is given by H vir cm−3 and T ≃ 105 K. M 1/3 Next, with the help of Eq. (5), we can evaluate the rvir = . (7) gas temperature and H2 fraction as a function of the (4/3)πρvir halo radius, i.e. T (r)andfH2 (r). Once these are estab- $ % r,v Here ρvir(z)=∆c(z)ρcr(z) is the virial density, ρcr(z)= lished, we can derive nH2 (r), nH(r)andΛH2 (r) which are 3H2(z)/(8πG) is the critical density at z, H(z) is the needed for the H2 luminosity calculation. In the right 2 2 Hubble parameter, and ∆c(z)=18π +82x−39x where panel of Fig. 2, we show the density profile of the gas x = Ω (z) − 1. and molecular hydrogen in blue solid and dashed lines M 6 We then derive the number density of gas by n(r)= for a dark matter halo with M =10 M⊙/hatz =15. nH(r)+nHe(r). Here nH(r)=fHρ(r)/mH is the num- The blue region shows the uncertainty of the H2 density ber density of hydrogen, where fH =0.739 is the hy- profile which is derived by the uncertainty of fH2 in the drogen mass fraction and mH is the mass of hydrogen left panel of Fig. 2. The gas density profile is a straight atom. Similarly, nHe(r)=(1− fH)ρ(r)/mHe is the num- line with a slope of -2.2 as indicated by Eq. (5). On the ber density of helium, where mHe is the mass of helium other hand, the density profile of molecular hydrogen has atom. Also, the temperature-density relation T (n)and a more complex shape which is dependent on the relation between f 2 and gas density n. For the outer layer the H fraction-density relation f 2 (n)=n 2 /n can be H 2 H H − derived from existing numerical simulations. Here we use of the gas cloud (r>10 2 pc), gas density is less than Molecular Hydrogen tracing primordial cooling sites/halos

Gong et al. 2012, ApJ arXiv:1212.2964 H2 0-0 (S3) 6 6 SPICA-like CALISTO-like deep field CALISTO-like lensed counts

Fig. 6.— The auto power spectrum of the 0-0S(3) line (left panel) and the cross power spectrum of the 0-0S(3)×0-0S(5) (right panel) Fig. 6.— The auto power spectrum of the 0-0S(3) line (left panel) and the cross power spectrum of the 0-0S(3)×0-0S(5) (right panel) at z =12withtheerrorsestimatedforaSPICA/BLISS-likeand10× better SPICA/BLISS-like surveys in each plot. The noise power × at z =12withtheerrorsestimatedforaSPICA/BLISS-likeand10spectrum and shot-noise power spectrum are shownbetter in long-da SPICA/BLISS-likeshed and dotted surveys lines, in respectively. each plot. The noise power spectrum and shot-noise power spectrum are shown in long-dashed and dotted lines, respectively. 0S(5), 0-0S(4), 0-0S(1)0S(5), and 0-0S(4),0-0S(2) are 0-0S(1) also and bright 0-0S(2) with are alsoand bright with and total intensities of ∼ 4.3, 2.0, 1.5 and 1.3 Jy/sr, respec- ∞ i j ∼ j total intensities of 4.3, 2.0, 1.5 and 1.3 Jy/sr, respec- ∞ ij i dn LH2 LH2 i 2 j 2 tively, at z = 15. The vibrational lines 1-0S(1),ij 1-0Q(1) dnP L =2 L 2 dM y (z)D y (z)D , tively, at z = 15. The vibrational lines 1-0S(1), 1-0Q(1) P = dM shot H H yi(z)D2 yj(z2 )D2 ,2 A A and 1-0O(3) have low mean intensities atshot the level of 2 Mmin2 dM A4πDL 4πADL and 1-0O(3) have low mean intensities at the level of min dM 4πD !4πD 0.83, 1.0 and 0.99 Jy/sr, respectively. The mean! biasM L L (14) 0.83, 1.0 and 0.99 Jy/sr,factor respectively. of these lines lies The between mean bias 2.6 (for 0-0S(0)) and 4.2 respectively. From these equations,(14) we find that the the factor of these lines lies(for between 0-0S(11)), 2.6 and(for the0-0S(0)) bias factors and 4.2 of therespectively. rotational lines From thesecross equations, power spectrum we find should that the have the a similar magnitude (for 0-0S(11)), and theat bias higher factors rotational of the energyrotational level lines are highercross than power that spectrumat to the should auto have power a spectrum.similar magnitude The clustering cross power at higher rotational energy level are higher than that at to the auto power spectrum.clus The clustering cross power lower level. This is because the lines with high Jclusare spectra Pij for several H2 lines at z =15areshown lower level. This is becausestronger the at higher lines with mass high halosJ whereare thespectra temperaturePij isfor severalin the H right2 lines panel at z of=15areshown Fig. 5. We choose the strongest stronger at higher masslarger. halos where the temperature is in the right panel of0-0S(3) Fig. 5. line We to choosecross correlate the strongest with the other 5 bright larger. In the left panel of Fig. 5, the clustering0-0S(3) auto line power to crosslines, correlate i.e. 0-0S(1), with the 0-0S(2), other 0-0S(4), 5 bright 0-0S(5) and 1-0Q(1). In the left panel ofspectrum Fig. 5, the of eight clustering H2 lines auto at z power=15areshown.Wefindlines, i.e. 0-0S(1), 0-0S(2),We find 0-0S(4), that the 0-0S(5) 0-0S(3) and× 0-0S(5) 1-0Q(1). is the largest cross × spectrum of eight H2 thatlines the at z shot=15areshown.Wefind noise power spectrum PshotWeis find relatively that the 0-0S(3)power spectrum0-0S(5) since is the they largest are crossbrightest two lines. At that the shot noise powersmall compared spectrum toPshot the clusteringis relatively power spectrumpower spectrumPclus, sincez ∼ they15, are then brightest we would two be lines. cross-correlating At the wave- small compared to theand clustering would not power affect spectrum the PclusPatclus the, scalesz ∼ of15, interest. then we wouldlength be regimes cross-correlating around 110 the and wave- 155 µm. A search for and would not affectThis the P isclus easyat to the understand scales of if interest. we notice thatlength the halo regimes mass aroundmid-IR 110 lines and revealed155 µm. no A astrophysical search for confusions from This is easy to understandfunction if we is notice dominated that bythe halos halo with mass low massesmid-IR which lines are revealedlow no redshifts astrophysical that overlap confusions in these from two wavelengths at function is dominatedmore by halos abundant. with low masses which are low redshifts that overlapthe same in redshift. these two Thus, wavelengths while low at redshift lines will more abundant. We also calculate the cross correlation betweenthe same two redshift. dif- easily Thus, dominate while low the redshift auto power lines spectra will of H2 lines, the We also calculate theferent cross H correlation2 lines. Such between a cross-correlation two dif- easily will reduce dominate the thecross auto power power spectrum spectra of will H2 lines, be independent the of the low- ferent H2 lines. Suchastrophysical a cross-correlation contamination will reduce from the the othercross sources, power such spectrumredshift will beconfusions. independent In addition of the low- to reducing the astro- astrophysical contaminationas low-redshift from the emission other sources, lines from such star-formingredshift galaxies, confusions.physical In addition confusions, to reducing the cross the power astro- spectra also have the as low-redshift emissionincluding lines from 63 µ star-formingm [OI] and 122galaxies,µm [NII],physical among others. confusions,advantage the cross power that it spectra can minimize also have instrumental the systematics ∼ including 63 µm [OI]At andz 12215,µ them [NII],dominant among rotational others. line 0-0S(3)advantage would that be it canand minimize noise, depending instrumental on systematicsthe exact design of an experi- At z ∼ 15, the dominantobserved rotational at a wavelength line 0-0S(3) of would 155 µm. be Suchand a line noise, would depending be ment. on the exact design of an experi- contaminated by, for example, z ∼ 0.3 galaxies emitting observed at a wavelength of 155 µm. Such a line would be ment. 5. DETECTABILITY contaminated by, for[NII]. example, Thusz ∼ the0.3 auto galaxies power emitting spectrum would be higher [NII]. Thus the autothan power what spectrum we have would predicted be higher given that the line intensi- 5. DETECTABILITYIn this section we investigate the possibility to detect than what we have predictedties of [NII] given are that higher the thanline intensi- H2 lines. To avoidIn this this section as- wethese investigate lines based the possibility on current to or detect future instruments. We trophysical line confusion we propose a cross-correlation assume a SPICA-like 1 survey with 3.5 m aperture diam- ties of [NII] are higher than H2 lines. To avoid this as- these lines based on current or2 future instruments. We trophysical line confusionbetween we propose two rotational a cross-correlation or rotational and vibrationalassume a SPICA-like lines eter,1 survey 0.1 deg withsurvey 3.5 m area, aperture 10 GHz diam- band width, R=700 between two rotationalof or the rotational H2 line emission and vibrational spectrum. lines eter, 0.1 deg2 surveyfrequency area, 10 resolution, GHz band 100 width, spectrometers,R=700 and 250 hours The cross clustering and shot-noise powerfrequency spectrum resolution, for total 100 integration√ spectrometers, time and and 250 noise hours per detector σpix = of the H2 line emission spectrum. 6 such two H2 lines i and j can be evaluated as 10 Jy s/sr at 100 µm. Such an instrument corresponds The cross clustering and shot-noise power spectrum for total integration√ time and noise per detector σpix = 6 to the latest design of the mid-IR spectrometer, BLISS, such two H2 lines i and j can be evaluated as 10 Jy s/sr at 100 µm. Such an instrument corresponds to the latest designfrom of the SPICA mid-IR (Bradford spectrometer, et al. 2010). BLISS, ij ¯i ¯j ¯i ¯j from SPICA (Bradford1 http://sci.esa.int/science-e/www/area/index.cfm?far et al. 2010). eaid=105 Pclus = IH2 IH2 bH2 bH2 Pδδ (13) ij ¯i ¯j ¯i ¯j 1 http://sci.esa.int/science-e/www/area/index.cfm?fareaid=105 Pclus = IH2 IH2 bH2 bH2 Pδδ (13) Architecture Consensus

70% 30% Source reconstructed: SMA/JVLA (red) and Keck/Hubble (green/blue)

CALISTO SPIRIT Some (old) new ideas appeared, though: • Polarimetry Interferometer should be able to resolve • Very high spectralindividual resolution starbursting clumps out to z of 2. 15 December 2016 3rd FISICA Workshop: London, UK 7 Resolve narrow-line regions of local AGNs extended out to 500 pc in [OIV]/26 & [NeV]/24

Interferometer can also separate AGN from starburst components.

Sensitivity needed to lines at the level of 1e-19 W/m2.

Far-IR Interferometer Science Case

Asantha Cooray, UC Irvine far-infrared 2016 AAS Wish list single aperture: Wedding cake survey from deep 1-3 deg2, medium tier of 100-300 deg2, and shallow wide 1000-2000 deg2, 60-600 microns R~300-600, 12 arcsec spatial resolution at 250 um

interferometer: line sensitivity below 1e-19 W/m2, probe ~100 range of AGN and starburst galaxies

New interesting sciences:

(a) Molecular Hydrogen pre-reionization at z ~15 (especially in a deep survey of lensing galaxy clusters for example).

(b)OI at z > 6 to combine with mm-wave CII etc from ground and 21-cm experiments such as SKA low-frequency

Summary

Asantha Cooray, UC Irvine Far-infrared 2016 AAS