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Change Log

Revision Author(s) Revision Date Description J. Glenn, F.Bertoldi, J. Carpenter, D. Lis, S. Padin, T. Phillips, R. Plume, P. Schilke, G. Stacey

Sept. 27, 2011 – v6 First version of document (except instrument requirements table)

J. Glenn, F.Bertoldi, J. Carpenter, S. Golwala, D. Lis, S. Padin, T. Phillips, R. Plume, P. Schilke, G. Stacey

Dec. 6, 2011 – v8 First complete version of document, with instrument requirements and science committee comments addressed

J. Glenn, D. Lis, R. Plume, P. Schilke

January 20, 2012 – v10 Heterodyne array requirements updated; this is the release version

   

  

Table of Contents

1. Introduction 1

2. Description of CCAT 2

3. CCAT Science Drivers and Requirements 3

4. Detailed Telescope Requirements 34

5. Detailed Instrument Requirements 39

6. Appendix A: Acronyms 40

7. Appendix B: Document History 42

 

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1.0 INTRODUCTION

In recent decades, major advances have been made in our understanding of the evolution of galaxies, the formation of stars and planetary systems, and the nature and evolution of the Universe itself. We now know that galaxies grow at least partly through mergers, that many stars have planetary systems, and the baryonic, dark matter, and dark energy content of the Universe have been precisely measured. However, important questions remain. When did the first galaxies form? What is the history of star formation (and supermassive black hole formation) in galaxies and what is the origin of interstellar dust? How does the astrophysics of feedback in galaxy clusters affect their utility for precision cosmology? What determines the masses of stars when they form?

Submillimeter astronomy is poised to directly address these questions, and others. Submillimeter observations provide the only opportunity to measure the bolometric luminosities of star-forming galaxies from z ~ 1 to the epoch of reionization – and beyond – by measuring the emission from the very dust that obscures young and forming stars. Observations of submillimeter atomic and molecular transitions enable astronomers to model the detailed physics of the interstellar media in galaxies from the nearby universe to z > 5, and in clouds in the Milky Way. High spatial resolution millimeter-wave observations of the Sunyaev-Zel’dovich effect in galaxy clusters, coupled with submillimeter observations to remove contamination from cluster members and lensed galaxies, will enable precise measurements of cluster scaling relations utilized for cluster mass estimates and thereby provide strong constraints on non-gaussianity in the dark energy equation of state. Deep submillimeter surveys will measure the mass function of protostellar cloud cores down to substellar (brown) dwarf masses within two kpc of the Sun.

The clear need for a submillimeter telescope with excellent sensitivity, high mapping speed, and good angular resolution to address major questions in astrophysics was recognized in the strong recommendation of CCAT by the Astro2010 Decadal Survey1. CCAT, a 25 meter diameter, large field-of-view (FoV) telescope to be constructed atop Cerro Chajnantor in Chile, will meet the needs identified by Astro2010 and be complementary to ALMA2.

CCAT will be a versatile observatory with unique and transformational capabilities. The need for CCAT is driven by several themes: high-z galaxy formation and evolution, characterization of the atomic and molecular interstellar medium in nearby galaxies, censuses of star formation regions in the Milky Way, and the physical state of baryonic matter in galaxy clusters.

This document addresses the science requirements for CCAT and the telescope and instrument requirements that follow from them. The definition of surveys to meet the science goals of the telescope will be described in future documents.

                                                            1 Astro2010:  The Astronomy and Astrophysics Decadal Survey; New Worlds, New Horizons in Astronomy and Astrophysics (http://sites.nationalacademies.org/bpa/BPA_049810). 2 https://almascience.nrao.edu/about‐alma.   

 

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2.0 DESCRIPTION OF CCAT CCAT will be a 25 meter diameter Ritchey-Chretien (RC) telescope operating in the 0.35–2.1

mm wavelength range (with a goal of operating in the 200 m – 3.5 mm range). CCAT will be

large enough to resolve the submillimeter background at λ = 350 μm, and its small beam (3.5 FWHM at λ = 350 μm with pointing error < 1/10th beamwidth) will yield source positions with subarcsecond accuracy to support follow-up observations. CCAT will have high efficiency (half wavefront error (HWFE) < 12.5 μm RMS for Strehl ratio > 0.8 at λ = 350 μm), and it will be located on an excellent site, at an altitude of 5,600 m, on Cerro Chajnantor in northern Chile

(long. 67 44 24 west, lat. 22 59 08 south). The cumulative precipitable water vapor above

Cerro Chajnantor is 200 m or less 10% of the time, 400 m or less 25% of the time, and 600

m or less 50% of the time, appreciably better than the ALMA plateau3. CCAT will achieve about the same continuum per-pixel sensitivity as ALMA at λ = 450 μm, and about twice the

sensitivity of ALMA at = 350 µm (owing to better telescope surface accuracy and a higher site). The RC design will support cameras and spectrometers with up to 1° FoV. The combination of an efficient telescope, an excellent site, and wide-field cameras and spectrometers will make CCAT a uniquely powerful survey instrument.

CCAT will have an active primary mirror to compensate for gravitational and thermal deformations. The primary mirror will be supported by a carbon-fiber-reinforced plastic (CFRP) spaceframe truss on an elevation-over-azimuth mount made of steel. The primary surface will be made of 162 keystone-shaped segments, each with 16 ~ ½ × ½ m machined aluminum tiles mounted on a ~ 2 × 2 m CFRP subframe. The surface control will be open loop, based on look-up tables for elevation and soak temperature. The primary mirror will be fast, f/0.4, to minimize the size and cost of the enclosure. The secondary mirror will be mounted on a hexapod positioner at the apex of a CFRP tripod support attached to the truss. The tertiary mirror will be mounted on a rotator at the intersection of the azimuth and elevation axes, just behind the primary.

Instruments will be located at the two f/6 Nasmyth foci, which will be inside the elevation axle near the tertiary. The active Nasmyth focus will be selected by rotating the tertiary. The Nasmyth focal surface will be 2.6 m in diameter for a 1° FoV, so it will be possible to support a few large instruments, or multiple smaller instruments. CCAT will likely operate with all instruments on line all the time, with a simple pointing change to switch between instruments. Table 1 shows a possible 1st light instrument complement. CCAT will be inside an enclosure to reduce wavefront and pointing errors due to wind forces and thermal deformation due to solar illumination. The enclosure will also protect the telescope from storms.

                                                            3Radford, S. J. E., et al. Proc. of 18th Int. Symp. On Space THz Tech., arXiv:0704.3031, (2007)   

 

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Table 1−Possible 1st light instruments

Instrument  Banda  Detectors FoV Notes 

     SWCam  (200), 350, 450, 

(620) m 

3 sub‐field cameras ×1282 

6.5’ at λ=350μm 3 ATACamera modules,½Fλ absorber‐coupled pixels 

LWCam  (750), 850, 1100, 

1300, 2100 m 

542 15’ at λ=850μm 5×MUSIC, (1−2)Fλ antenna‐coupled pixels 

D‐Spec  350−1300 m  TBD TBD Direct‐detection “redshift” spectrometer 

H‐Spec  460, 660, & 809 GHz 

> 100 channels, > 50 pixels (dual polarization) 

TBD Heterodyne array spectrometer 

aBands in parentheses might not be included in a 1st light camera. 

3.0 CCAT SCIENCE DRIVERS AND REQUIREMENTS

3.1 Measuring the History of Star Formation in Galaxies from the Epoch of Reionization, Through the Peak of Activity, to Today

3.1.1 Galaxy Formation and Evolution Background

A major goal of contemporary astrophysics is to develop a comprehensive model of the formation and evolution of galaxies from the earliest epoch of galaxy formation to the present day. The submillimeter wavelength regime will play a crucial role in realizing this goal by measuring the star formation history of galaxies. In the 1990s, the Cosmic Background Explorer (COBE) revealed an (unresolved) far-infrared extragalactic background radiation (CFIRB) with integrated flux approximately equal to all the visible and ultraviolet extragalactic starlight reaching Earth4,5, which indicated that averaged over the entire Universe, dust reprocessing accounted for 50% of the energetics from nucleosynthesis and accretion. This forced a dramatic reappraisal of astronomers’ picture of galaxy formation and evolution: models must account for bolometric galaxy luminosities and star formation rates, as a function of redshift.

“Submillimeter” galaxies, those that have been discovered by SCUBA, MAMBO, AzTEC, LABOCA, Bolocam, BLAST, and Herschel Space Observatory at submillimeter and millimeter wavelengths, are the extreme-luminosity (~1013 Lsun) tip of the dust obscured galaxy population. Tens of thousands of submillimeter galaxies have been discovered, but they only account for ~10% of the CFIRB, and all studies to date are severely source confusion limited. Not only has source blending been a major problem for these surveys, but robust counterpart identification at

                                                            4 Puget, J. ‐L., et al., A&A, 308, 5 (1996) 5 Fixen, et al., ApJ, 473, 576 (1996) 

 

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other wavelengths has only been possible for a fraction of the catalogued galaxies. Much deeper

(from 350m = 0.1 mJy, the confusion level, to 0.5 mJy), large-area (a few square degrees integrated down to the confusion level to hundreds of square degrees for shallower surveys), multicolor submillimeter-millimeter surveys, with few-arcsecond-scale angular resolution, will be required to characterize the galaxy population and establish the role of submillimeter galaxies in galaxy formation and evolution scenarios. With large surveys to draw from, ALMA can be used optimally to study galaxies in detail.

We will now take “submillimeter galaxies” to mean galaxies that emit an appreciable fraction of their luminosity at submillimeter wavelengths, indicating embedded star formation and/or nuclear accretion by supermassive black holes. Thus, submillimeter galaxies will henceforth comprise all galaxies that contribute to the CFIRB. Basic questions must be answered about submillimeter galaxies and their role in the evolution of galaxies: What fraction of the energy produced by submillimeter galaxies is generated by mergers? What fraction of submillimeter galaxies are blends? Have most galaxies undergone short-lived, extreme spheroid-forming starbursts, or is the galaxy assembly process more quiescent and long term? What is the star formation history (SFR as a function of redshift) of the Universe? Does it match the redshift distribution of QSOs, implying that central supermassive black hole formation and the build-up of stellar populations are coeval, or does it perhaps precede it indicating that QSOs quench star formation? What fraction of the bolometric energy radiated by submillimeter galaxies is produced by accretion onto supermassive black holes in active galactic nuclei (AGN) as opposed to star formation?

3.1.2 Galaxy Formation and Evolution Goals for CCAT

There are several specific scientific goals that must be met by CCAT to address these basic questions about galaxy evolution.

Measure the bolometric luminosities and SFRs of star-forming galaxies as a function of redshift. Only far-infrared (FIR) and submillimeter observations can reveal the bolometric luminosities of star-forming galaxies because (1) interstellar dust obscures the ultraviolet (UV) and optical light of young and forming stars and (2) at the characteristic temperatures of interstellar dust (10 to 100 K) the spectral energy distributions (SEDs) peak in FIR and submillimeter in the observed frame, or possibly millimeter-wave for the highest-z galaxies (z > 5). Sizable samples of galaxies will be required to characterize the population across redshift and free of sample variance6.

                                                            6As an example, using the models of Bethermin, M., et al., A&A, 529, 4 (2011), an assumed  = 350 m NEFD of 14 

mJy (see Section 3.1.3), a 6.5  6.5 FOV, and 25% observing time for precipitable water vapor (PWV)  400 m (Radford, S. J. E., et al. Proc. of 18th Int. Symp. On Space THz Tech., arXiv:0704.3031, (2007)), yielding 2,192 hours 

of observing per year, 56,000 galaxies would be detected in one square degree at or above 5 in 470 hours, where  = 0.1 mJy is the confusion noise. In 1,720 hours remaining in one year for PWV  400 m, 90 square degrees 

could be observed to  = 0.5 mJy, yielding 1.8  106 galaxies at or above 5.  

 

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Place submillimeter galaxies into the context of galaxy evolution models. The properties of galaxies detected at other wavelengths (X-ray, UV, optical, near-infrared, radio) are well understood compared to galaxies dominated by submillimeter emission. By characterizing the submillimeter galaxy population and correlating it with other populations, the interstellar medium (ISM), SFRs, stellar populations, and nuclear activity can be combined into comprehensive models of galaxy evolution. One of the largest limitations of current submillimeter data sets is the lack of a large number of redshifts, which make it very difficult to infer physical properties of the galaxies. CCAT must be able to obtain redshifts for as many of the galaxies it discovers in surveys as practical with submillimeter/millimeter spectrometer technology.

Establish what types of galaxies comprise the CFIRB. That is, what are their SFRs, how large are their gas reservoirs, are they mergers? Separate the energy generated by star formation and AGN (using redshifted far-infrared emission lines and molecular rotational lines to infer the gas excitation mechanisms).

Measure the clustering properties of submillimeter galaxies from the smallest angular scales that can be proved (i.e., of tens of arcseconds) to many degrees. Comparison of

clustering two-point correlation functions and P(k) to Cold Dark Matter (CDM) structure formation models will establish the halo masses and occupation numbers of submillimeter galaxies, and therefore their formation environments, and by extension their likely modern-day environments (e.g., field, groups, or clusters). Small-scale clustering (arcminute-scale) can reveal protoclusters.

3.1.3 Galaxy Formation and Evolution Requirements

To attain the galaxy evolution goals, the following requirements must be satisfied:

Wavelength Range: 350 µm – 2.1 mm, with a goal of 200 µm – 2.1 mm. Dust emission from galaxies typically has a modified blackbody spectrum with observed-frame flux density peaks ranging from short of 100 µm for warm, nearby galaxies, to a few hundred µm or longer for cooler and higher-redshift galaxies. Thus, measuring the bolometric luminosities of galaxies for 0 < z < 10 requires coverage from the shortest wavelengths allowable by the atmospheric transmission to 2.1 mm. Coverage from 350 µm to 1.3 mm is necessary to identify high-z (z > 4) galaxy candidates by color selection and for [CII] (158 µm), [NII] (122 µm), [OIII] (88 µm), and CO spectroscopy for accurate redshifts and gas diagnostics7.

Diffraction-Limited Beam Size and Aperture: A FWHM beam size of 3.5 at 350 m is required for four reasons: (1) Based on the most up-to-date number counts measurements

and models8,9, 3.5 FWHM beams are required to resolve the majority (~80%) of the

                                                            7 The 200 µm to 600 µm range also allows for the possibility of observing the earliest galaxies in redshifted (z = 12 – 20)  28 µm and 17 µm H2 lines, which are believed to be the primary coolants of low‐metallicity (proto)galaxies. 8 Glenn, J., et al., MNRAS, 732, 35 (2010) 

 

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CFIRB10 at = 350 µm (where it contains much more energy than = 1.3 mm). (2) To resolve mergers of Milky Way-sized galaxies with radii of ~15 kpc requires 30 kpc FWHM beams. For 1.5 < z < 2, approximately the peak epoch of galaxy formation, 30

kpc corresponds to 3.5. For z < 1.5 and z > 2, 3.5 resolves structures smaller in size than 30 kpc, although it is a slow function of redshift for 2 < z < 10. (3) This angular resolution is also required to identify counterparts in crowded fields at shorter-wavelengths to measure stellar content, morphologies, and redshifts of galaxies. (4) To detect ultraluminous-infrared-galaxy (ULIRG) luminosity galaxies (LIR = 1012 Lsun) at z =

6 at five times the expected 1 confusion noise11 (5conf), requires a FWHM beam size of

3.5. Thus, with 3.5 resolution CCAT will be able to probe the epoch of reionization (ending at approximately z = 7) for super-ULIRG galaxies. A diffraction-limited 25 m

aperture yields 3.5 FWHM beams at 350 m.

Sensitivity and Aperture: (1) To enable strong synergy between CCAT and ALMA in a scenario in which CCAT performs surveys to identify objects to be followed-up at high angular resolution with ALMA, CCAT must detect galaxies in the continuum much faster than ALMA. This will be possible if CCAT meets or exceeds the per-beam sensitivity of ALMA in CCAT’s

prime submillimeter bands closest to the peak in the CFIRB (350 m and 450 m) and if CCAT has large focal plane arrays of detectors. For CCAT to match the per-beam continuum sensitivity of ALMA at = 450 µm, a noise-equivalent flux density (NEFD) of 14 mJy s1/2 is needed12, requiring an aperture diameter of 25 meters. This corresponds

to 1.0 mJy/beam RMS at 450 µm, 5, in a one-hour integration. The 350 m NEFD for a 25 m aperture is also 14 mJy s1/2, exceeding ALMA’s predicted sensitivity. The sizes of the focal plane arrays (in terms of numbers of pixels) should be as large as possible when collectively considering the detector array technology, the budget available for instrumentation, and the instrument suite required to meet CCAT’s science goals. (2) CCAT must be able to measure redshifts for as many galaxies discovered in its continuum surveys as possible using atomic fine-structure lines and molecular rotational lines. This requires two things: excellent sensitivity and multiobject spectroscopic capability. Excellent sensitivity is required because the number of detectable galaxies depends sensitively on integration depth. Using the Bethermin, et al.13, models and

                                                                                                                                                                                                9 Bethermin, M., et al., A&A, 529, 4 (2011) 10 Galaxy number counts based on the deepest data available (Herschel Space Observatory) only reach down to 2 

mJy; current models cannot be used to accurately or precisely predict what fraction of the 350 m background radiation will be resolved by CCAT, but it will likely constitute the bulk of the background.  11Within the context of current galaxy number count models, detecting sub‐ULIRG luminosity z = 6 galaxies at 

5conf is not practical for a 25 m class, single‐dish, submillimeter telescope.  12For assumptions, see Section 4.2.4 of the Feasibility/Concept Design Study at http://www.submm.org/doc/2006‐01‐ccat‐feasibility.pdf.  13 Bethermin, M., et al., A&A, 529, 4 (2011) 

 

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assuming CO line luminosities and LCO/LFIR from M8214, L[CI]158m = 10-3LFIR, 0.4 mm

PWV, a 25 m aperture, 10 hour integration times (with secondary chopping), 3 line

detections15, and 1 sensitivities of 1x10-20 W m-2, 6x10-21 W m-2, 1.0x10-21 W m-2, and

6x10-22 W m-2 for the 350 m, 450 m, 850 m, and 1.1 mm atmospheric windows, respectively, approximate detection rates of galaxies would be approximately one

thousand per square degree in [CI] 158 m at redshifts of z ~ 1.2, 1.7, 4.3, and 6, one

hundred per square degree in CO in the 350 m and 450 m atmospheric windows (at z <

1), and one thousand per square degree for z 1 in the 850 m and 1.1 mm windows. Reducing the diameter of CCAT by 20% would reduce the detection rates by a factor of ~ 0.6 – 0.8 for [CI] and factors from 0.1 to 0.8 for CO, depending on redshift and J transition, with a factor of ~ 0.5 typical. The number of multi-object spectrometer (MOS) beams coupling to galaxies should be maximized16.

Mapping Speed and FoV: (1) To enable large-scale surveys, (2) to take advantage of megapixel-scale submillimeter cameras that are anticipated within the next 1 – 2 decades, and (3) to enable MOS of as many objects as possible simultaneously, CCAT must have the largest FoV possible. Practically, the FoV should not be limited by mechanical interference.

Emissivity: So that the telescope emission is less than the atmospheric emission, the combined emissivity of the telescope, comprised of the primary mirror, secondary mirror, secondary mirror support legs, and tertiary mirror, shall be less than 10% at = 350 µm.

Sensitivity and HWFE: The HWFE of the system comprised of the primary mirror, secondary mirror, and tertiary mirror shall be less than 12.5 µm RMS, such that integration times will be < 1.5 times that which would be achieved in the absence of any wavefront error.

Az/El Range: The Az range shall be 360 plus 90 in either direction, totaling 540, centered on Az = 90 to prevent cable wrap when tracking sources through zenith. The

El range shall be as large as possible17, at least 20 < El < 90. Offset Pointing: Photometry and spectrophotometry of point sources within 1 of a

pointing source should not be affected by more than 5% pointing errors. For a Gaussian approximation to the telescope beam and a point source, offsets of 0.1 FWHM, 0.15 FWHM, and 0.2 FWHM yield amplitudes of 0.973, 0.940, and 0.895, respectively. So, the offset pointing within 1 of the most recent pointing measurement should be 0.35 ( / 350 m) RMS.

                                                            14 Panuzzo, P., et al., A&A, 518, 37 (2010) 15 3 sigma will only be sufficient for line detections if multiple lines are detected; for a single line greater significance will be required. 16 A trade study is underway to compare steered‐beam MOS options and fully filled focal plane spectrometer technologies considering the likely sky density of detectable galaxies and technology maturities. 17 The elevation is constrained to  20 by mechanical interference in the mount and truss designs. 

 

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Point Knowledge During Scan and Pointing Reconstruction: So that photometry of point sources in maps is not affected by more than 10% by pointing errors, the pointing knowledge during the scan and the pointing reconstruction should be less than 0.2

FWHM , or 0.70 ( / 350 m) RMS.

Pointing Stability: To enable post-processing pointing of 0.70 ( / 350 m) RMS and to enable integration times of up to one hour for high observing efficiency, the

pointing stability must be 0.1 FWHM, or 0.35 ( / 350 m) RMS.

Scanning Speed: The telescope must scan fast enough to modulate the astrophysical signals at a rate greater than the expected rate of atmospheric fluctuations18 to minimize susceptibility to atmospheric 1/f noise. This is satisfied by a wavelength-dependent

speed given by (0.33 s-1) x ( / 350 µm) in Az, and half that in El (for scanning at El up

to 60, Az scan rate 2 El scan rate).

Beam Modulation for Spectroscopy, Secondary Nutation, Nodding, and MOS: Either: (1) multiobject capability must be enabled by the spectroscopic instrumentation for time-stream differencing to remove atmospheric signal, (2) beam modulation – switching – must be provided by the spectroscopic instrumentation to remove atmospheric signal, or (3) secondary mirror nutation and telescope nodding must be provided by the telescope for this purpose. Options (1) and (2) are preferred to minimize cost and complexity of the telescope and one or both are likely realizable. For option (1), the MOS beams can be differenced to remove background. In the case of option (2), background will be subtracted using off positions. Should options (1) and (2) not be possible and option (3) selected, nodding of 10 beams, or 3.5' at = 2.1 mm, at 0.1 Hz, with a duty cycle (settling to within 0.1 beams) of > 80% is required. Secondary nutation of 7' peak-to-peak will be required at 1 Hz with a duty cycle (settling to within 0.1 beams) of > 80%.

Scan Patterns: To enable both large and small maps and to maximize observing efficiency (e.g., minimize lost turn-around time), raster scans with scans and cross-scans at arbitrary angles (e.g., Az/El, RA/Dec, arbitrary position angles to match other observations), Lissajous, and Box-scans are required.

Calibration: A continuum flux density calibration goal of 10% error is likely achievable

(for the 350 m and longer wavelength bands) and is necessary to minimally impact galaxy SED measurements.

                                                            18 See Section 6.6 of the Feasibilty/Concept Design Study, http://www.submm.org/doc/2006‐01‐ccat‐feasibility.pdf. 

 

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3.2 Characterizing the Atomic, Molecular, and Solid-State Interstellar Media of Nearby Galaxies

3.2.1 Nearby Galaxies Background

To understand the processes that lead to evolution of star formation activity over cosmic time, it is important to study nearby, spatially resolved systems so that one can relate the astrophysical probes and lessons learned to the high redshift, unresolved systems. Most of the star-forming galaxies in the early universe are dusty, and often extremely dusty, so that it is often impossible to directly observe the star-formation regions themselves using rest-frame optical or UV probes. The overwhelming presence of dust also ensures that much, and often most, of the energy of these sources – be it released through nuclear or gravitational process – emerges in the rest-frame FIR bands.

Stars form in the cores of molecular clouds. In the local universe the rate of star formation in galaxies like the Milky Way can be described -- within a large scatter -- by the “Kennicutt-

Schmidt” relationship19 between the surface density star formation rate, SFR, and the gas surface

density, gas; SFR = A{gas}N, where N ~ 1.4. Recent studies of star forming galaxies at

redshifts up to 3 indicate that this relationship may hold at early times as well20,21, suggesting that there is a universal star-forming “main-sequence” for galaxies. However, the intense starbursts triggered by the mergers of large, gas-rich galaxies, both in the local universe and at high-z, produce stars at rates up to 100 times that of the main-sequence galaxies, so that there appears to be at least two modes of the star formation process in galaxies. The merger-induced mode appears to be more intense, which might reflect top heavy initial stellar mass functions, but it is more likely the result of the highly compressed ISM and short dynamical timescales found in the cores of merging gas-rich galaxies.

CCAT’s exquisite sensitivity will be used to extend current knowledge of star formation in galaxies to both lower luminosity and higher redshift systems. With its high spatial resolution, mapping speed, and surface brightness sensitivity, we will explore the applicability of the Kennicutt-Schmidt law in extended low surface brightness regions of nearby galaxies, as well as high surface brightness regions of low-z merger-induced star bursts. With CCAT, we are able to study both the dust-reprocessed starlight and AGN activity in the submillimeter at low-z and in the rest-frame FIR at high z, as well as the molecular gas (the fuel for star formation) through its CO rotational line emission and molecular and atomic gas with [CI] emission. Dense molecular gas tracers with millimeter and submillimeter rotational transitions, such as HCN and CS, indicate quantities of gas associated directly with star formation. For nearby galaxies, CCAT studies will benefit greatly from complementary imaging in the FIR continuum from Herschel/PACS, and images in the low-J CO lines obtained with millimeter-wave                                                             19 Kennicutt, R., ARAA, 36, 189 (1998) 20 Genzel, R., et al., MNRAS, 407, 2091, (2010) 21 Daddi, E., et al. ApJ, 714, L118 (2010) 

 

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interferometers. A goal is for CCAT to work well in the 200 µm telluric windows. 200 µm is close to the dust emission peak, and CCAT 200 µm images will be exceptionally high spatial resolution. Additionally, opening this atmospheric window enables imaging in the [NII] 205 µm line at uniquely high spatial resolution. This line traces the physical conditions of the warm ionized medium, and for ionization-bounded, low-density (n < 50 cm-3), low-ionization (softer than O9 stars) HII regions, tracks the production of Lyman continuum photons. Thus, taken together, molecular line, atomic line, and dust observations with CCAT will characterize the ISM constituents related to star formation in galaxies in an extinction-free manner: CO observations will reveal the molecular gas reservoirs, dense gas tracers will yield the gas fraction involved in star formation, rest-frame FIR (30-200 µm) continuum emission will track star formation rates, rest-frame submillimeter continuum will track dust column density, and [NII] observations will yield the production rate of ionizing photons and the distribution of the ionized gas. Furthermore, access to the 200 µm window will enable studies of the hot molecular tori thought to envelop AGN through their high-J CO line emission. 3.2.2 Nearby Galaxies Goals for CCAT

Several specific scientific goals must be met by CCAT to address star formation in nearby galaxies:

Map nearby galaxies in the dust continuum, CO rotational lines, and [CI] (and [NII]) fine-structure lines at the scales approaching those of individual giant molecular clouds (~40 pc) to resolve the interplay between star formation and the natal ISM. CCAT will use these key diagnostics to infer the physical conditions in spiral arms, bars, and merging galaxies. Where Herschel/PACS or SOFIA imaging in the FIR fine-structure lines (e.g., [CII], [OI], [OIII], [NII], and [NIII]) or the FIR continuum exist, models will be further constrained, addressing questions about the origins and destruction of molecular clouds in star-forming systems and measure variations in star formation efficiency both within resolved galaxies (ranging from low-surface-density to high-surface-density environments) and between galaxies. Of particular interest are the nearby, low metallicity dwarf galaxies (e.g., the LMC) which may serve as templates for low metallicity systems in the early universe. ALMA will address these issues as well, but will not obtain large angular scale maps of the nearby galaxies – those that have ancillary data (FIR fine structure lines and continuum) obtained at the (near) appropriate spatial resolution for good comparison. The properties of these resolved regions will serve as templates to understand the global properties of unresolved galaxies at high redshift, and

test interpretations of the star formation law. Due to their angular size (~ 10 for the

LMC, ~ 1 for M33, and 1/4 for M83), complete maps are impossible or very unlikely with ALMA for these galaxies.

Survey many hundreds of local FIR luminous galaxies at small redshifts in these same lines and other important lines, such as those of water vapor. Though these galaxies will

 

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be only marginally resolved, CCAT will be able to survey the preponderance of warm molecular gas in a wide variety of systems. Is this gas predominately heated by the energy of stars through stellar far-UV light, micro-turbulent shocks, or cosmic rays? Or is it heated by the energy of AGN in X-ray dissociation region (XDR) environments? Do these gas-heating processes influence the formation of the next generation of stars? Does AGN activity inhibit star-formation activity or change the properties of the young stars (initial mass function upper mass cutoff, etc.)? The CO (J = 7 – 6) and [CI] 370 µm lines, which lie only 1000 km s-1 apart, will be particularly interesting for these studies as they are excellent discriminators of photon-dominated regions (PDRs) versus shock excited regions22.

Survey many hundreds of ULIRG or AGN-dominated systems in their CO line emission, extending through the high-J (J ~ 13 – 12) lines accessible in the 200 µm telluric window. These high-J lines may be a “smoking gun” for emission from XDRs formed by AGN X-rays impinging on molecular tori23. XDRs will also strongly emit in H2O lines; with sufficient redshift (z > 1) these diagnostic lines are observable24.

3.2.3 Nearby Galaxies Requirements

The telescope requirements are the same as those identified in Section 3.1.3 except the following (duplicative requirements generally will not be repeated in this document):

Wavelength Range: The requirements are the same as in Section 3.1.3, but the goal of 200 µm is more important because (1) it traces the dust SED for local galaxies closer to

its peak near = 100 µm and (2) it would enable [NII] 205 µm and high-J CO line observations.

Diffraction-Limited Beam Size and Aperture: Same as in Section 3.1.3, but it is desirable

to achieve diffraction-limited imaging at 200 µm. Diffraction-limited (3.5 FWHM) imaging at 350 µm (dust continuum, CO, [CI]]) resolves 0.9 pc structures at the LMC (dist. ~ 50 kpc) and 80 pc structures at nearby spiral galaxies, such as M83 (distance ~

4.5 Mpc). However, the 200 µm diffraction limit of 2 resolves 45 pc (giant molecular cloud-sized) structures at the distance of M83 and helps resolve the nuclear torus from regions of star formation at the distance of nearby AGN-hosting galaxies, such as NGC

1068 (distance ~ 20 Mpc, for which a 2 beam at = 200 µm corresponds to 200 pc).

Sensitivity: In the submillimeter continuum, the bright, nearby (distance ~ 4.5 Mpc) resolved galaxy M83 has surface brightness in the spiral arms in excess of 3 µJy/arcsec2 as measured in the 25” beam of Herschel/SPIRE at 350 µm25. This corresponds to 40 µJy per CCAT beam at 350 µm, well below the confusion limit. We therefore take the

                                                            22 Nikola, T., et al., ApJ 742, 88 (2011) 23 van der Werf, P., et al., A&A., 518, L42 (2010) 24Gonzalez‐Alfonso, E., et al., A&A., 518, L43 (2010) 25 Bendo, G. J., et al., submitted to MNRAS 

 

‐ 12 ‐  

confusion limit as our sensitivity requirement: 5conf = 0.45 mJy/beam at 350 µm.

Assuming 20 K grey-body emission in the dust ( = 2) this corresponds to a total FIR

luminosity ~ 2.7104 L

-- equivalent to detecting dust heated by a single early B star per CCAT beam at 4.5 Mpc.

Heterodyne Arrays: The spectral lines in Local Group galaxies are narrow – they are only 10 km s-1 wide in the LMC. Therefore, much of the Local Group spectroscopy is best done with arrays of heterodyne receivers. The array format should be substantial (> 100 beams) to permit imaging of these systems. The broader lines expected for the more distant, massive galaxies, such as M83 and NGC 1068, may be better addressed with direct-detection systems, depending on the velocity resolution of the spectrometer, the line width, and the line frequency.

3.3 Probing the Astrophysics of Galaxy Clusters with the Sunyaev-Zel’dovich Effect 3.3.1 Sunyaev-Zel’dovich Effect Background

Galaxy clusters are the largest gravitationally bound structures in the Universe and are still forming. Most of their baryons are in the hot intra-cluster medium (ICM), which is in a low density, diffuse, ionized state filling up the space between cluster galaxies. This hot intra-cluster medium (kT ~ 2-10 keV) emits in the X-rays through thermal Bremsstrahlung emission, and is observable at submillimeter to centimeter wavelengths through the Sunyaev-Zel’dovich (SZ) effect, which is a distortion of the Cosmic Microwave Background (CMB) spectrum due to Compton scattering of CMB photons on the ICM electrons26.

The abundance of galaxy clusters and its redshift evolution offer a direct probe of the large-scale structure and information on the structure and dynamics of the universe. In addition, galaxy clusters are important laboratories for studying various astrophysical phenomena in great detail, such as the process of infall and virialization and the relative importance of thermal, turbulent, and bulk motions, mergers of very massive systems, heat transport and instabilities in magnetized plasmas, the evolution of galaxies in dense environments, cosmic circulation of heavy elements, etc. Surveys with the South Pole Telescope27 (SPT) and the Atacama Cosmology Telescope28 (ACT) are delivering catalogs of hundreds of clusters with an approximately mass-limited selection function to high redshift. The Planck Surveyor is surveying the entire sky for thousands of clusters, though with a selection function weighted more to lower redshift29. SZ and X-ray imaging, along with gravitational lensing and dynamical studies, provide a powerful means to measure cluster masses and thereby use large cluster samples as cosmological probes. Detailed multi-wavelength studies of sufficiently large cluster sub-samples

                                                            26 Sunyaev, R.A., and Zel’dovich, Ia., B.  ARAA, 18, 537 (1980) 27 Vanderlinde, K., et al., ApJ, 722, 1180 (2010) 28 Marriage, T.A., et al., ApJ, 737, 61 (2011) 29 Ade, P.A.R., et al., A&A, 536, 8 (2011) 

 

‐ 13 ‐  

at different redshifts are necessary to calibrate these mass measurements and explore the physical processes governing cluster formation and evolution.

CCAT will make unique contributions to SZ studies of galaxy clusters. The large FoV and sensitivity of CCAT will produce high angular resolution cluster images over fields comparable to or larger than the cluster virial radii, down to faint surface brightness limits, which will provide accurate constraints on cluster masses. The angular resolution and wavelength coverage will be sufficient to separate the emission of dusty galaxies and AGN from the cluster thermal SZ, to measure cluster temperatures using relativistic corrections to the thermal SZ effect, and to detect the kinetic SZ effect to yield cluster peculiar velocities and measure internal kinematics. CCAT will bridge the observational gap between Planck/SPT/ACT and ALMA, combining angular resolution with sensitivity to extended emission, thus delivering accurate integrated cluster properties for samples large enough to be of cosmological significance.

3.3.2 SZ Goals for CCAT

CCAT will have a number of objectives in the study of galaxy clusters. CCAT will:

Efficiently detect the SZ effect out to cluster virial radii to measure ICM pressure profiles for large samples of clusters over a wide range in cluster mass and redshift. A 20' FoV enables imaging z > 0.3 clusters efficiently (requirement FoV > R200)

30. The sufficiently stable atmosphere above Cerro Chajnantor may also enable imaging lower redshift clusters with novel sky-noise removal methods. This will enable the study of the baryon pressure, density, and temperature from cluster cores, through the virialized regions, out to the radii at which material is infalling and thermalizing. Such measurements will also enable the characterization of deviations from electron-ion thermal equilibrium and from hydrostatic equilibrium and self-similarity due to incomplete virialization, feedback from star formation and AGN, radiative cooling, and non-thermal pressure support, e.g., from turbulent bulk gas motions31,32,33.

Resolve structure in the ICM pressure and velocity without filtering away diffuse SZ signal. This will enable:

o The study of the apparent pressure dichotomy near the cluster cores (cool core vs. non-cool cores) observed in X-rays (see, e.g., Arnaud et al. 201034).

o The detection of structure encoding the merger history, for example merger shocks and the gas / dark matter offset obtained through SZ35.

                                                            30  R200 here is the radius at which the enclosed mean mass exceeds the critical density by a factor of 200.  Clusters are expected to be virialized within these regions.   31 Lau, E.T., et al., ApJ, 705, 1129 (2009) 32 Shaw, L.D., et al., ApJ, 725, 1452 (2010) 33 Battaglia, N., et al., ApJ, 725, 91 (2010) 34 Arnaud, M., et al., A&A, 517, 92 (2010) 35 Molnar, S.M., et al., arXiv:1201.1533 (2012) 

 

‐ 14 ‐  

o Measurement of the ICM turbulence from the power spectrum of pressure fluctuations. The expected amplitude for pressure fluctuations at the CCAT resolution of 40-400 kpc in low-z clusters is ~10%36, thus the SZ signal must be

detected at ~ 50 per beam. o The search for nonthermal electron populations, especially in correlation with

cluster radio halo measurements and -ray observations37. o Locate the hottest gas in the shock-heated regions, complementing X-ray data38,39. o The study of bulk flows and potentially residual nonthermal velocities inside

clusters via the kinetic SZ effect40.

Use relativistic SZ (rSZ) corrections to measure the cluster electron temperature directly.

Measure the integrated Comptonization for use in measurements of scaling relations between integral quantities, which may be used in calibrating cluster masses for measurement of cosmological parameters using other surveys to define a selection function (Planck, eROSITA, DES, and LSST; e.g., as was done with X-ray data41).

Obtain higher angular resolution information on clusters detected in the SPT, ACT, and Planck surveys and complement eROSITA X-ray data, both to measure the evolution of star formation in cluster environments and to measure the spatial distribution of the ICM.

Measure the SZ signal from optical/NIR-detected clusters (e.g., those from Spitzer and WISE), and thereby probe 3 to 5 times lower masses than current SZ surveys, enabling the study of the role of feedback processes42.

Accurately measure the thermal SZ (tSZ) anisotropy power spectrum at high angular multipole number, l ~ 2,000 - 20,000, and a possible turn-over around l ~ 5000, as has been initially done by SPT43 and ACT44. This measurement will test the validity of different ICM models because simulations (e.g., Battaglia et al 2010, Shaw et al 2010) indicate that the amplitude and position of the peak is determined by feedback from AGN and star formation. CCAT’s angular resolution and mapping speed, combined with the ability to remove contaminating point sources that dominate the anisotropy signal at these angular scales, make it the ideal instrument to trace the tSZ power spectrum for use as a cosmological probe.

Use the kinetic SZ (kSZ) effect to measure cluster peculiar velocities and potentially map cluster velocity fields45.

                                                            36 Schuecker, P., et al., A&A, 426, 387 (2004) 37 Pfrommer, C., et al., MNRAS, 378, 385 (2007) 38 Mason, B.S., et al., 716, 739 (2010) 39 Korngut, P.M., ApJ, 734, 10 (2011) 40 Diego, J.M., et al., ApJ, 597, 1 (2003) 41 Mantz, A., et al., MNRAS, 406, 1759 (2010) 42 Hand, N., et al., ApJ, 736, 39 (2011) 43 Reichardt, C., et al., arXiv1111.0932 (2011) 44 Dunkley, J., et al., ApJ, 739, 52 (2011) 45 Mak, D.S.Y., et al., ApJ, 736, 116 (2011) 

 

‐ 15 ‐  

Measure the kSZ power spectrum from patchy reionization and thereby constrain the reionization history of the universe46,47.

CCAT’s excellent angular resolution and broad wavelength coverage at trans-millimeter wavelengths are key for this work: these are crucial for identifying dusty, star-forming and radio-loud galaxies and to study star formation in cluster environments. ALMA’s small FoV (1') and angular filtering renders it unsuitable for most of the above measurements, but can provide complementary measurements of small scale structures detected but not resolved by CCAT.

3.3.3 SZ Requirements

To attain the galaxy cluster goals, the following additional requirements must be satisfied:

Wavelength Range: 870 µm – 2.1 mm required, with a goal of 740 µm – 3 mm. The multiple bands are required for simultaneous measurement of thermal, kinetic, and relativistic SZ effects and for removing contamination by primary CMB, dusty star-forming galaxies, and radio galaxies. Without five to six spectral bands, outside information would be required for separation of the various components.

Discussion of 3 mm (90 GHz) goal: o Ninety GHz may allow for better measurement of the kSZ effect and electron

temperatures48. o Identification of radio point sources: CCAT will have sufficient resolution to

resolve radio sources in galaxy clusters at 150 GHz (FWHM beam 26'' – recent MUSTANG observations by Korngut et al. 201139). Although 90 GHz observations would provide additional frequency leverage in distinguishing spectral components, with a 40" FWHM beam at 90 GHz, a proper analysis of the kSZ in combination with radio point sources might require a degradation of the 2 mm resolution and a simultaneous fit to point-source SED and normalization along with the cluster peculiar velocity. ALMA would be better to identify flat-spectrum radio sources if observations can be done contemporaneously.

o Ninety GHz imaging would aid in the removal of the CMB at large cluster-centric radii.

o A 90 GHz band should be implemented only if it carries a small or negligible incremental cost, which seems likely in the LWCam multi-scale pixel designs

                                                            46 Dvorkin, C., et al., PhRvD, 79, 7302 (2009) 47 Zahn, O., et al., arXiv1111.6386 (2011) 48 However, Holder, G., ApJ, 602, 18 (2004) and Sehgal, N.,  et al., ApJ, 635, 22 (2005) show that the best three‐

frequency setup to recover r and/or Te with SZ (X‐ray independent) measurements is 30, 150, and 300/350 GHz. Sommer (PhD thesis 2010) showed that using 90 GHz instead of 30 GHz at the same signal‐to‐noise ratio doubles 

the projected uncertainty on Te, i.e., 90 GHz does not significantly improve a determination of r or Te if the 150 GHz measurement is available.   

 

‐ 16 ‐  

being considered due to the nature of the architecture and the small added pixel count.

FoV: To image clusters to the virial radius CCAT must have the largest FoV possible. Cluster sizes R200 in arcmin for three masses are shown in Figure 1 and Table 2.

Figure 1 and Table 2−Cluster sizes in R200 in arcminutes for various cluster masses.

Scanning speed: With slow scans, sky noise is eliminated by removing the average signal over the FoV, which attenuates astrophysical signal on scales comparable to the FoV. This places a lower limit on the redshift at which one can detect appreciable signal at a cluster virial radius for a given mass. A robust capability for observing large-scale signal at the virial radius and beyond – a key aspect of the cluster science – requires rapid

scanning. Assuming a characteristic sky noise timescale of 0.5 – 1 s, scanning at 1 s-1 will enable high fidelity imaging on scales of 30' to 60'. Correspondingly large

acceleration (1 s-2) is required to ensure high efficiency during turnarounds. These

needs are met by the requirement (0.33 s-1) x ( / 350 µm) in Az, and half that in El (for

scanning at El up to 60, Az scan rate 2 El scan rate) previously specified.

Emissivity: The emissivity requirements on the telescope are discussed in a separate note “Long-Wavelength Emissivity Requirement for CCAT” by S. Golwala (CCAT Science Memo No. 102). The emissivity requirements are driven by instrumental considerations on what is feasible rather than by a science consideration to obtain the best possible sensitivity. This results in requirements of 5% emissivity for 1 mm to 3 mm and 10% at 750 µm and 850 µm.

Diffraction-Limited Beam Size and Aperture: The same as in Section 3.1.3: Beam size at the shorter wavelengths determines the confusion limit and thus the precision for removal of dusty galaxies. This will determine the ultimate sensitivity limit for SZ measurements.

Redshift Mass (Msolar)   Z 3.5x10

141.0x10

15  3.5x10

15

R200 (') R200 (')  R200 (') 

0.1 12.7 18.1  27.4 0.2 6.9 9.8 14.8 0.3 4.9 7.0 10.6 0.4 3.9 5.6 8.4 0.5 3.3 4.7 7.1 0.6 2.9 4.1 6.3 0.7 2.6 3.7 5.6 0.8 2.4 3.4 5.2 0.9 2.2 3.1 4.8 1.0 2.1 2.9 4.5 1.1 2.0 2.8 4.2 1.2 1.9 2.6 4.0 1.3 1.8 2.5 3.8 1.4 1.7 2.4 3.7 1.5 1.6 2.3 3.5 

 

‐ 17 ‐  

Thus, the telescope (and instrument) should be diffraction-limited at submillimeter wavelengths. At longer wavelengths, the angular resolution should be considered as part of the tradeoff with total optical loading and sensitivity: The best possible angular resolution favors uniform primary mirror illumination by the instrument (in the time-reverse sense), but this may yield substantial spillover at a cold stop or at the primary mirror, which will degrade the background-limited sensitivity.

Sensitivity: The goal is background-limited sensitivity in all bands, where the background is set by the combination of the atmosphere, the telescope emissivity, and the instrument loading. See the Emissivity requirement above.

Scan Patterns: Raster scans and Lissajous pattern required. Flexibility is required to develop additional scan patterns.

3.4 Galactic Molecular Cloud and Star Formation

There are two major areas of Galactic science in which CCAT can contribute significantly: the formation of molecular clouds and star formation. To make progress in these areas will require large-scale surveys, at high spatial resolution, with a) heterodyne array receivers to study the kinematics, dynamics, physical conditions and chemical complexity of the gas on large to small scales and b) large format submillimeter continuum cameras to trace the extended structure of the filaments and the cold, dense cores within them that are the potential sites of star formation.

3.4.1 The Formation of Giant Molecular Clouds, Filaments, and Quiescent Cloud Cores

3.4.1.1 Background

At any instant, only a few percent of molecular cloud masses have achieved sufficiently high density to form stars, as most of the cloud masses are contained in a low-density molecular state that is devoid of active star formation. The transition from a low-density state to a star-forming state is the first and biggest step on the path to stars. Understanding how this happens and why such a small fraction of the cloud mass is contained in dense gas is key to determining the dominant physical processes that control the star formation rate in molecular clouds.

Most modern theories of molecular cloud formation describe the process as being rather fast: starting from atomic clouds that are compressed, either through cloud-cloud collision, or through the potential of Galactic spiral arms49,50,51. As a result of this compression, the gas quickly becomes molecular, forms dense turbulent structures, some of which collapse gravitationally. In this picture, only a small fraction of the gas is able to proceed to star formation, since most is in transient structures that eventually disperse.

                                                            49 Vázquez‐Semadeni, E., et al., ApJ, 657, 870 (2007) 50 Heitsch, F., & Hartmann, L., ApJ, 689, 290 (2008) 51 Ballesteros‐Paredes, J., et al., ApJ, 637, 384 (2006) 

 

‐ 18 ‐  

A key feature of this model is that the development is very dynamic, and clouds never reach a quiescent, stationary state. As a result, molecular clouds should exhibit a rich chemical diversity that will be manifested on all spatial scales, from the large-scale cloud to individual star forming clumps. Dust continuum images from the Herschel Space Observatory have revealed the large-scale structure in spectacular detail, and have especially emphasized the importance of filamentary structures within molecular clouds.

3.4.1.2 Science Goals for CCAT

CCAT has two critical attributes that will enable it to connect the large-scale structure of molecular clouds to individual star forming sites. With the large FoV to mosaic regions several degrees in diameter, and an angular resolution of a few arcseconds at the highest frequencies, CCAT will determine how molecular material flows from the large-scale diffuse cloud to the compact dense regions destined to form stars. The main scientific goals are:

To study the extended, diffuse ISM outside the large-scale filaments through the presence of hydrides, deuterated hydrides, and select molecular ions. In many cases, the sub-millimeter continuum from background sources will be resolved by CCAT and, therefore, it will be possible to map absorption lines from many of these species52 and, therefore, determine their abundances.

To trace the transition from diffuse atomic gas to giant molecular clouds. Theoretical models predict the formation of CO, the classical tracer of molecular gas, lags behind the formation of H2, which is generally unobservable directly. A consequence of this formation lag is the existence of "dark" or "invisible" molecular gas, containing H2 but no CO gas. This "dark" molecular gas is expected to be traced by atomic transitions from carbon (both ionized and neutral carbon) and light hydrides like CH+, 13CH+, SH+, etc. which have been shown to be good tracers of overall molecular mass53. Two neutral carbon fine structure transitions are accessible to CCAT, at 492 GHz and 809 GHz. Specifically, CCAT should be able to map Galactic Center region (1 degree) in the 860 GHz atmospheric band to an RMS sensitivity of 0.1 K with a spectral resolution of 0.5 km/s in approximately 100 hours of integration time.

To study the formation of filaments and compact dense cores. Herschel continuum observations have identified a complex web of filaments that pervade the interstellar medium. These filaments further break up into dense, compact clumps some of which make a phase transition from turbulence to coherence and, then, become the formation sites of stars. A key aspect to understanding star formation is determining how the molecular gas transitions from the diffuse interstellar medium, to the filaments, and then to individual cores. High velocity resolution and high angular resolution molecular spectroscopy carried out over large spatial scales is the key to tracing the flow of material

                                                            52 Naylor, D. et al., A&A, 518, L117 (2010) 53 Falgarone, E. et al., A&A, 521, L15 (2010) 

 

‐ 19 ‐  

over the full range of angular scales from large molecular clouds to individual star forming regions.

To study the chemical diversity of star forming regions. As molecular gas transitions from diffuse material to dense clumps to newly formed stars, the molecular composition of the gas will change dramatically with the temperature and density of the gas. This leads to an intriguing, and complex, time-dependent chemistry that contains information on both the evolutionary and physical state (density, temperature, etc.) of the molecular gas. Spectral mapping surveys of hundreds of dense cores over many molecular transitions are needed to understand empirically the physical and chemical differences between different cores, and to establish a chronological sequence to core formation.

3.4.1.3 Heterodyne Spectroscopic Requirements

No single molecule can probe cloud structure over all spatial scales. Each molecular transition probes different volume densities and temperatures, and time-dependent chemistry favors the production of different molecules at distinct times. Therefore, multiple receivers are needed to cover the different atmospheric bands.

Heterodyne receivers: The spectral lines in molecular clouds will have linewidths on the order of 1 km/s. However, in the cold cores that have lost their turbulent support, linewidths are much narrower (on order 0.2 km/s). Thus, heterodyne receivers are required to achieve the necessary spectral resolution. The IF bandwidth should be at least 8 GHz to allow efficient spectral mapping of the atmospheric windows. The receiver in a given band should be a focal plane array to enhance imaging speed. The mapping speed of the heterodyne receiver for a single spectral line (factoring number of array elements and system temperature) should exceed that of ALMA. A naïve calculation just comparing collecting area would indicate that the current ALMA, with 16 elements and 3.6 times the collecting area of CCAT is14 times faster than CCAT for heterodyne mapping, taking into account that an ALMA beam encompasses approximately 4 CCAT beams. However, this crude calculation ignores the fact that CCAT will have a better surface than the ALMA antennae and will be at a better site, so CCAT should perform much better than the ratio of the collecting area suggests - particularly at high frequencies. The reduction in telescope gain as a function of rms surface accuracy is given by the

Ruze equation where is the rms surface error (10 m for CCAT and 25

m for ALMA) and is the wavelength. Thus we can estimate the sensitivity

improvement of CCAT as being ,which can effectively be approximated as

a reduction in the system temperature (Tsys). The improved atmosphere at the CCAT site also decreases Tsys, since / (ignoring gains and efficiency

factors). Thus, we can estimate the reduction in CCAT system temperatures by the ratio

 

‐ 20 ‐  

of between the ALMA and CCAT sites. We calculate using the online CSO atmospheric plotter. ALMA system temperatures are obtained from the ALMA sensitivity calculator. Results are given in Table 3. Table 3—System temperature decrease at CCAT accounting for improved atmospheric transmission and better surface accuracy .

Frequency 

(GHz) Tsys 

ALMA (K) 

ALMA    

1 mm pwv 

CCAT       0.6 mm pwv 

CCAT Gain 

Improvement 

CCAT  Atmosphere Improvement 

Tsys CCAT (K) 

345  177  0.19 0.12 1.12 1.07  148460   1057  0.84 0.52 1.22 1.38  631660  1223  1.16 0.73 1.49 1.55  527860  2242  1.33 0.84 1.98 1.63  696

Given these improved system temperatures and ignoring, for the time being, that ALMA will filter out extended spatial structure, we find that at 345 GHz we only need 10 pixels to match the mapping speed of ALMA with 16 antennae, and 30 pixels to match ALMA’s mapping speed with 50 antenna. Due to the vastly improved system temperatures at 860 GHz, however, we would only need 5 pixels to match ALMA’s mapping speed. To estimate the mapping times, we follow the approach outlined by P. Goldsmith at the Cologne CCAT conference (2011) and assume that we are performing a large survey

(surv > b) with Nyquist sampling and reconstructing the final image to an angular

resolution of FWHM. For a single pointing, the radiometer equation is:

∆

where surv is the frequency resolution in Hz, tint is the integration time in

seconds, Tsys is the system temperature, and Trms is the noise level. For a large survey,

ΩΩ Δ

where surv is the solid angle of the survey, b is the solid angle of the beam

(1.132FWHM), Npix is the number of pixels in the array, tsurv is the total survey time, and

is the efficiency of the mapping scheme (including calibration, slewing, off positions, etc) which we assume to be 0.5. Therefore, the survey time is given by:

Ω

Ω Δ

Assuming that we want to map a 1 degree area to a sensitivity of 0.1 K with a spectral resolution of 0.5 km/s, the above equation provides the required survey times. In Table 4 we compare the ALMA survey times (assuming the full complement of 50 antennae) with those of CCAT, for heterodyne arrays of 50 and 100 pixels. However, since

 

 

Aa

Tr

H

Wttsia

F

ALMA will array of 4 12

Table 4—Apprresolution of 0.

Frequen

345 GH460 GH660 GH860 GH

High Priority

Wavelength tunable betwthe receiver shows the atinterest for tand transition

Figure 2—Spec

resolve out 2-m antennae

roximate time 5 km/s.

cy  Time ALMAantenn

z  5.0 hz  237 hz  455 hz  1,992 

y Bands: 460

coverage foween 385 and

be tunable ttmospheric tthe 460 GHzn frequencie

ctral lines and

extended stre.

to map 1 squa

for  (50 nae) 

Time fTota

A

hrs  6hrs  2,6hrs  5,6hrs  24,9

0, 660, & 86

r the 460 GHd 530 GHz tto cover COtransmissionz atmospheres in the win

CCAT atmosp

‐ 21 ‐ 

ructure, we

are degree of s

for ALMAal Power Array 

T

2 hrs691 hrs687 hrs905 hrs

60 GHz

Hz window: to cover the

O J=4-3 at 46n from the Cric window.ndow.

pheric transmis

also compar

sky to a sensit

Time for CCAT(50 pixels) 

3.2 hrs78 hrs78 hrs177 hrs

Ideally, thee atmospheri62 GHz andCCAT site a Table 5 lis

sion for 400 G

re with ALM

tivity of 0.1K

T Time for (100 pix

1.6 hr39 hr39 hr89 hr

e 460 GHz ric window. Id [CI] at 492and various sts various m

GHz to 550 GH

MA’s total p

(1) with a sp

CCAT xels) 

rs rs rs rs 

receiver wouIt is required2 GHz. Figspectral lin

molecular sp

Hz.

power

pectral

uld be d that

gure 2 nes of pecies

 

‐ 22 ‐  

Table 5—Some potentially interesting species in the 460 GHz window.

Species  Frequency (GHz) Species Frequency (GHz)

CH3OH  385 – 530 HCO+ 5‐4 445.902 13CO 4‐3  440.765 HDCO 386.165 – 529.597

CO 4‐3  461.039 N2H+ 5‐4 465.824 

C18O 4‐3  439.088 N2D+ 5‐4, 6‐5  385.516, 462.603

CCS  388.1 – 521.7 NH2D 391.349 – 526.107

CH3OCH3  385 – 530 OCS 389.044 – 522.594

CN   449.212 – 453.638 S34O 392.511 – 521.799

CS 8‐7, 9‐8, 10‐9  391.8, 440.8, 489.8 S34O2  

D2CO  385 – 530 SO  387.328 – 527.941

D2O  393.3 – 525.4 SO2 385 ‐ 530 

DCN 6‐5, 7‐6  434.4, 506.8 [CI] 1‐0 492.1 

DCO+ 6‐5, 7‐6  432.2, 504.2 SiO 9‐8,10‐9,11‐10 390.7, 434.1,520.8

DNC 6‐5  457.776 CF+ 4‐3, 5‐4 410.304, 512.846

H2CO 6‐5, 7‐6  389 – 525 CD+ 1‐0 453.521 

H2S   392.6 – 521.1 H2Cl+ 1‐0 395.076 – 485.420

H3O+ 3‐2  388.458, 396.272 NO+ 4‐3 476.733 

HC3N  391.045 – 527.266  

HCN 5‐4  443.116  

Wavelength coverage for the 660 GHz window: Ideally, this receiver should be tunable over the range 600 to 720 GHz to cover the atmospheric window. Figure 3 shows the atmospheric transmission from the CCAT site and various spectral lines of interest for the 660 GHz atmospheric window. Table 6 lists various molecular species and transition frequencies in the window.

 

 

F

S

C1

C

C

C

C

C

D

D

D

D

H

H

H

H

Woc

Figure 3—Spec

Species 

CH3OH 3CO 6‐5 

CO 6‐5 

C18O 6‐5 

CN  

CH3OCH3 

CS 13‐12, 14‐13

D2CO 

D2O 

DCN 9‐8 

DCO+ 9‐8 

H2CO 

H2S  

HC3N 

HCN 7‐6, 8‐7 

Wavelength over the rancompelling a

ctral lines and

Table 6—S

Fr

6

6

6

6

6

6

3  6

6

6

6

64

6

6

6

6

coverage fonge 787 toarguments fo

CCAT atmosp

ome potentiall

requency (GHz

00.202 – 719.6

61.067

91.473

58.553

73.779 – 680.2

00.448 – 719.7

36.532, 685.43

01.300 – 718.9

07.349 – 714.0

51.565

48.193

00.330 – 716.9

11.441 – 708.4

08.919 – 717.6

20.304, 708.87

or the 860 G 960 GHz

or this receiv

‐ 23 ‐ 

pheric transmis

y interesting sp

z) Sp

664 HC

HD

HD

HN

296 N2

722 N2

35 SH

933 SO

087 SO

SiO

CF

938 H2

470 NO

678 H2

77

GHz windowto cover t

ver are 13CH

sion for 580 G

pecies in the 6

pecies

CO+ 7‐6, 8‐7

D2+ 4‐3, 1‐1

DCO

NC 7‐6

2H+ 7‐6

2D+ 7‐6, 8‐7

H+ 1‐0

O 

O2

O

F+ 6‐5, 7‐6

2Cl+

O+ 6‐5

2D+ 3‐3

w: Ideally, thithe atmosph

H+ 1 – 0 and

GHz to 740 GH

60 GHz windo

Freq

624

635

603

634

652

616

683

601

600

603

615

621

715

646

 

is receiver sheric windothe [CI] 2 –

Hz.

ow.

quency (GHz)

4.208, 713.341

5.679, 691.660

3.595 – 719.985

4.510 

2.095 

6.750, 693.806

3.336 – 683.448

1.258 – 719.751

0.534 – 719.870

3.382 – 694.293

5.365, 717.857

1.932 – 713.521

5.019 

6.430 

should be tunow. Two m– 1 transition

5

8

1

0

3

1

nable major n that,

 

 

cdmcFls

F

S

S1

C

D

H

S

[

C

D

H

N

S

combined wdensity. Botmolecular gachallenges inFigure 4 sholines of interspecies and t

Figure 4—Spec

Species 

SH+ 2‐1 3CH+ 1‐0 

CS 17‐16,18‐17

DCO+ 11‐10,12

HCN 9‐8, 10‐9 

SO 

CI] 2‐1 

CH3OH 

D2CO 

H2CO 

N2D+ 11‐10,12‐

SO2 

with the [CIh 13CH+ andas. This recn implemen

ows the atmorest for the 8transition fre

ctral lines and

Table 7—So

Fr

8

8

7,19‐18  8

2‐11,13‐12  7

7

8

8

7

7

7

‐11 84

7

I] 1-0 transd [CI] haveceiver, therenting low noospheric tran860 GHz atmequencies in

CCAT atmosp

ome potentially

requency (GHz

93.047 – 893.1

30.216

32.1, 880.9, 92

92.1, 864.1, 93

97.433, 885.97

00.685 – 946.7

09.341

87.342 – 959.9

87.0 – 957.134

89.815 – 959.7

47.877, 924.88

87.301 – 959.5

‐ 24 ‐ 

sition, uname been showefore, has a oise receivensmission frmospheric wthe window

pheric transmis

y interesting sp

z) Sp

133 CO13C

29.7 C1

36.0 CC

70 D2

785 H2

N2

900 CF

4 CH

706 DC

89 HC

507 OH

mbiguously wn to be tra

high sciencers at these from the CCwindow. Taw.

sion for 760 G

pecies in the 86

pecies

O 7‐6, 8‐7

CO 8‐718O 8‐7

CH

2O

2Cl+

2H+ 9‐8, 10‐9

F+ 8‐7, 9‐8

H+ 1‐0

CN 11‐10,12‐1

CO+ 9‐8, 10‐9

H+ 1‐0

gives one tacers of the ce priority bfrequencies

CAT site andable 7 lists v

GHz to 980 GH

60 GHz window

Freq

806

881

877

872

836

789

838

820

835

1,13‐12  796

802

909

the [CI] coso-called “d

but the techs are recognd various spevarious mole

Hz.

w.

quency (GHz)

6.651, 921.799

1.272 

7.921 

2.435 – 959.581

6.330 – 952.242

9.671 – 912.758

8.307, 931.385

0.317, 922.740

5.137 

6.2, 868.6, 950.

2.458, 891.557

9.045, 909.158

olumn dark” hnical nized. ectral ecular

1

2

8

.9

 

 

L

Wtaissra

F

C

C

C

C

C

C

D

D

Sas

Lower Priori

Wavelength tunable betwand N2H

+ J =in this bandsubmillimetesensitivity anreceiver wouareas of the s

Figure 5—Spec

CH3OH 

CCH 4‐3 

CCS 

CN 3‐2 

CO, C18O 3‐2 

CS 6‐5, 7‐6 

D2CO 

D2O 1‐1 

Spectrometerachieve a vspectrometer

ity Band: 34

coverage foween 272 GH= 4 – 3), see d have been,er telescopend high altituld be of prsky.

ctral lines and

Table 8 – S

DC

H2

H2

H2

HC

HC

H

H

r: The spectvariety of r should be a

5 GHz

or the 345 GHHz (to observ

Figure 5 and, and will cs, this wavetude site. Brime use for

CCAT atmosp

ome potentiall

CO+  4‐3, 5‐4

2CO 4‐3, 5‐4

2D+ 1‐1

2S 3‐3, 4‐4

C3N  

CN 4‐3

DCO 5‐4

NCO 

trometer shoobservation

able to achie

‐ 25 ‐ 

Hz window:ve N2H

+ J =d Table 8. H

continue to beband does ecause of itsr mapping e

pheric transmis

y interesting sp

N2

N2

S34

S34

SH

SO

SO

SiO

ould be flexal goals. T

eve 0.1 km/s

: Ideally, the= 3 – 2) andHowever, givbe, performnot take a

s low systemextremely la

sion for 320 G

pecies in the 34

H+ 3‐2, 4‐3

D+ 4‐34O

4O2 

+ 1‐0

O 

O2 

O 7‐6, 8‐7

xible (i.e., reTo image qspectral res

e 345 GHz rd 373 GHz (ven the fact

med relativelydvantage ofm temperatuarge (> 100

GHz to 400 GH

45 GHz windo

CF+ 3

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econfigurablquiescent dolution (with

receiver wou(to observe that observay easily at f CCATs unure, however0 square deg

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ow.

3‐2 

+ 

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uld be H2D

+ ations other nique r, this grees)

it can , the s as a

 

‐ 26 ‐  

goal) and thereby resolve the thermal linewidth in molecular clouds. Mapping surveys of the Galactic plane should be able to cover the entirely velocity range of the Galaxy in a single tuning for a given spectral line; this will require a minimum bandwidth of 1,000 km/s to cover the wide range of velocities observed toward the Galactic Center at 0.5 km/s resolution. The IF bandwidth should be 4 GHz – sufficient for any line and for observing CO J = 7 – 6 and CI J = 2 -1 simultaneously. Finally, to enable efficient spectral mapping in an atmospheric window, the goal is to have a spectrometer that can cover the entire IF receiver bandwidth at 0.5 km/s resolution. To achieve the desired flexibility the spectrometer would likely need to be a next-generation signal processing system based on Field Programmable Gate Array (FPGA) based computing. The spectrometer should meet the following specifications: 1) > 100 channels (> 50 pixels, dual polarizations), 2) Sampling bandwidth of 3.6 Giga-samples/second with 4-8 bit sampling, 3) at least 1.8 GHz bandwidth (1,000 km/s bandwidth at all wavelengths) with a goal to cover the entire IF bandwidth, 4) Sub-banding to look at multiple narrow lines in each channel.

Map Sizes: The nearest giant molecular cloud (GMC) complexes span tens of square degrees on the sky. A Galactic plane survey will require mapping hundreds of square degrees. The telescope needs to be able to perform on-the-fly mapping in equatorial and galactic coordinates. In order to map entire GMC complexes, CCAT needs scan patterns and stability suitable to obtain Nyquist sampled maps that are > 100 square degrees in size. Using the same calculation as above, at 345 GHz 100 square degrees of sky could be mapped to a sensitivity of 1 K in 293 hours with a 100 pixel array.

Aperture size: The goal is to resolve individual protostellar envelopes in the nearest massive GMC: the Orion cloud complex at a distance of 400 pc. This requires an angular resolution < 7, and ideally 3.5 to have a spatial resolution of 3,000 AU at Orion (protostellar envelopes are typically 5,000 AU in diameter, so the resolution must exceed this to measure sizes). This can be achieved with a 25 m telescope for frequencies greater than 430 GHz.

Calibration: Absolute calibration of 20% is a desired goal. This would likely require the use of both hot and cold loads and rapid monitoring of the atmospheric conditions in the direction of the observation. Thus, the calibration goals will probably require a PWV monitor similar to the 183 GHz WV monitors used by ALMA in order to obtain accurate measurements of the atmospheric opacity concomitantly with the astronomical observations. Such monitors are becoming relatively commonplace and easy to obtain.

Frequency Switching: The ability to frequency switch is also a desired goal since, if it can be made to work properly, it can result in significant time savings. This may require occasional nodding to an off-position to calibrate out the bandpass and also places conditions on bandpass stability.

 

‐ 27 ‐  

3.4.2 Star Formation and the Clump Mass Function

3.4.2.1 Background

The distribution of stellar masses that form in a local volume of space, also known as the stellar Initial Mass Function (IMF), shows remarkable consistency in environments ranging from globular clusters, to starburst clusters, and to sparse star-forming regions54,55. The physical processes that lead to a seemingly invariant IMF in the local universe have long been the subject of theoretical conjecture. Gravitational or turbulent fragmentation, feedback from stellar winds and outflows, competitive accretion, ejection of protostellar cores, and stellar mergers have all been proposed to explain various aspects of the IMF shape56.

Direct evidence to indicate which of these mechanisms, if any, play a dominant role in governing the IMF shape remains elusive. One intriguing observational development is emerging evidence that the mass function of dense “clumps'' in molecular clouds is similar in shape to the stellar IMF57. These clumps are the final stage in the flow of material from the diffuse ISM to molecular clouds/filaments. Since many of these clumps have sizes and masses needed to form individual stars, the inference is that the clump mass function directly translates into the stellar IMF. If this is indeed the case, it will provide compelling evidence that the stellar IMF is imprinted in the fragmentation structure of molecular clouds. Otherwise, alternative mechanisms or a combination of processes may be needed to explain the origin of the IMF.

3.4.2.2 Science Goals for CCAT

A number of submillimeter continuum surveys of the Galactic plane have been conducted to measure the clump mass function in clouds, both by ground-based (ATLASGAL, Bolocam, SCUBA-2) or space-based (Hi-GAL, Herschel Gould belt survey, HOBYS) facilities. Despite the intriguing observational results to date, the results are far from conclusive. Two observational tests can establish a definitive link between the clump mass function and the stellar IMF.

Measure the clump mass function to masses as low as 0.01 Msun. If the clump mass function follows that of the IMF, it should peak around 0.5-1 Msun and decline toward lower masses assuming that 30% of the clump mass is converted into stars. CCAT should detect clumps with masses as low as 0.01 Msun at the 10 level to make a definitive determination that the clump mass function declines toward the brown dwarf regime.

Measure the clump mass function in low and high mass star forming regions. If the clump mass function traces the stellar mass function, then the shape of the clump mass function should be invariant in low and high mass star forming regions, as is the stellar mass function. This requires detection 0.03 Msun clumps out to distances of at least 400

                                                            54 Kroupa, P. Science, 295, 82 (2002) 55 Chabrier, G., PASP, 115, 763 (2003) 56 See review by Larson, R. B., RPPh, 66, 1651 (2003) 57 Ward‐Thompson, D., et al., PPV, 951, 33 (2007) 

 

‐ 28 ‐  

pc, which will encompass both high mass (Orion) and low mass (e.g., Chamaeleon, Ophiuchus, and Lupus) star forming regions.

3.4.2.3 Requirements

Measuring the clump mass function is best performed in the submillimeter continuum, which is generally optically thin and can probe trace amounts of dust. Additionally, to get the actual mass, one needs an estimate of the distance to the clumps which can be obtained kinematically through a combination of heterodyne observations and a model of Galactic rotation.

Wavelength range: 350 m continuum observations are required to achieve high sensitivity and angular resolution. As a goal, 850 m observations are useful to measure the dust spectral index, and 200 m observations can be used to measure dust temperatures.

Angular resolution: Protostellar envelopes are approximately 5,000 AU in diameter. To resolve individual protostars out to a distance of 400 pc, with at least 3 resolution elements across the diameter of the clump, requires an angular resolution of 3.5. This can be achieved at a wavelength of 350 m with a 25 m diameter telescope. This high angular resolution will tie in nicely with the results from the Herschel Space Observatory at shorter wavelengths58.

Sensitivity: For a dust opacity of = 0.31 (/230.0 GHz) cm2 g-1, the 350 m continuum observations need to reach a RMS noise level of 0.9 mJy to detect a 0.01 Msun clump at a distance of 400 pc at the 10 level. To observe one square degree to this depth in ten hours requires a 6.5' FoV for an NEFD of 14 mJy s1/2.

3.5 Solar System Studies

The primary submillimeter Solar System studies enabled by CCAT fall into two categories: Trans-Neptunian Objects (TNOs) and isotopic composition of comets.

3.5.1 Trans-Neptunian Objects

3.5.1.1 TNO Background

Trans-Neptunian Objects59,60 are the most primitive remnants left from the planetesimal building stage of the solar nebula. Over a thousand of these small bodies have now been discovered beyond the orbit of Neptune, in the region referred to as the Kuiper Belt, and the total population of objects brighter than 24th optical magnitude is estimated at approximately 30,000. The total mass of the Kuiper Belt, a local analog of the debris disks seen around nearby stars, is estimated

                                                            58 Molinari, S., et al., A&A, 518, L100 (2010) 59 Morbidelli, A., et al., in The Solar System Beyond Neptune, ed. M.A. Barucci, et al. UIP, Tucson, 275 (2008) 60 Müller, T. G., et al. EMP, 105, 209 (2009) 

 

‐ 29 ‐  

at 0.03–0.3 Earth masses61. The dynamical and physical properties of TNOs provide key constraints for the Solar System formation and evolution models. While these objects can be easily discovered, and their orbits determined, from observations at visual wavelengths, the determination of their physical properties (i.e., size, albedo, density, as well as thermal properties) remains challenging and requires measurements of their thermal flux at far-infrared or submillimeter wavelengths.

A sample of 47 TNOs was observed with Spitzer MIPS at 24 and 70 m62, showing correlations between geometric albedo and perihelion distance and albedo and size. However, given their large heliocentric distances, TNOs are cold (20–50 K) and their thermal emission is best studied in the far-infrared/submillimeter. Herschel is observing the thermal emission of 139 TNOs, including 25 known multiples63,64, although only the brightest 15 will be observable at submillimeter wavelengths.

CCAT offers unprecedented capabilities to extend observations of TNOs into the submillimeter regime. The CCAT confusion noise65 at 15 beams per source is approximately 0.25 mJy and 0.6

mJy at 350 m and 450 m, respectively. This sensitivity can be achieved in a 1-hour

observation at 350 m and in 6 min at 450 m, enabling detection at 350 m of a 1,000 km diameter TNO up to 100 AU distance and a 100 km TNO up to approximately 20 AU (Fig. 6).

For a comparison, a typical RMS sensitivity of SPIRE at 350 m is only ~2 mJy66. The ALMA

sensitivity is 0.64 mJy in 1 hour at 350 m and 0.84 mJy in 6 min at 450 m, assuming the same elevation and observing conditions67. CCAT is therefore a factor of ~2.5 more sensitive than

ALMA for TNO thermal flux measurements at 350 m (at 450 m confusion becomes the limiting factor). After the Herschel mission is complete in early 2013, JWST will provide continuum measurements of TNOs on the Wien side of the SED peak. However, submillimeter fluxes and light curves are needed to constrain the thermal and albedo models of the dwarf planet surfaces, going beyond simple gray-body fits. Therefore, there is a compelling science case for

CCAT 350 m (and longer-wavelength ALMA) follow-ups of TNOs detected at optical wavelengths (e.g., with the LSST).

                                                            61 Trujillo, C.A., et al. AJ, 122, 457 (2001) 62 Stansberry, J., et al. in The Solar System Beyond Neptune, ed. M. A. Barucci, et al. UIP, Tucson, 161 (2008)  63 Müller, T.G., et al., EMP, 105, 209 (2009) 64 Müller, T.G., et al., A&A, 518, L146 (2010) 65 This confusion limit will be recast into the extragalactic confusion noise for consistency with the rest of this section. 66 Lim, T.L., et al. A&A, 518, L148 (2010) 67 http://almascience.eso.org/call‐for‐proposals/sensitivity‐calculator 

 

 

3.5.1.2 T

CCAT wnew cons

Msiafcocasiev

Mstthinroco

Mona ap

Mdal

                  68 Stansber

TNO Goals fo

will enable rastraints on S

Measure TNOize distributffected by composition, arbon) and hize, albedo, volutionary

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                       rry, J., et al., in

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‐ 30 ‐ 

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‐ 31 ‐  

due to shape changes will be correlated with the optical light curve. If the optical variability is caused primarily by the albedo variations, the two light curves will be anti-correlated.

3.5.1.3 TNO Requirements

To achieve the TNO science goals outlined in the previous sections, the following requirements have to be satisfied:

Diffraction-Limited Beam Size and Aperture69: A 25 m aperture with diffraction-limited beams at 350 µm is required to provide a 0.25 mJy confusion limit (at 15 beams per source).

Sensitivity: A continuum sensitivity of 0.25 mJy in 1 hour at 350 m is required to detect a 100 km diameter TNO out to 20 AU or a 1,000 km TNO out to 100 AU.

Flux Calibration (Goal): Relative flux calibration of 2% (1) is required to detect 20%

variability in the thermal light curve at 10 level.

3.5.2 Comets

3.5.2.1 Comet Background

Together with TNOs, comets are among the most primitive bodies left over from the planetesimal building stage of the solar nebula, and so their physical and chemical composition provides an important link between nebular and interstellar (or outer disk) processes. Here, submillimeter observations are extremely powerful, having provided the most accurate estimates of the composition of the nucleus of comets and the isotopic ratios in key species (water, HCN70). The most complex molecules seen in comets (methyl cyanide, methyl formate, ethylene glycol) have also been detected at (sub)millimeter wavelengths71. Heterodyne spectroscopy can be used to probe of the composition of the nucleus itself, particularly if the observations are combined with simultaneous infrared observations of species without permanent dipole moments, such as CO2, C2H2, and CH4.

Deuterium fractionation in interstellar water, as compared to that measured in Solar System objects (in particular in comets), is of great interest. The CSO provided the first ground-based observation of HDO in a comet and the corresponding measurement of the D/H ratio in cometary water72,73. Since then, a handful of additional measurements have been obtained from the

                                                            69 This confusion limit should be recast in terms of the confusion noise for consistency with the rest of the document. 70 Bockelée‐Morvan, D., et al., A&A, 353, 1101 (2000); Bockelée‐Morvan, D., et al. In Comets II, eds. M. C. Festou, et al. Univ. Arizona Press, 391 (2005), see also recent review by Mumma, M. J., & Charnley, S. B. ARAA, 49, 471 (2011) 71 Crovisier, J., et al., A&A, 418, L35 (2004)  72 Lis, D.C., et al., Icarus, 130, 355 (1997) 73 Bockelée‐Morvan, D., et al.,  Icarus, 133, 147 (1998) 

 

 

ground—

derived Dsolar valuThis sugthe D/H r

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                  74 Drake, M75 See discu76 Gomes, 77 Morbide

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                       M.J., & Righter,ussion in MumR., et al., Natuelli, A., et al., N

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                    K., Nature, 41

mma, M.J., & Chre, 435, 466 (2ature, 435, 46

ort-cloud co

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16, 39 (2002) harnley, S.B., A2005) 2 (2005) 

‐ 32 ‐ 

omet—at ra

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‐ 33 ‐  

during the Late Heavy Bombardment and could have contributed a portion of Earth’s oceans78. Recent dynamical models79 also suggest that up to 90% of comets in the Oort cloud may have been captured from other stars in Sun’s birth cluster. All these calculations predict that significant variations in the cometary D/H ratio are expected and demonstrate the increasing emphasis on classifying comets based of their composition and isotopic ratios rather that orbital dynamics80.

The high cometary D/H ratios, twice the Standard Mean Ocean Water (SMOW) value, measured in Oort cloud comets appeared to have ruled out the delivery of the majority of terrestrial water from cometary sources – given the temperature and density gradients in the protosolar disk, Jupiter-family comets were expected to have even higher than Oort cloud comets. Therefore the recent Herschel/HIFI measurement of the D/H ratio in comet Hartley 2 came as a great surprise81. Hartley 2 is the first Jupiter-family comet, coming from the Kuiper belt, in which the

D/H ratio has been measured. The value implied by the HIFI measurements, (1.610.24)10-4, is a factor of 2 lower than the earlier measurements in Oort cloud comets and agrees with the

Vienna Standard Mean Water (VSMOW) ratio of 1.5610-4 in the Earth oceans. Hartley 2 potentially traces a different, large reservoir of water ice rich material in the outer Solar system – the Kuiper belt. The HIFI observations demonstrate that the earlier high D/H values are not representative for all comets—this is consistent with predictions of latest models of radial transport of pre-cometary material from the inner solar system outward to the comet-forming region. Consequently, a much high fraction – possibly all – of Earth’s ocean water could have been delivered by comets, which also might have seeded the early Earth with organics.

3.5.2.2 Comets Goals for CCAT

The question of isotopic gradients in the protosolar nebula, in the 5–40 AU range, is a key issue when considering cometary delivery of water to terrestrial planes. The HIFI observations of comet Hartley 2 clearly demonstrate the need for increasing the sample of comets with accurate measurements of the D/H ratio, in particular those coming from the Kuiper belt. CCAT offers outstanding capabilities to carry out such measurements by observing the 110–101 rotational transition of HDO at 509 GHz (the same line that has been observed by HIFI in comet Hartley 2). In a typical comet at 1 AU, this line is 5 times stronger than the 464 GHz transition accessible to ALMA. Atmospheric transmission models predict ~55% zenith transmission at 509 GHz at the CCAT site in the best 25 percentile weather conditions. With a 65 K (DSB) receiver temperature, the expected system temperature at an airmass of 1.4 is ~ 1,000 K, about the same as ALMA at 464 GHz. The strength of the 509 GHz HDO line thus more than compensates for

                                                            78 Gomes, R., et al., Nature, 435, 466 (2005) 79 Levison, H.F., et al., Science, 329, 187 (2010) 80Mumma, M.J., & Charnley, S.B. ARAA, 49, 471 (2011)  81 Hartogh, P., et al., Nature, 478, 218 (2011)

 

 

‐ 34 ‐  

the larger ALMA collecting area compared to CCAT. Assuming a 75% beam efficiency, the 509 GHz HDO line would have been detected by CCAT in comet Hartley (at 0.2 AU) with a signal-to-noise ratio of 5 for the integrated line intensity in a 4-hour integration. This clearly demonstrates that CCAT will be capable of detecting HDO in weak Jupiter-family comets, bringing significant contributions to our understanding of deuteration in different Solar System reservoirs. For brighter Oort cloud comets, HDO will be detectable at much larger geocentric distances and a heterodyne array will allow for mapping the structure of the coma in comets making a close approach to the Earth.

In order to determine the D/H ratio, an accurate measurement of the water production rate is needed. The submillimeter rotational transitions of water isotopologues are not observable from

the ground. Therefore, infrared (3–5 m) observations of water or OH prompt emission will have

to be relied on, as well as radio observations of the 18 cm OH line. In addition, the 118 m ground-state rotational transition of OH will soon be accessible from SOFIA, using the SHASTA or GREAT heterodyne instruments.

3.5.2.2 Comets Requirements

To carry out measurements of the D/H ratio in comets described above the following requirements have to be satisfied:

Frequency coverage: A heterodyne instrument covering the 110–101 rotational transition of HDO at 509.3 GHz is required.

Backend: A minimum velocity resolution of 0.1 km s-1 is required to spectrally resolve the narrow cometary lines (typical outflow velocities of order 1 km s-1 at 1 AU).

Sensitivity: A goal for the receiver temperature is ~65 K (DSB) at 509 GHz, resulting in a ~ 1,000 K (SSB) system temperature on the sky in the best 25 percentile weather conditions on the CCAT site. This will ensure detection of HDO in Jupiter-family comets with a water production rate similar to Hartley 2.

Flux Calibration: Absolute flux calibration of 5% is a goal, matching the modeling uncertainties for the HDO production rate.

4.0 DETAILED TELESCOPE REQUIREMENTS

Table 9 summarizes the CCAT telescope and instrument requirements set by the science requirements. Requirements and goals are listed separately, along with relevant notes and the primary science driver for each requirement. The definitions of the parameters follow.

 

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Table 9—CCAT Telescope Requirements

Parameter  Requirement  Goal Notes Primary Driver

Wavelength range  0.35 ‐ 2.1 mm  0.2 ‐ 3.5 mm Submm  galaxies,  star formation  and molecular cloud evolution, SZ 

Aperture  25 m  Same 0.45 mm continuum sensitivity as ALMA  

Resolves >80% of submm background radiation 

Field of view  20'  1  Limited by field curvature; not to be mechanically limited 

Surveys 

Emissivity   10% for 350 m <  < 850 m and   5% for  > 1 

mm 

Small c.f. atmosphere

   < 20% @  = 200 m   

Polarization  None  0.1% reconstructed On axis, after calibration  Molecular cloud B‐fields

Half Wavefront Error  < 12.5 m RMS  <9.5 µm RMS <1.5x integration time  

      c.f. no wavefront error  

Azimuth range  ‐180to +360  Centered on AZ = 90  Avoids cable wrap

Elevation range  20 ‐ 90  Maximum sky coverage

Zenith blind spot  < 2 zenith angle  < 1 zenith angle  Minimize AZ slews

Slew speed  > 3 s‐1  < 2 minutes for AZ slew

Slew acceleration  > 2 s‐2  1.5 s to max speed

Tracking speed  7' s‐1  Sidereal rate/sin(2)   

Blind pointing  2 RMS  0.5 RMS  <0.5 FWHM at 0.35 mm  To  arrive  on‐source  for local pointing after a slew 

Offset pointing  0.35 x (µm) RMS  0.1 beams within 1  Hourly pointing

Pointing stability  0.35 x (µm) RMS hr

‐1 

0.1 beams per hour Hourly pointing

Maximum scan speed  (0.33 s‐1) x (µm) in Az, half that in El 

200 Hz timestreams at all wavelengths 

Modulate  above  the atmospheric 1/f noise 

Maximum scan acceleration 

(0.33 s‐2) x (µm)  2 s turnaround time

Following error during scan 

< 1.8 x ( / 350 µm) RMS  Half FWHM beamwidth  

Pointing knowledge during scan 

 x ( / 350 µm) RMS  0.1 beams  

Pointing reconstruction   0.2 beam FWHM  0.1 beam FWHM  N.B. Involves instruments & data pipeline  too 

Counterpart identification  at  other wavelengths,  minimize photometric errors 

Nodding  1 in 2 s, 1 s settling  1 FoV in 2 s  

Secondary nutation  5' p‐p at 1 Hz, 0.1 s settling; 

Not in baseline design  

   None if MOS w/ beams separated by > 5' 

Chopping not necessary if there is MOS capability 

Must  be  able  to difference  off  of  nearby galaxies 

Mapping patterns  Arbitrary scan patterns  (1) Raster in Az/El, RA/Dec, Galactic coordinates, and arbitrary position angles, (2) Lissajous, (3) Spiral, (4) Box 

Optimized  field  shapes for  efficient  observing; toolkit  for  planning observations  prior  to runs 

Telescope/Dome Min. Clearance 

Under investigation   Under investigation  

 

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Requirements and their Connections to Science Goals

Wavelength Range: The wavelength range requirement is the range through which the observatory will perform routine scientific operations. The short-wavelength requirement (350 µm) is primarily driven by the requirement to measure the bolometric luminosities of galaxies, to use the submillimeter colors to identify high redshift galaxies, to avoid source confusion from galaxies, and for observations of Galactic molecular cloud cores and nearby galaxies near the peak of the thermal emission. The long-wavelength requirement (2.1 mm) is primarily driven by the spectrum (decrement) of the SZ effect for galaxy clusters. The 200 µm goal is driven by: (1) the desire to measure the dust continuum near its flux density peak for both nearby galaxies and star formation regions within the Milky-Way to best determine source luminosity and dust temperature, and (2) to observe high-J CO rotational line emission in AGN and in Galactic star formation regions, and to observe the [NII] 205 µm line from nearby galaxies.

Aperture: The aperture requirement is 25 m. This is driven by several science requirements

needing the 3.5 angular resolution at 350 m and NEFD of 14 mJy s1/2 at 350 m and 450 m.

Field of View: The FoV should not be limited by mechanical blockage. The 20' FoV

requirement is to maximize mapping speed. The goal of 1 is driven by the expectation that detector array formats will continue to grow over the lifetime of CCAT.

Emissivity: The telescope emissivity requirement (including the primary mirror, secondary

mirror, tertiary mirror, and secondary mirror support structure) is 10% for 350 µm to 800 m

and 5% for 1 mm to 3 mm.

Polarization: There is no polarization requirement; there is a polarization goal of < 0.1% on-axis, after post-observation corrections. This applies to the telescope and a polarimeter in combination. The key goal is to have stable (repeatable) systematic polarization, which can be measured and removed. This will enable B-fields angles projected onto the sky to be measured robustly with continuum observations (polarizations range from 0.1% to several percent, peaking at 2%).

Half-Wave-Front Error: The overall HWFE error budget for the telescope system is 12.5 µm rms. This HWFE limit is required to ensure that the integration time to a given flux density noise level does not grow by more than a factor of 1.5 over a perfect optical system at the shortest high-redshift survey wavelength, 350 µm. This increase in integration time would increase the overall time required by a survey by a factor of 1.5.

Azimuthal Range: The azimuthal range for the telescope requirement is -180 to +360, centered

on +90. This range will enable tracking of any source from rising, through transit, to setting without unwinding cables.

 

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Elevation Range: The 90 peak elevation range requirement is driven by the need to approach

the zenith for sources that transit near the zenith. The 20 minimum is driven by the desire to access all of the southern sky, and as much of the northern sky as is reasonable (above 3

airmasses). The 20 lower elevation limit also ensures a good stretch in airmass for elevation tips (skydips) of the telescope designed to measure opacity.

Zenith Blind Spot: The zenith blind spot should be as small as practically possible, with a

reasonable goal of 2 degrees (the requirement to track sources to within 2 of the zenith results in an 8 minute loss in integration time for sources at the declination equal to the telescope latitude).

Slew Speed: The Az requirement of 3 s-1 slew speed ensures that the entire sky can be accessed from any other given position on the sky in less than 2 minutes. Only half that speed is needed in El to meet the same requirement.

Slew Acceleration: The requirement of 2 s-2 acceleration on slews yields full slew speed within 1.5 seconds, minimizing time lost in slewing to, for example, sky reference positions during deep integrations.

Tracking Speed: The tracking speed requirement is 7' s-1. This rate is sufficient to track on a

source when it approaches within 2 of the zenith, where the AZ rate is 15 s-1 / sin(2).

Blind Pointing: The blind pointing requirement is 2 RMS from any region on the sky to any other region. This requirement ensures that a point source anywhere on the sky will be within

half beam power at 350 µm following a slew. The goal is 0.5 RMS, which is nearly sufficient to couple all of the power from a point source into a diffraction limited single-beam spectrometer at 350 µm following a slew.

Offset Pointing: The 1 offset pointing requirement, 0.35 ( / 350 µm) RMS is primarily driven by spectroscopy and photometry of point sources. To maximize sensitivity, a point source should be centered and within 1/10th of a diffraction limited beam (RMS).

Pointing Stability: Once 1/10th beam offset pointing is achieved, this should be stable for a 1 hour integration on source, so that checking pointing offsets needs to be done no more frequently than once per hour.

Scan Speed: To observe with cameras, the telescope must scan the sky at a speed of 0.33 s-1

µm) in Az, and half this value in El (for scanning at El up to 60, Az scan rate 2 El scan rate). This ensures that pixels in the array move more than a beamwidth in a timescale short compared with fluctuations in atmospheric transmission, thereby modulating the signal faster than atmospheric fluctuations.

Scan Acceleration: The scan acceleration is to be able to switch directions at the end of a scan within 2 seconds to maximize observing efficiency.

 

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Following Error During Scan: The difference between the commanded pointing position and the actual pointing position must be less than ½ beams at 350 µm and longer wavelengths.

Pointing Knowledge During Scan: The pointing knowledge (where the telescope was at any given time) during a scan must be better than 1/10th beam to enable diffraction-limited images to be reconstructed from data streams. This pointing knowledge refers to relative pointing between the telescope beam position and a nearby pointing source, where “nearby” refers to pointing sources used for offset pointing.

Reconstructed Pointing: The absolute pointing knowledge must be better than 1/5th beam after processing of images. This will ensure good registrations with images at other wavelengths for source cross-identification. The goal is 1/10th beam.

Nodding: The telescope should be capable of nodding from one position (e.g., source) to another

position (e.g., reference) 1 away in 2 seconds, with a settling time of 1 second. This ensures

high (90%) efficiency nod-cycles with integration times as short as 30 seconds. The choice of 1 fits well with Galactic science and Local Group galaxies, and is driven by the desire to move to

an off position for the 1 FoV cameras that are envisioned for the future.

Secondary Nutation: There is no requirement for secondary nutation. Nutation is often the method used to overcome short time-scale fluctuations in sky transparency, but it is not required for fast-scanning cameras or for multi-object spectrometers that can have one feed placed on “blank sky” off positions to enable rapid sky subtraction.

Mapping Patterns: The telescope must be able to perform arbitrary scan patterns, including rasters in Az/El, RA/Dec. (e.g., for extended objects and uniform map sensitivities), arbitrary position angles (e.g., for edge-on, nearby galaxies or the Galactic plane), Lissajous (e.g., to eliminate turn-around losses for small fields), Spiral (e.g., to control sampling or for pointing), and Box patterns. These different patterns optimize field shapes for efficient observing for the various types of science envisioned with CCAT. The maximum scan amplitude requirement is

10.

Telescope/Dome Minimum Clearance: The instrument beams must not intersect the dome slit at a substantial level, which would introduce background loading modulation, degrading sensitivity. This will be avoided by ensuring adequate clearance between the telescope beam and the dome slit. This is under investigation.

 

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5.0 DETAILED INSTRUMENT REQUIREMENTS

Table 10 lists the instrument requirements. Additional spectrometer backend goals and requirements are listed in section 3.1.4.3.

Table 10—CCAT Basic Instrument Requirements

Instrument  Bandpasses  Format Spectral Resolution 

Flux Calibration Uncertainty 

(Goal)  

Sensitivity (Goal) 

     SWCam  Matched to 200a, 350, 

450, & 620a m atmospheric windows 

3 sub‐field cameras ×1282 detectors, 6.5’ FOV at λ = 350 μm 

N/A 10%  NEFD 14 mJy s1/2 at 350 & 

450 m 

LWCam  Matched to 750a, 850, 

1100, 1300, & 2100 m 

542 detectors, 15’FOV at λ = 850 μm 

N/A 10%  BLIP

D‐Spec  TBD  TBD TBD TBD  TBD 

H‐Spec high priority bands 

462 – 510 GHz, 600 – 720 GHz, 787 – 960 GHz ; spectrometer must process >1.8 GHz IF bandwidth 

> 100 channels, > 50 pixels (dual polarization) 

0.1 km s‐1 20%  TSYS:  1000 K, 1200 K, 2200 K, respectively  

H‐Spec low priority band 

272 – 373 GHz; spectrometer must process >1.8 GHz IF bandwidth 

> 100 channels, > 50 pixels (dual polarization) 

0.1 km s‐1 20%  TSYS:  200 K

aBands noted with this symbol are goal bands, not required bands.

 

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APPENDIX A: ACRONYMS AND ABBREVIATIONS

ACT – Atacama Cosmology Telescope

AGN – active galactic nuclei

Az – azimuth

CMB – Cosmic microwave background (radiation)

COBE – COsmic Background Explorer

CFIRB – cosmic far-infrared background radiation

CFRP – carbon-fiber-reinforced-plastic

El – elevation

FoV – field-of-view

FWHM – full width half maximum

GMC – giant molecular cloud

HWFE – half wavefront error

IMF – (Stellar) Initial mass function

ICM – intracluster medium

ISM – interstellar medium or interstellar media

kSZ – kinetic Sunyaev-Zel’dovich (effect)

LWCam – Long-Wavelength Camera

MOS – multiobject spectroscopy or multiobject spectrometer, including integral-field spectrometer

NEFD – noise equivalent flux density

NIR – near-infrared

PDR – Photon-dominated regions

PWV – precipitable water vapor

RC – Ritchey Chretien

 

‐ 41 ‐  

RMS – root mean square

rSZ – relativistic Sunyaev-Zel’dovich effect

SED – spectral energy distribution

SFR – star-formation rate

SPT – South Pole Telescope

SWCam – Short-Wavelength Camera

SZ – Sunyaev-Zel’dovich

tSZ – thermal Sunyaev-Zel’dovich (effect)

TNO – Trans Neptunian Object

ULIRG – ultraluminous infrared galaxy

UV – ultraviolet

XDR – X-ray dissociation region

CDM -- cold dark matter

 

‐ 42 ‐  

APPENDIX B: DOCUMENT HISTORY

July, 2011: Document commissioned

September 27, 2011: First draft compiled for comments – instrument requirements table not yet included

December 6, 2011: First complete draft with comments addressed in instrument requirements included

January 17, 2012: Included minor, editorial comments and updated the heterodyne array section with new sensitivities and integration times calculated by Rene Plume.

January 20, 2012: Document released