The Science Case for Building a Band 1 Receiver Suite for ALMA
J. Di Francesco, D. Johnstone, B. Matthews, N. Bartel, L. Bronfman, S. Casassus, S. Chitsazzadeh, M. Cunningham, G. Duchene, A. Hales, M. Houde, D. Iono, P. M. Koch, R. Kothes, S.-P. Lai, S.-Y. Liu, B. Mason, T. Maccarone, G. Schieven, A. M. M. Scaife, D. Scott, H. Shang, S. Takakuwa, J. Wagg, A. Wootten, F. Yusef-Zadeh
VVersion: November 11, 2018
The Science Cases for Buildinga Band 1 Receiver Suite for ALMA
J. Di Francesco , , D. Johnstone , , B. Matthews , , N. Bartel , L. Bronfman ,S. Casassus , S. Chitsazzadeh , , M. Cunningham , G. Duchˆene , , J. Geisbuesch ,A. Hales , P.T.P. Ho M. Houde , D. Iono , F. Kemper , P.M. Koch , K. Kohno ,R. Kothes , S-P. Lai , K.Y. Lin , S.-Y. Liu , B. Mason , T.J. Maccarone ,N. Mizuno , O. Morata , G. Schieven , A.M.M. Scaife , D. Scott , H. Shang ,S. Takakuwa , J. Wagg , , A. Wootten , F. Yusef-Zadeh National Research Council Canada, 5071 West Saanich Rd, Victoria, BC, V9E 2E7, Canada Dept. of Physics & Astronomy, University of Victoria, Victoria, BC, V8P 1A1, Canada Dept. of Physics and Astronomy, York University, Toronto, M3J 1P3, ON, Canada Dept. de Astronom´ıa, Universidad de Chile, Casilla 36-D, Santiago, Chile Dept. of Physics and Astronomy, The University of Western Ontario, London, ON, N6A 3K7, Canada School of Physics, University of New South Wales, Sydney, NSW 20152, Australia Astronomy Dept., University of California, Berkeley, CA 94720-3411, USA Universit´e Joseph Fourier - Grenoble 1/CNRS, LAOG UMR 5571, BP 53, 38041 Grenoble, France National Research Council Canada, P.O. Box 248, Penticton, BC, V2A 6J9, Canada National Radio Astronomy Observatory, 520 Edgemont Road, Charlottesville, Virginia 22903, USA Academia Sinica, Institute of Astronomy and Astrophysics, P.O. Box 23-141, Taipei 10617, Taiwan National Astronomical Observatory of Japan, 2-21-1 Osawa, Mitaka, Tokyo, 181-8588, Japan Institute of Astronomy, The University of Tokyo, 2-21-1 Osawa, Mitaka,Tokyo 181-0015, Japan Institute of Astronomy and Dept. of Physics, National Tsing Hua University, Taiwan School of Physics and Astronomy, University of Southampton, Southampton, Hampshire, S017 1BJ, UK Dept. of Physics and Astronomy, University of British Columbia, Vancouver, BC, V6T 1Z1, Canada European Southern Observatory, Alonso de Cordova 3107, Vitacura, Casilla 19001, Santiago 19, Chile Astrophysics Group, Cavendish Laboratory, JJ Thomson Avenue, Cambridge, CB30HE, UK Dept. of Physics and Astronomy and Center for Interdisciplinary Research in Astronomy, NorthwesternUniversity, Evanston, IL 60208, USA a r X i v : . [ a s t r o - ph . I M ] O c t
1. Executive Summary
We present a set of compelling science cases for the ALMA Band 1 receiver suite. Forthese cases, we assume in tandem the updated nominal Band 1 frequency range of 35-50 GHzwith a likely extension up to 52 GHz; together these frequencies optimize the Band 1 sciencereturn. The scope of the science cases ranges from nearby stars to the re- ionization edge ofthe Universe. Two cases provide additional leverage on the present ALMA Level One ScienceGoals and are seen as particularly powerful motivations for building the Band 1 Receiversuite: (1) detailing the evolution of grains in protoplanetary disks, as a complement to the gaskinematics, requires continuum observations out to ∼
35 GHz ( ∼ < z <
10, also requires Band 1 receiver coverage. The range of Band 1 scienceis wide, however, and includes studies of very small dust grains in the ISM, pulsar windnebulae, radio supernovae, X-ray binaries, the Galactic Center (i.e., Sgr A*), dense cloudcores, complex carbon-chain molecules, masers, magnetic fields in the dense ISM, jets andoutflows from young stars, distant galaxies, and galaxy clusters (i.e., the Sunyaev-Zel’dovichEffect). A comparison of ALMA and the Jansky VLA (JVLA) at the same frequencies ofBand 1 finds similar sensitivity performance at 40–50 GHz, with a slight edge for ALMA athigher frequencies (e.g.,., within a factor of 2 for continuum observations). With its largernumber of instantaneous baselines, however, ALMA Band 1data will have greater fidelitythan those from the JVLA at similar frequencies. 3 –
2. Introduction
The Atacama Large Millimeter/submillimeter Array (ALMA) will be a single researchinstrument composed of at least fifty-four 12-m and twelve 7-m high-precision antennas,located at a very high altitude of 5000 m on the Chajnantor plain of the Chilean Andes.The weather conditions at the ALMA site will allow transformational research into thephysics of the cold Universe across a wide range of wavelengths, from radio to submillimeter.Thus, ALMA will be capable of probing the first stars and galaxies and directly imagingthe disks in which planets are formed. ALMA will be a complete astronomical imaging andspectroscopic instrument for the millimeter/submillimeter regime, providing scientists withcapabilities and wavelength coverage that complement those of other research facilities of itsera, such as the Jansky Very Large Array (JVLA), James Webb Space Telescope (JWST),30-m class Giant Segmented Mirror Telescopes (GSMTs), and the Square Kilometer Array(SKA). ALMA will be the pre-eminent platform for astronomical research at millimeter andsubmillimeter wavelengths for decades to come.ALMA will revolutionize many areas of astronomy and an amazing breadth of sciencehas already been proposed (see for example the ALMA Design Reference Science Plan). Thetechnical requirements of the ALMA Project are, however, driven by three specific LevelOne Science Goals: (1)
The ability to detect spectral line emission from CO or CII in a normal galaxy like theMilky Way at a redshift of z = 3, in less than 24 hours of observation. (2) The ability to image the gas kinematics in a solar-mass protostellar/ protoplanetary diskat a distance of 150 pc (roughly, the distance of the star-forming clouds in Ophiuchus orCorona Australis), enabling one to study the physical, chemical, and magnetic field structureof the disk and to detect the tidal gaps created by planets undergoing formation. (3)
The ability to provide precise images at an angular resolution of 0.1 (cid:48)(cid:48) . Here the term“precise image” means an accurate representation of the sky brightness at all points wherethe brightness is greater than 0.1% of the peak image brightness. This requirement appliesto all sources visible to ALMA that transit at an elevation greater than 20 ◦ .ALMA was originally envisioned to provide access to all frequencies between 31 GHzand 950 GHz accessible from the ground. During a re-baselining exercise undertaken in2001, the entire project was scrutinized to find necessary cost savings. The two lowestreceiver frequencies, Bands 1 and 2, covering 31–45 GHz and 67–90 GHz respectively, wereamong those items delayed beyond the start of science operations. Nevertheless, Band 1 wasre-affirmed as a high priority future item for ALMA.In May 2001, John Richer and Geoff Blake prepared a document Science with Band 1 (31–45 GHz) on ALMA as part of the re-baselining exercise. Key arguments for maintainingthe receiver were put forward and included: (1) exciting science opportunities, bringing ina wider community of users; (2) a significantly faster imaging and survey instrument thanthe upgraded VLA (now known as the Jansky VLA or JVLA), especially due to the largerprimary beam; (3) access to the southern sky at these wavelengths; (4) excellent sciencepossible even in “poor” weather; (5) a relatively cheap and reliable receiver to build andmaintain; and (6) a very useful engineering/debugging tool for the entire array given thelower contribution of the atmosphere at many of its frequencies relative to other bands.The Richer/Blake document was followed by an ASAC Committee Report in October2001 in which the addition of Japan into the ALMA project re-opened the question ofobserving frequency priorities for those receiver bands which had been put on hold duringre-baselining. The unanimous recommendation of the ASAC was to put Band 10 as toppriority, followed by a very high priority Band 1. The key science cases for Band 1, at thattime, were seen to be (1) high-resolution Sunyaev-Zel’dovich effect (SZE) imaging of clustergas at all redshifts; and (2) mapping the cold ISM in Galaxies at intermediate and highredshift.The scientific landscape has changed significantly since 2001 and thus it is time toupdate the main science drivers for Band 1, even reconsidering its nominal frequency rangeto optimize its science return. In addition, the ALMA Development process has begun, andnow is the time to put forth the best case for longer wavelength observing with ALMA. InOctober 2008, two dozen astronomers from around the globe met in Victoria, Canada todiscuss Band 1 science. This paper represents a brief summary of the outstanding casesmade possible with Band 1 that were highlighted at that meeting and since. In Section 3, wedescribe the new nominal Band 1 frequency range of 35-50 GHz, and its possible extensionto 52 GHz. In Section 4, we present two science cases that reaffirm and enhance the alreadyestablished ALMA Project Level One Science Goals. In Section 5, we provide a selection ofscience cases that reinforce the breadth and versatility of the Band 1 Receiver. Section 6discusses both weather considerations at the ALMA site and the complementarity of ALMAversus the EVLA at Band 1 frequencies. Finally, Section 7 briefly summarizes the report.
3. The Band 1 Frequency Range
Band 1 was originally defined as 31.3–45 GHz. The lower value was the set at thelower edge of a frequency range assigned to radio astronomy and the upper value was set toinclude SiO J =1–0 emission at 43 GHz. Receiver technology advances, however, have made itpossible to widen and shift the Band 1 range and optimize science return. For example, a shift 5 –to higher frequencies for Band 1 will improve (slightly) the angular resolution of continuumobservations and better exploit the advantages of the dry ALMA site. Furthermore, a widerrange and shift to higher frequencies will allow molecular emission from galaxies at a widerrange of (slightly lower) redshifts to be explored, and allow molecular emission from severalnew species in our Galaxy to be probed. (Of course, this shift does in turn remove the abilityto detect molecular emission from some higher redshift galaxies and some other Galactictransitions.)A review of the nominal frequency range by the Band 1 Science Team (i.e., severalauthors of this document) in June 2012 resulted in a proposed new Band 1 frequency rangedefinition, nominally 35–50 GHz with an extension up to 52 GHz encouraged. The shift up to50 GHz will allow the important line CS J =1–0 at 48.99 GHz to be observable with ALMA.In addition, the nominal range of 35-50 GHz alone is itself ∼
10% wider than before. As itwill provide the highest sensitivities, the nominal range will be preferred for high-redshiftscience. The extension to 50–52 GHz, which the JVLA cannot observe, may be somewhatadversely affected by atmospheric O , resulting in lower relative sensitivity. Since numeroustransitions of other interesting molecules have rest frequencies at 50-52 GHz, however, thisextension will allow such emission to be probed toward sources in our Galaxy. This documenthas been updated in September 2012 to reflect the new nominal frequency range and theextension. See Section 6 for a comparison of the sensitivities and imaging characteristics ofALMA and the JVLA over their frequencies in common.
4. Level One Science Cases for Band 1
In this section, we present two science cases that reaffirm and enhance the alreadyestablished ALMA Project Level One Science Goals: Evolution of Grains in Disks AroundStars (4.1) and The First Generation of Galaxies (4.2).
Planet formation takes place in disks of dust and gas surrounding young stars. Itis within these gas-rich protoplanetary disks that dust grains must agglomerate from thesub-micron sizes associated with the interstellar medium to larger pebbles, rocks and plan-etesimals, if planets are ultimately to be formed. The timescale of this agglomeration processis thought to be a few tens of Myr for terrestrial planets, while the process leading to the 6 –formation of giant planet cores remains uncertain. Core accretion models require at leasta few Myr to form Jovian planets (Pollack et al. 1996), while dynamical instability modelscould form giant planets on orbital timescales ( t (cid:28) where and when dust coagulation occurs are critical to constrain currentmodels of planetary formation. Growth from sub-micron to micron-sized particles can betraced with infrared spectroscopy and imaging polarimetry. The next step, growth beyondmicron sizes, is readily studied by determining the slope of the spectral energy distribution(SED) of the dust thermal emission at submillimeter and millimeter wavelengths. The dustmass opacity index at wavelengths longer than 0.1 millimeter is approximately a power-lawwhose normalization depends on the dust properties, such as composition, size distribution,and geometry (Draine 2006). The index of the power law is β . The presence of large grainsis detectable through a decrease in the index β , which can be derived directly from the slopeof the Rayleigh-Jeans tail of the SED, α , where β = α −
2, when the emission is opticallythin. Studies reveal that the β values of disks are substantially lower than the typical ISMvalue of ∼ β occurs when the disk is not resolvedspatially. The amount of flux detected at a given wavelength is a function of both β and thesize of the disk (Testi et al. 2001). Resolving the ambiguity therefore is truly a matter ofresolution, and sufficient resolution is only offered at these wavelengths by interferometers.Among the three high level science goals of ALMA is the ability to detect and image gaskinematics in protoplanetary disks undergoing planetary formation at 150 pc. At ALMA’s 7 –observing wavelengths, its capability for imaging the continuum dust emission in these disksis also second-to-none. At present, however, ALMA will only reach a longest wavelength of3.6 mm. Given that dust particles emit very inefficiently at wavelengths longer than theirsizes, the present ALMA design will not be sensitive to particles larger than ∼ mm. Thissituation negates ALMA’s potential ability to follow the dust grain growth from mm-sized tocm-sized pebbles in protoplanetary disks. Figure 2 shows the SEDs for three different circumstellar disk models, computed usingthe full dust radiative transfer MCFOST code (Pinte et al. 2006; Pinte et al. 2009). Themodel parameters are representative of protoplanetary disks (although there is substantialobject-to-object variation). The circumstellar disk is passively heated by a 4000 K, 2 L (cid:12) central star and the system is located 160 pc away. The dust component of the disk isassumed to be fully mixed with the gas and the latter is assumed to be in vertical hydrostaticequilibrium. The disk extends radially from 1 AU to 100 AU. The total dust mass in themodel is 10 − M (cid:12) (the gas component is irrelevant for continuum emission calculations, soits mass is not set in the model, though a typical 100:1 gas:dust ratio is generally assumed).The dust population is described by a single power-law size distribution N ( a ) ∝ a − . witha minimum grain size of 0.03 µ m and extending to 10 µ m, 1 mm or 1 cm depending on themodel. The dust composition is taken to be the “astronomical silicates” from Draine (2003).Figure 2 reveals that observations in the ALMA Band 1 spectral region are crucial fordetermining whether grain-growth to cm-sizes is indeed occurring. The 1 cm flux density ofthe max size =1 cm disk model is ∼ µ Jy, comparable to the 1-sigma sensitivities provided byALMA’s Band 1 with 1 minute integration. Besides ALMA, there are no existing or plannedsouthern astronomical facilities capable of observing to such depths at these frequencies.Therefore, if Band 1 is not built there will be no way of putting ALMA observations ofprotoplanetary disks in the context of coagulation of dust grains to centimeter sizes .By complementing observations in other ALMA Bands, Band 1 will provide a cruciallonger wavelength lever to minimize the uncertainty in α . Evidence for small pebbles hasbeen detected in several disks (Rodmann et al. 2006). The prime example is TW Hya, aprotoplanetary disk 50 pc from the Sun (Wilner et al. 2000). Its SED is well matched byan irradiated accretion disk model fit from 10s of AU to an outer radius of 200 AU andrequires the presence of particle sizes up to 1 cm in the disk (see Figure 3). The measured β is 0 . ± . β has been detected with stellarluminosity, mass or age (Ricci et al. 2010). Lower α values are associated with less 60 µ mexcess, however, suggesting that settling or agglomeration processes could be removing thesmallest grains, decreasing the shorter wavelength emission (Acke et al. 2004).At the resolution provided by its longest baselines at ∼
40 GHz ( ∼ (cid:48)(cid:48) ), ALMA will 8 –easily resolve protoplanetary disks at the distance of the closest star-forming regions (50–150pc). These resolved images will provide the most accurate determination of the disk’s dustmass. The dust distribution at centimeter wavelengths can then be compared to millimeterand submillimeter images, revealing where in the disk dust coagulation is occurring. Forexample, previous investigations of the radial dependency of dust properties in disks byGuilloteau et al. (2009) and Isella et al. (2010) were conducted at 1 mm and 3 mm, andas such they were sensitive to only millimeter-sized grains. Note, however, that Melis et al.(2011) used the Jansky VLA to map the 7 mm emission from the protoplanetary disk aroundthe young source L1527 IRS at ∼ (cid:48)(cid:48) and tentatively detected a dearth of “pebble-sized”grains. ALMA Band 1 will help clarify this situation.As described above, Band 1 data will be sensitive to larger grains. Moreover, throughdetection of concentrations of such large grains, protoplanets in formation can be identified.These condensations are predicted by simulations of gravitational instability models (seeFigure 4a; Greaves et al. 2008) and have been detected in the nearby star HL Tau (Figure 4b;Greaves et al. 2008). ALMA’s high resolution is clearly critical to the detection of suchsubstructures within protoplanetary disks.Detecting dust emission at centimeter wavelengths also requires high sensitivity, becauseits brightness is several orders of magnitude lower than in the submillimeter. In addition,at wavelengths longer than 7 mm (i.e., ν <
45 GHz), the contribution from other radiativeprocesses, such as ionized winds, can contribute significantly to the total flux and complicatethe interpretation of detected emission. Rodmann et al. (2006) found that the contribution offree-free emission to the total flux is typically 25% at a wavelength of 7 mm. Observations ofcontinuum emission at the 35-50 GHz (9 mm–6 mm) spectral range enabled by Band 1 wouldincrease substantially the sampling rate in the region where emission is detected from boththe free-free and thermal dust emission components. The synergy with the JVLA will providea longer wavelength lever for sources observed in common, providing an estimate of the free-free contribution to the Band 1 flux. Such data would not be essential, however, given widefrequency coverage within Band 1 alone. For example, multiple continuum observationscould be used to quantify accurately the relative amounts of free-free and dust emissionthrough changes in spectral slope, and thereby determine precisely the contribution fromlarge dust grains (i.e., protoplanetary material).
Around main sequence stars, pebble-sized bodies are produced differently than in disksaround pre-main-sequence stars. Here, destructive, collisional cascades from even larger 9 –planetesimals through to centimeter, millimeter, and then micron-sized particles providesongoing replenishment of the debris population (Wyatt 2009; Dullemond & Dominik 2005).The methods of detection of large (i.e., centimeter-sized) grains is the same as in proto-planetary disks, despite their origin in destructive, rather than agglomerative, processes. Ineach case, the longer the wavelength at which continuum emission is detected, the larger thegrains that must be present in the system.Boley et al. (2012) detected the debris disk of Fomalhaut using Band 7, and noted itssharp inner and outer boundary. Band 1 images , however, could show higher contrast fea-tures in debris disks compared to other ALMA bands, due to the longer resonant lifetimesof the larger particles that dominate the emission. This sensitivity in turn will help detectany edges and gaps in the disks. Dramatic changes in the morphology of debris disks as afunction of wavelength have already been observed (e.g., Maness et al. 2008), but not yetat the long wavelengths Band 1 will probe. When observed, such structures are often con-sidered signposts to the existence of planets. These detections will be challenging comparedwith detecting forming condensations in protoplanetary disks. Debris disks typically haverelatively low surface brightnesses and large spatial distributions 100s of AU in radii. Theyalso can be found much closer to the Sun than protoplanetary disks. Indeed, the closestdisks could subtend as much as ∼ (cid:48)(cid:48) on the sky (assuming a 300 AU diameter disk at 2pc). Therefore, ALMA’s larger field of view than other long wavelength instruments, suchas the JVLA, will be very advantageous for imaging these objects, although mosaicking willstill be required to image the largest ones on the sky. ACA observations will be instrumentalin this research. The first generation of luminous objects in the Universe began the process of re-ionizingthe intergalactic medium (IGM). The detection of large-scale polarization in the cosmicmicrowave background (CMB), caused by Thomson scattering of the CMB by the IGMduring re-ionization, suggests that the Universe was significantly ionized as far back as z ≈ ± z (cid:38) α break is due to moderate amounts of neutral hydrogen in the IGM,suggesting re-ionization was complete by z ≈
6. The Gunn-Peterson effect also insures thatat these redshifts the Universe is opaque at wavelengths shorter than ∼ µ m. 10 –To study the first generations of galaxies, and to understand the origins of the black hole-bulge mass relation, it will be necessary to study the star-formation properties of galaxiesin the 6 (cid:46) z (cid:46)
11 range. Quasar hosts and other sources are rapidly being discovered atthe near end of this range (e.g., Cool et al. 2006; Mortlock et al. 2008; Glikman et al. 2008;Willott et al. 2009), and searches are underway for even more distant objects (e.g., Ota etal. 2008; Bouwens et al. 2009).Recently, CO has been detected in galaxies at redshifts > . These and other observa-tions in the cm/mm of z > z > FIR > L (cid:12) . Only a small fraction of galaxies are this luminous. Thebest-studied such object is J1148+5251 with a redshift of z = 6 .
419 (see Carilli et al. 2008).For example, Walter et al. (2004) imaged the CO J =3–2 emission (Figure 5) using the VLA,from which they were able to infer the dynamical mass. Walter et al. (2009) were not ableto detect the [NII] line at 205 µ m,, but did detect the CO J =6–5 transition. More recently,Wang et al. (2011b) detected the lower-energy CO J =2–1 transition and Reichers et al.(2009) imaged CO J =7–6 and CI ( P – P ) emission towards this source. These and other(dust continuum) observations show that there was already a significant abundance of metalsand dust by this epoch.Figure 6 shows the observable frequency of rotational transition sof CO, from J =1–0through J =10–9, as a function of redshift. Also shown are the frequency ranges of the ALMABands (excluding Band 2 for clarity). Note that this Figure shows the new nominal range ofBand 1 of 35-50 GHz, as this range will yield the highest sensitivities. As the Figure shows,Band 1 receivers will be able to detect galaxies in J =3–2 at 6 (cid:46) z (cid:46)
9, i.e., in the redshiftsof the era of re-ionization ( z > ∼ J lines thatmay be less excited. (For example, Band 3 receivers would be able to detect J =6–5 emissionin the range 4 . (cid:46) z (cid:46) . J =2–1and J =1–0 emission at 3 . (cid:46) z (cid:46) . . (cid:46) z (cid:46) .
3, respectively. Assuming a 150 µ JyCO J =2–1 line of width ∼
600 km s − at z = 5 .
7, a 5 σ detection would take less than 4hours with 50 ALMA antennas.Band 1 will also allow multiline observations toward certain subsets of redshifts. Forexample, galaxies at 1 . (cid:46) z (cid:46) . J =4–3 and J =3–2 Note that interferometers in general have an advantage over single-dish telescopes when detecting molec-ular emission at high redshift since their high-resolution imaging capabilities provide the spatial informationneeded to associate a detection with a specific object.
11 –in Band 4 (NB: a small gap exists at z ≈ P / – P / line can be observed toward a subset of these galaxies at 1 . (cid:46) z (cid:46) . . (cid:46) z (cid:46) . z ≈ z sources. An instantaneous ∼ The discovery of molecular gas in quasar host galaxies at z ∼
6, when the Universewas less than 1 Gyr old (Walter et al. 2003; Bertoldi et al. 2003; Carilli et al. 2007), hasopened a new window on the study of gas in systems that contributed to the reionizationof the Universe. Studies of how the molecular gas properties should evolve, and how theycan be used to reveal the dynamics of these massive systems, have recently prompted a newgeneration of semi-analytic models with the further aim of understanding how high-redshiftquasars fit within the context of large-scale structure formation. Li et al. (2007, 2008) haveused state of the art N-body simulations to show that the observed optical properties of high-redshift quasars can be explained if these objects formed in the most massive dark matterhalos ( ∼ × M (cid:12) ) early on. These models predict that the most luminous quasarsshould evolve due to an increase of major mergers, which one would expect to find evidencefor in the CO line profiles and the spatial distribution of the molecular gas (Narayananet al. 2008). Detailed radiative transfer models of the FIR spectral energy distributionof these systems have been driven by the observations of one z = 6 .
42 quasar (namelyJ1148+5251; Walter et al. 2003, 2004). Larger samples of CO-detected quasars are neededto provide better constraints on the models and constrain dynamical masses to compare withinfrared measurements of black-hole masses (e.g., from MgII lines) and explore the (possible)evolution of the relation between the masses of central black holes and bulges. Current 3 mmsurveys of high-J CO line emission in z ∼ α Emitters
The rarity of the luminous quasars at early times suggests that their UV emission isunlikely to contribute significantly to the reionization of the Universe (e.g., Fan et al. 2001).A more important type of galaxy in the context of cosmic re-ionization are the Lyman- α emitters (hereafter LAEs). These galaxies were discovered through their excess emission innarrow-band images centered on the redshifted Lyman- α line (e.g. Hu et al. 1998; Rhoadset al. 2000; Taniguchi et al. 2005), and constitute a significant fraction of the star-forminggalaxy population at z ∼
6. While the star-formation rates in LAEs inferred from theirUV continuum emission are a few tens of solar masses per year (e.g. Taniguchi et al. 2005),their number density and the shape of the Lyman- α emission line provide important probesof physical conditions in the Universe around the epoch of reionization. As such, it is offundamental importance that we understand the properties related to their star-formationactivity. In particular, we need to quantify the amount of molecular gas available for fuel.Wagg, Kanekar & Carilli (2009) used the Green Bank Telescope to search for CO J=1–0line emission in z > . z = 6 .
96 (Iye et al. 2006). The limits to the CO line luminosityimplied by the non-detections of CO J=1–0 in these two objects suggest modest molecular gasmasses ( < ∼ M (cid:12) ). This conclusion, however, is based on observations of only two objects,and future studies would benefit from the sensitivity gained by observing higher-order COtransitions, whose flux density may scale as ν due to a contribution to the molecular gasexcitation by the cosmic microwave background radiation (19 K at z = 6). With otherfacilities, it has been proven challenging to detect even the higher energy J =2–1 line fromLyman- α -emitting galaxies at these redshifts, using existing facilities (Wagg & Kanekar2011). At these redshifts, such studies would require ALMA, including the Band 1 receivers.
5. A Broad Range of Science Cases
Along with the two science cases presented above, there is a wealth of scientific oppor-tunity available to the ALMA community if the Band 1 Receiver is built. Here we highlighta selection of science cases which would significantly benefit from Band 1 on ALMA.
The astrophysical continuum radiation at wavelengths of ∼ The last decade has seen the discovery of surprisingly bright cm-wavelength radio emis-sion in a number of distinct galactic objects but most notably in dark clouds (e.g., Finkbeineret al. 2002; Casassus et al. 2008 (see Figure 7); Scaife et al. 2009a). The spectrum of this newcomponent of continuum radiation can be explained by electric dipole radiation from rapidlyrotating (“spinning”) very small dust grains (VSGs), as calculated by Draine & Lazarian(1998; DL98). It has also been seen as a large-scale foreground in CMB maps, spatiallycorrelated with thermal dust emission and having a spectrum peaking at ∼
40 GHz.All of the existing work aimed at diagnosing this continuum emission is derived fromCMB experiments on large angular scales, where the bulk of the radio signal occurs, e.g.,recently by the
Planck satellite. Details on small angular scales are crucial, however, forprobing star formation and circumstellar environments. Simply, progress in the understand-ing of the solid and gaseous states of the ISM requires sufficient resolution to separate thedistinct environments. Directly measuring the VSG abundance and solid state physics isvery exciting because VSGs play a central role in the chemical and thermal balance of theISM. For example, the smallest grains account for most of the surface area available forcatalysis of molecular formation.DL98 proposed that the grain size distribution in their spinning dust model would bedominated by VSGs, thought to be mostly PAH nanoparticles. The size distribution of VSGsis poorly known as studies of interstellar extinction are relatively insensitive to its details.The existence of VSGs is supported by assertions that the existence of a significant amountof carbonaceous nanoparticles in the ISM could explain observations of the unidentifiedIR emission features and the strong mid-infrared emission component seen by IRAS whichmust result from starlight reprocessing of ultrasmall grains. The fraction of the ISM carboncontent proposed to exist in VSGs considerably exceeds that implied by the MRN dust sizedistribution, which is known to underestimate it.Observationally determining PAH content in dust clouds is not straightforward. Wherethere is a strong source of UV flux present it is possible to identify PAHs by their spectralemission features. In the case of pre-stellar and Class 0 clouds, however, these features are 14 –absent. With microwave observations from ALMA Band 1 constraining the spinning dustSED at similar resolution to, e.g., Spitzer or the forthcoming MIRI instrument on the JWST,it will be possible to measure the size distribution of VSGs directly from the data. This workwill also be important in the context of circumstellar and protoplanetary disks, where theproposed population of VSGs may have important implications for disk evolution. Certainly,spinning dust emission will provide a better measure of the small grain population withincircumstellar disks than PAH emission since favorable excitation conditions for PAHs existonly in the outermost layers of the disk. Since all the VSGs in the disk should contributespinning dust emission, such emission will provide a much better probe of the mass in VSGs.Combining this information with the PAH emission features would then also give us a usefulmeasure of sedimentation in disks.Theoretically, spinning dust emission from a VSG population has been shown to dom-inate the thermal emission from stellar disks (around HAeBe stars) at significant factorsat frequencies ≤
50 GHz (Rafikov 2006). The existence of these VSGs has been confirmedobservationally using PAH spectral features as seen in the disks of Herbig Ae/Be stars (Acke& van den Ancker 2004) but it has not been detected in protoplanetary disks due to a lackof strong UV flux. Since spinning dust emission has been observed to be spatially corre-lated with PAH emission (Scaife et al. 2009b) spinning dust may provide a unique windowon the small grain population of these disks. In the context of disk evolution, these recentmeasurements conflict with the established view that dust grains are expected to grow asdisks age. It may be the case that dust fragmentation is important in disks (Dullemond& Dominik 2005), or there exists a separate population of very small carbonaceous grainsdistinct from the MRN distribution (Leger & Puget 1984; Draine & Anderson 1985). Thissecond proposition has not only important implications for the study of circumstellar disksbut also for the complete characterization of dust and the ISM more generally.The arcsecond resolution necessary for these measurements will be achievable with sev-eral ALMA configurations and Band 1. From the models of Rafikov (2006), the differencebetween a thermal dust spectrum with β ≈ σ in a matter of minutes with ALMA Band1. With longer observation times and consequently higher sensitivity, it will also be possibleto distinguish between different grain size distributions and physical conditions within thedisk (such as grain electric dipole moments, rotational kinematics, optical properties andcatalysis of molecule formation).Spinning dust emission provides a unique insight into the VSG population under con-ditions where it is not possible to observe using mid-IR emission. The high resolution andexcellent sensitivity of ALMA are ideal for differentiating the distinct environments where 15 –the VSG population resides and will be crucial for probing star formation and circumstellarregions. Band 1 will allow routine surveys of the new continuum component at its spectralmaximum. Pulsars generate magnetized particle winds that inflate an expanding bubble called apulsar wind nebula (PWN) whose outer edge is confined by the slowly expanding supernovaejecta. Electrons and positrons are accelerated at the termination shock some 0 . Chandra satellite (e.g., Helfandet al. 2001).ALMA has the sensitivity and resolution necessary to detect PWNe features at highradio frequencies, where we can detect the emission from relativistic particles that havemuch longer lifetimes than in X-rays. At cm/mm-wavelengths, flat-spectrum synchrotronPWNe stand out over steep-spectrum SNRs [e.g. as seen in the Vela PWN, Hales et al. (2004)and also as discussed in Sec. 4.1.3 below, and illustrated in Fig. 8, taken from Bietenholzet al. (2004)] with minimal confusion from the Rayleigh-Jeans tail of submm dust. ALMABand 1 will allow observations in the frequency regime where PWNe dominate, and bridgean important gap in frequency coverage, where spectral features such as power-law breaksoccur and linear polarization observations do not suffer from significant Faraday rotation.
Radio supernovae occur when the blast wave of a core-collapse supernova (SN) sweepsthrough the slowly expanding wind left over from the progenitor red supergiant. Particle ac-celeration and magnetic field amplification lead to synchrotron radiation in a shell boundedby the forward and reverse shocks (Chevalier 1982). In general, free-free absorption of theradiation in the ionized foreground medium coupled with the expansion of the SN causes 16 –the radio light curve first to rise at high frequencies and subsequently at progressively lowerfrequencies while the optical depth decreases. When the optical depth has reached approxi-mately unity, the radio light curve peaks and decreases thereafter (e.g., Weiler et al. 2002).These characteristics allow estimates to be made of the density profiles of the expandingejecta and the circumstellar medium and also of the mass loss of the progenitor. Resolvedimages of SNe provide information, e.g., on the structure of the shell, size, expansion ve-locity, age, deceleration, and magnetic field, in addition to refined estimates of the densityprofiles and the mass loss (Bartel et al. 2002). Radio observations of SNe can be regardedas a time-machine where the history of the mass loss of the progenitor is recorded tens ofthousands of years before the star died. Finally, the SN images can be used to make a movieof the expanding shell of radio emission and to obtain a geometric estimate of the distanceto the host galaxy (Bartel et al. 2007).ALMA Band 1 will allow exciting science to be done in the areas of radio light curvemeasurements, imaging of a nearby SN and, in conjunction with VLBI, imaging of moredistant SNe. Depending on the medium, the delay between the peak of the radio light curveat 20 cm and 1 cm can be as long as 10 years, as for instance was the case of SN 1996cr(Bauer et al. 2008). Absorption can also occur in the source itself. In case of SN 1986J,a new component appeared in the radio spectrum and in the VLBI images about 20 yearsafter the explosion and then only at or around 1 cm wavelength. The component is locatedin the projected center of the shell-like structure of the SN and may be emission from a verydense clump fortuitously close to that center, or possibly from a pulsar wind nebula in thephysical center of the shell (Figure 8, Bietenholz et al. 2004, 2010). Observations in Band 1minimize the absorption effect relative to observations at longer wavelengths and thus allowinvestigations of SNe at the earliest times without compromising too much on the signal tonoise ratio of a source with a steep spectrum. ALMA with Band 1 has the sensitivity tomeasure the radio light curves of 10s to 100 SNe. In addition, ALMA with Band 1 may beparticularly sensitive in finding SN factories in starburst galaxies (e.g., Lonsdale et al. 2006)where relatively large opacities would otherwise hinder or prevent discovery.ALMA with Band 1 will allow high-dynamic range images of SN 1987A in the LargeMagellanic Cloud with a resolution of about 300 FWHM beams across the area of the shellin 2014. Such data would be a significant improvement over presently obtainable images(Gaensler et al. 2007; Laki´cevi´c et al. 2012). Also, since the size of the SN increases byone Band 1 FWHM beam width per 3 years, the expansion of the shell can be monitoredaccurately and in detail, making this SN an important target for ALMA.ALMA with Band 1 could also be an important element of VLBI arrays. SN VLBIobservations at 1 cm wavelength have provided clues about physical conditions at the earliest 17 –times after the transition from opaqueness to transparency, and SN VLBI with Band 1 willfocus on this area of research.
X-ray binaries (i.e., binary star systems with either a neutron star or a black holeaccreting from a close companion) frequently show jet emission. Most of these systemsare transients. Typically, 1-2 black hole X-ray binaries undergo a transient outburst peryear, while neutron stars outburst at a slightly higher rate. Outbursts typically last severalmonths (although there are some which are both considerably longer or shorter), and duringoutbursts, X-ray luminosities can change by as much as 7 orders of magnitude. The radioluminosities of systems seen to date correlate well with the hard X-ray luminosities (i.e., theluminosity above ∼
20 keV), albeit with considerable, yet poorly understood scatter.When the X-ray spectra become dominated by thermal X-ray emission, the radio emis-sion often turns off (e.g., Tananbaum et al. 1972; Fender et al. 1999), but the extent to whichthe flux turns down is still poorly constrained. This turndown is not seen in neutron starX-ray binaries (Migliari et al. 2004). The reduced radio emission in black hole X-ray binarieswhen they have soft X-ray spectra can be explained by models of jet production in whichthe jet power scales with the polodial component of the magnetic field of the accretion flow(e.g., Livio, Ogilvie & Pringle 1999), and may have implications for the radio loud/quietquasar dichotomy (e.g., Meier 1999; Maccarone, Gallo & Fender 2003). The still-presentradio emission from neutron stars in their soft state may be indicating that the neutronstar boundary layers play an important role in powering jets (Maccarone 2008). The softstates of X-ray transients are short-lived. During them, there may be decaying emission fromtransient radio flares launched during the state transitions. Therefore, to place better upperlimits on the radio jets produced during the soft state, a high sensitivity, high frequencysystem with a very high duty cycle is needed.The radio properties of X-ray binaries with neutron star primaries are much more poorlyunderstood than those of black hole X-ray binaries. This situation is partially because theneutron star X-ray binaries are fainter in X-rays than are the black hole X-ray binaries.There is, however, additionally some evidence that neutron star X-ray binaries show a steeperrelation between X-ray luminosity and radio luminosity than do the black hole X-ray binaries,with L R ∝ L . X for the black holes and L R ∝ L . X for the neutron stars. This difference maybe explained if the neutron stars are radiatively efficient (i.e., with the X-ray luminosityscaling with the accretion rate) while the black holes are not (i.e., with the X-ray luminosityscaling with the square of the accretion rate, as has been proposed by Narayan & Yi 1994) 18 –– see Koerding et al. (2006). Radio/X-ray correlations for neutron star X-ray binaries are,to date, based on small numbers of data points from few sources, and the most recent work(Tudose et al. 2009) indicates that the situation may be far more complex than the picturepresented above. Near-IR and radio observations provide compelling evidence that the compact nonther-mal radio source Sgr A* is identified with a 4 × solar mass black hole at the center of theGalaxy (Reid and Brunthaler 2004; Ghez et al. 2008; Gillessen et al. 2009). It is puzzling,however, that the bolometric luminosity of Sgr A* due to synchrotron thermal emission fromhot electrons in the magnetized accretion flow is several orders of magnitude lower than ex-pected from the accretion of stellar winds. There have been two different approaches toaddress this puzzling issue. One is to search for the base of a jet from Sgr A* and identifyinteraction sites of a jet with the ionized and molecular material surrounding Sgr A*. Theother is to study the correlations of the variable emission from Sgr A* at centimeter andmillimeter bands. Studies of images and variability are well suited using ALMA’s Band 1and will be complementary to each other in addressing the key question as to why Sgr A* isso underluminous. A Jet from Sgr A*
Recent JVLA observations at radio wavelengths presented a tantalizing detection of a jet-likelinear feature appearing to emanate from Sgr A* (Yusef-Zadeh et al. 2012). Figure 9 showsa grayscale 23 GHz image of the inner 30 (cid:48)(cid:48) of Sgr A*. A new linear feature is noted runningdiagonally crossing the bright N and W arms of the mini-spiral, along which several blobs(b, c, d, h1 and h2) are detected. What is interesting about the direction in which the linearfeature is detected is that several radio blobs have X-ray and FeII/III counterparts also alongthe axis of the linear structure. In addition, the extension of the linear feature appears to bepolarized at 8 GHz, suggesting that this feature is a synchrotron source. The radio-polarizedlinear jet-like structure is best characterized by a mildly relativistic jet-driven outflow fromSgr A*, and an outflow rate γ ˙ M ∼ − M (cid:12) yr − .The linear arrangements of antennas in the JVLA configurations can lead to linearstructures in the residual beam pattern due to deconvolution errors. ALMA’s configurations,however, should lead to data with better, more-uniform uv coverage and will establish thereality of the linear structure. In particular, Band 1 will be most effective in studying the 19 –faint jet-like feature from Sgr A*. Dust emission from the immediate environment of SgrA* dominates fluxes at shorter wavelengths relative to optically thin non-thermal emissionfrom the jet with a steep energy spectrum. Thus, observations with Band 1 are critical formeasuring properly the morphology, spectral index and polarization characteristics of thejet emanating from Sgr A*. Although, Sgr A* is a unique object in the Galaxy, similarmotivations also apply to other non-thermal radio continuum sources such as microquasars,e.g., 1E1740.7-2942, having faint radio jets and are located in the inner Galaxy. Time Delay at Centimeter Wavelengths
Recent radio measurements have detected a time delay of ∼ ±
10 minutes between thepeaks of 7 mm and 13 mm radio continuum emission toward Sgr A* (Yusef-Zadeh et al.2006). This behavior is consistent with a picture of a flare in which the synchrotron emissionis initially optically thick. Flaring at a given frequency is produced through the adiabaticexpansion of an initiallyoptically thick blob of synchrotron-emitting relativistic electrons.The intensity grows as the blob expands, then peaks and declines at each frequency that theblob becomes optically thin. This peak first occurs at 43 GHz and then at 22 GHz about30 minutes later. Theoretical light curves of flare emission, as shown in Figure 10, showthat flare emission occurs at high near-infrared frequencies first and is increasingly delayedat successively lower ALMA frequencies that are initially optically thick.The limited time coverage of JVLA observations at radio wavelengths means that therecan be a large uncertainty in determining the underlying background flux level of a par-ticular flare, as well as difficulty identifying flares in different bands. Observations of SgrA* with a long time coverage using ALMA’s Band 1 can fit the corresponding light curvessimultaneously to place much tighter constraints on the derived physical parameters of theflare emission region. Two of the parameters that are of interest are the expansion speed ofthe hot plasma and the initial magnetic field. These quantities characterize the nature ofoutflow and cooling processes relevant to millimeter emission. The fitting of a light curveat one frequency will automatically generate models for any other frequency. We should beable to test the time delay between the peaks of flare emission within Band 1.What has emerged from past observing campaigns to study Sgr A* is that radio, sub-millimeter, near-infrared, and X-ray emission can be powerful probes of the evolution of theemitting region since they are all variable. We now know that flare emission at infraredwavelengths is due to optically thin synchrotron emission that is detected when a flare islaunched (Eckart et al. 2006). The relationship between radio and near-infrared/X-ray flareemission has remained unexplored due the very limited simultaneous time coverage betweenradio and infrared telescopes. The continuous variations of the radio flux on hourly time 20 –scale also makes the identification of radio counterparts to infrared flares difficult. In spite ofthe limited coverage in time, the strong flaring in near-infrared/X-ray wavelengths has givenus an opportunity to examine if there is a correlation with variability at radio frequencies.A key motivation for observing Sgr A* is to compare its flaring activity with the adiabaticexpansion picture. One of the prediction of this model is a time delay between the peaks ofoptically thin near-infrared emission and optically thick radio emission, as discussed above.From this model, a near-infrared flare of short duration of 0.5-1 hr is expected to have aradio counterpart shifted in time by 3-5 hr of duration of ∼ ± Close Encounters of Gas Clouds with Sgr A*
A 3 M
Earth cloud of ionized gas and dust named G2 has been recently determined to beon a collision course with Sgr A*. VLT observations indicate that the G2 cloud approachespericenter in mid-2013 and it will be disrupted and portions will likely be accreted by themassive black hole residing there (Gillessen et al. 2012). At the pericenter distance, thevelocity of the gas cloud will be 5400 km s − . Accordingly, the cloud is expected to producea bow shock that can easily accelerate electrons into a power-law distribution of index p =2 . − .
5, assuming standard shock conditions (Narayan et al. 2012). Depending on p ,the expected additional emission from Sgr A* ranges from 0.6 Jy to 4 Jy, over a dynamicaltimescale of ∼ Much of what we know about galaxy clusters has come from X-ray observations of ther-mal bremsstrahlung emission of the intra-cluster medium (ICM). For example, the angularresolution of
Chandra has been crucial to advancing our understanding in this area and hasresulted in a renaissance in astrophysical studies of galaxy clusters. In recent years, theSunyaev-Zel’dovich Effect (SZE) has provided an increasingly important view of these cos-mic structures (Birkinshaw 1999). Since the SZE signal is proportional to the product of theelectron density and its temperature ( ∼ n e T e , compared to n e √ T e for the x-rays), it givesa complementary view of the physical state of the ICM, one more sensitive to hot phasesthat also directly measures local departures from thermal pressure equilibrium. To date,the majority of SZE observations have been carried out at comparatively low angular reso-lution (beams > (cid:48) in size), yielding information about the overall bulk cluster properties.Advances in instrumentation have begun making higher angular resolution measurements ofthe SZE possible, revealing previously unsuspected shock-heated gas in the ICM of clusterspreviously thought to be dynamically relaxed (Komatsu et al. 2001, Kitayama et al. 2004,Mason et al. 2010, Korngut et al. 2011, Plagge et al. 2012). These 10 (cid:48)(cid:48) to 20 (cid:48)(cid:48) SZE imagesare the current state of the art. A Band 1 receiver suite on ALMA will surpass this bench-mark, making possible detailed studies of the ICM using the SZE on larger samples and withgreater sensitivity than before.ALMA Band 1 will be capable of addressing a wide range of basic questions about theobserved structure and evolution of clusters. For example, what is the structure of ICMshocks and the mechanism(s) responsible for converting gravitational potential energy intothermal energy in the ICM (Markevitch et al. 2007, Sarazin et al. 1988)? What is the influ-ence of Helium ion sedimentation within the cluster atmosphere (Ettori et al. 2006)? What isthe nature of the AGN-inflated “bubbles” seen in the cores of some clusters (Pfrommer et al.2005), and what is the role of cosmic rays in the ICM? What is the nature of the underlyingICM turbulence (e.g., Kolmogorov versus Kraichnan)? A particularly rich area will be thedetailed study of ICM shocks, which are common since infalling sub-clusters are typicallytranssonic. Several galaxy cluster mergers have been observed recently with Chandra andXMM in X-rays with resolutions at the arcsecond level where substructures become visible(Markevitch, et al. 2000, 2002). The features of interest for these studies will typically fitwithin one or a few ALMA Band 1 fields-of-view and require longer integrations (several to 22 – ∼
10 hours per pointing). Note that Band 1 also may have the sensitivity to detect the SZEfrom the halos of massive individual ellipticals or massive groups.Another important area where high-resolution SZE imaging will have an impact is theinterpretation of SZE survey data. ACT (Dunkley et al. 2011), SPT (Williamson et al.2011), and PLANCK (Planck Collaboration, 2011) have all conducted 1000 + deg surveysto detect and catalog galaxy clusters via the SZE. These surveys provide unique and valuableinformation about cosmology but their interpretation depends upon assumptions about therelationship between the SZE signal and the total virial mass of the halos observed. It isknown that both gravitational (cluster merger) and non-gravitational processes (AGN andsupernova feedback, bulk flows , cosmic ray pressure) give rise to considerable scatter andpotential biases (e.g., Morandi et al. 2007) in this relationship. Cluster mergers have aparticularly dramatic effect on the SZE, typically generating transsonic (Mach ∼ (cid:48)(cid:48) − (cid:48)(cid:48) resolution) galaxy clusters discovered in the low-resolution ( ∼ (cid:48) ) surveys, detecting shocksand mergers and identifying ICM substructure, and providing a direct, phenomenologicalhandle on important survey systematics. Indeed, the sensitivity and resolution of an ALMABand 1 receiver suite allows for efficient follow-up observations of cluster detections madeby blind southern hemisphere SZE surveys. Thus a study of the cluster selection of thesesurvey experiments in a statistical manner becomes feasible and new important insights intothe mass-observable relation and its scatter and dependence on cluster physics can poten-tially be obtained. The ability to understand cluster selection in detail is essential to derivereliable constraints on cosmological models from SZE cluster surveys (see e.g., Geisbueschet al. 2005; Geisbuesch & Hobson 2007).The coming decade will also see an explosion of optical and X-ray cluster data. TheGerman/Russian satellite eRosita , due to launch in 2014, will carry out the first all-sky X-raysurvey since ROSAT (Merloni et al. 2012). Among other things, it is expected to catalog By bulk flow, we refer to the motion of a cluster itself through its surrounding medium, producinga kinematic contribution to the observed SZE signal; in theory, this contribution has a different spectraldependence than the thermal SZE and may be distinguishable with good spatial coverage.
23 – ∼ z = 1 . ,
000 deg , mostly southern sky surveyalso expected to find ∼ z clusters, such as the ACT-discovered SZEcluster “El Gordo” at z = 0 .
89, weighing in at M = (2 . ± . × M (cid:12) (Menanteau et al.2011), offer leverage on so-called “pink elephant” tests capable of constraining cosmologicalor gravitational theories based on the existence of individual extreme objects, i.e., providedtheir properties are accurately determined. It is worthy of note that in addition to the high-resolution capability, a Band 1 equipped ALMA Compact Array (ACA) will be comparablein capability to the OVRO/BIMA arrays which have been used in the current decade tomeasure the bulk SZE properties of large northern hemisphere cluster samples (Bonamenteet al. 2008). Extending this capability to the southern hemisphere over the next decade isimportant to realize the full potential of these rich cluster samples.Given the large number of ALMA baselines, the resulting high image fidelity and dy-namic range of the data will be advantageous to SZE studies, in particular the detailed ones.In addition, long baseline data from ALMA can be used to remove accurately intrinsic andbackground (i.e., gravitationally lensed) discrete source populations. These latter objectsare a signal of substantial interest from another point of view, but they also set a significant“confusion noise” floor to millimeter single-dish observations, especially considering the fac-tor of 2 − D ∼ (cid:48) ) of a moderatelymassive SZE cluster with a merger shock (Figure 15). We considered a hypothetical projectaiming to detect a feature with a Compton y = 10 − – a characteristic of strong shocksin major mergers – over the virial radius of a cluster, with a characteristic feature sizeof 5 (cid:48)(cid:48) − (cid:48)(cid:48) . The required flux density sensitivity is similar in both cases after allowing forresolution effects, about 8 − µ Jy RMS (1 σ ) in both instances. We find that a clear detectionis achieved in only 1 . (cid:48)(cid:48) FWHM beam. Yamada et al. (2012) findsimilar results in a detailed study of SZE imaging with ALMA and the ACA at λ ≈ Previous single-dish millimeter molecular-line observations have found that moleculardistributions differ significantly between individual dark cloud cores. A widely accepted inter-pretation of this chemical differentiation is that there exists non-equilibrium gas-phase chem-ical evolution through ion-molecule reactions within dark cloud cores. Younger cores are richin “early-type” carbon-chain molecules such as CCS and HC N, while more evolved cores,closer to protostellar formation via gravitational collapse, are rich in “late-type” moleculessuch as NH and SO (Suzuki et al. 1992). Recent high-resolution millimeter-line observations,however, have revealed that there are even finer variations of molecular distributions withincores down to ∼ –2 N 4–3 36.392332 GHzHCS + N 17–16 45.264721 GHzCCS 4 –3 N 5–4 45.490316 GHzCCCS 8–7 46.245621 GHzC H , –2 , S 1–0 48.206956 GHzCH OH 1 –0 , –3 , N 19–18 50.58982 GHzDC N 6–5 50.65860 GHzO N=35-35, J=35-34 50.98773 GHzCH CHO 1(1,1)-0(0,0) 51.37391 GHzNH D 1(1,0)–1(1,1) 51.47845 GHzCH CHCHO 1 –0 H , –0 , ALMA Band 1 will provide a unique opportunity to search for new complex organicmolecules, including the amino acids and sugars from which life on Earth may have originallyevolved. In addition, these complex molecules provide a powerful tool for understanding starformation and the processes surrounding it.There are several reasons why Band 1 is the best place to search for complex molecules.First, the heavier a molecule, the lower will be its rotational transition frequencies. The manyabundant lighter molecules (e.g., CO, HCN, CN) have their lowest transitions in Band 3, andso do not appear at all in Band 1. Therefore, Band 1 does not suffer from contamination from 26 –these common molecules, and so line confusion is much less of a problem. Second, the systemtemperatures in Band 1 are significantly lower than in higher bands, giving extra sensitivityto detect weak transitions from less abundant complex molecules, such as glycolaldehyde,the simple sugar known to exist in the interstellar medium. Table 2 lists some complexcarbon-chain molecules whose transitions have been already detected in the ISM. Note thatsearches for complex molecules can be made with Band 1 using emission lines, but also usinglines in absorption against bright background objects like, e.g., young stars or quasars.There is now a significant body of evidence to suggest that complex biological molecules,such as amino acids and sugars needed for evolution of life on Earth, evolved in the interstellarmedium (e.g., see Holtom et al. 2005; Hunt-Cunningham & Jones 2004; Bailey et al. 1998).Band 1 will be one of the best instruments in the world to test this hypothesis observationally.Table 2: Some detected ISM complex carbon chain moleculesCH CHCN propenitrileCH CNH ketenimineCH C H methyldiacetyleneCH CCCN methyl cyanoacetyleneCH CH CN ethyl cyanideCH CHO acetaldehydeCH CONH acetamideCH OCH ethyl butyl etherCH OCHO methyl formateC H − hexatriyne anionC H octatetraynylH CCCC cumulene carbeneHCCCNH + · · · Radio recombination lines (RRLs) are powerful, extinction-free diagnostics of the ionizedgas in young, star-forming regions. Dozens of galaxies have been detected in RRLs so far,and improved sensitivity ALMA Band 1 and the JVLA will open new capabilities in galacticand extragalactic RRL studies (e.g., Peters et al. 2012). ALMA and the JVLA will also eachhave a significant impact on the study of photoevaporating protoplanetary disks using RRLs 27 –(Pascucci et al. 2012). RRLs do not suffer from dust obscuration and thus provide a powerfulmethod for studying the kinematics, structure, and physical properties of ionized gas suchas the ionizing photon flux, density, filling factor, and electron temperature in the nuclearregions of dusty galaxies. A comparison of the predicted strengths of NIR recombinationlines (e.g., Br α and Br γ lines) based on RRLs with actual observations also leads to revisedestimates of the dust extinction, which are essential for probing the heavily dust-enshroudednuclei of galaxies such as Arp 220 (e.g., Kepley et al. 2011).RRLs are observed over a wide range of frequency, from a few GHz up to a few 100GHz (e.g., Puxley et al. 1997). It is important to observe RRLs at multiple, widely spacedfrequencies because they trace different parts of the ionized medium (e.g., Anantharamaiahet al. 2000). The observed strength of the millimeter-wavelength RRLs indicates the presenceof a higher density ( n e ≈ cm ) ionized gas component (e.g., Zhao et al. 1997) and leads toinformation about star formation at recent epochs ( t ≈ yr). On the other hand, the lowerfrequency lines provide constraints on the amount of low density ( n e ≈ cm ) ionized gas.Therefore, adding a Band 1 receiver suite to ALMA is useful for studying the moderatelydense ionized medium through observations of lines such as H53 α (43.3094 GHz). Accuratemeasurements of both the RRL and continuum emission are necessary to model the RRLemission, and ALMA gives an ideal combination of the large instantaneous band width andhigh spectral resolution. To make an accurate determination of the free-free component,multi-wavelengths continuum measurements are also essential, so adding a Band 1 receiversuite is also beneficial for a better modelling of RRLs. Masers (Microwave Amplifications by Stimulated Emission of Radiation) frequently oc-cur in regions of active star formation, from molecular transitions whose populations areeither radiatively or collisionally inverted. A photon emitted from this material will interactwith other excited molecules along its path, stimulating further emission of identical photons.This process leads to the creation of a highly directional beam that has sufficient intensityto be detected at very large distances.Masers are observed from a variety of molecular and atomic species and each serves as asignpost for a specific phenomenon, a property which renders masers powerful astrophysicaltools (Menten 2007). More precisely, masers are formed under specific conditions, and thedetection of maser emission therefore suggests that physical conditions (e.g., temperature,density, and molecular abundance) in the region where the maser forms lie within a definedrange (c.f., Cohen 1995, Ellingsen 2004, and references therein). Therefore, interferometric 28 –blind and targeted surveys of maser species can lead to the detection of objects at interestingevolutionary phases (Ellingsen 2007).Table 3: ALMA bands with known maser lines (Menten 2007)Species ALMA BandsH O B3, B5, B6, B7, B8, B9CH OH B1, B3, B4, B6SiO B1, B2, B3, B4, B5, B6, B7HCN B3, B4, B6, B7, B9Theoretical models of masers strongly depend on physical conditions as well as the ge-ometry of the maser source. A successful model should be able to reproduce observationalcharacteristics of observed maser lines but also to predict new maser transitions (e.g., themodels of Sobolev 1997 for Class II methanol masers and Neufeld 1991 for water masers).In that respect, interferometry is essential for the successful search of candidate lines andconfirmation of their maser nature. ALMA, in particular, will resolve closely spaced maserspots and help further establish precise models of masing sources by determining if the de-tected maser signals are associated with thermal emission (Sobolev 1999), which is essentialfor improving theoretical models. With Band 1, ALMA will cover a wider frequency range,making it ideal for multi-transition observations of various maser species across the millime-ter and submillimeter windows. Examples of species with observed maser radiation in thedifferent ALMA bands are given in Table 3, while Tables 4 & 5 list SiO and methanol masertransitions that have been observed or predicted to be within Band 1.Maser radiation can be linearly or circularly polarized depending on the magnetic prop-erties of the molecule. Polarimetric studies of maser radiation with interferometers cantherefore yield information on the morphology of the magnetic field threading the region onsmall scales, with the plane-of-sky and line-of-sight components of the field being probedusing linear and circular polarization measurements, respectively (e.g., see Harvey-Smith2008, Vlemmings 2006). Polarization data are essential for improving on the theory of maserpolarization first introduced by Goldreich (1973a), which applies to a linear maser region, aconstant magnetic field, the simplest energy states for a masing transition, and asymptoticlimits. Observations at higher spatial resolution are needed to verify and improve on morerealistic and extensive models (Watson 2008). 29 –Table 4: Observed SiO maser lines in the Band 1 of ALMA (Menten 2007).Transitions Frequency (GHz)v=0 (J= 1 →
0) 42.373359v=3 (J= 1 →
0) 42.519373v=2 (J= 1 →
0) 42.820582v=0 (J= 1 →
0) 42.879916v=1 (J= 1 →
0) 43.122079v=0 (J= 1 →
0) 43.423585Table 5: Observed (Menten 2007) and predicted (designated with a star, Cragg et al. 2005)methanol maser lines in Band 1Transitions Frequency (GHz)4(-1) → → → + → − → + → ∗ → ∗ Magnetic fields are believed to play a crucial role in the star formation process. Varioustheoretical and numerical studies explain how magnetic fields can account for the supportof clouds against self-gravity, the formation of cloud cores, the persistence of supersonic linewidths, and the low specific angular momentum of cloud cores and stars (McKee & Ostriker2007). The standard model suggests that the initial mass-to-(magnetic) flux ratio, M/Φ init ,is the key parameter governing the fate of molecular cores. Namely, if M/Φ init of a core isgreater than the critical value, the core will collapse and form stars on short time scales,but for cores with M/Φ init smaller than the critical value the process of ambipolar diffusionwill take a long time to reduce the magnetic pressure (Mouschovias & Spitzer 1976; Shu et 30 –al. 1987). On the other hand, recent MHD simulations suggest that turbulence can controlthe formation of clouds and cores and in such cases the mass-to-flux ratio in the centerof a collapsing core will be larger than that in its envelope, the opposite of the ambipolardiffusion results (Dib et al. 2007). Therefore, measuring the magnetic field strengths and themass-to-flux ratios in the core and envelope provide a critical test for star formation theories.Despite its central importance, the magnetic field is the most poorly measured parameterin the star formation process. The main problem is that magnetic fields can be measuredonly via polarized radiation, which requires extremely high sensitivity for detections. Asa result, the observed data on magnetic fields is sparse compared with those related todensities, temperatures, and kinematics in star-forming cores. The large collecting area ofALMA provides the best opportunity to resolve the sensitivity problem for magnetic fieldmeasurements.The key to determining mass-to-flux ratios is the measurement of the strength of mag-netic fields. This measurement can be made directly through detection of the Zeeman effectin spectral lines. Observations of Zeeman splitting involve detecting the small differencebetween left and right circular polarizations, which is generally very small in interstellarconditions (with the exception of masers). Successful non-maser detections of the Zeemaneffect in molecular clouds have only been carried out with HI, OH, and CN lines becausethese species have the largest Zeeman splitting factors ( ∼ µ G) among all molecularlines (Crutcher et al. 1996, 1999; Falgarone et al. 2008). Thermal HI and OH lines, how-ever, probe relatively low-density gas ( n (H) < cm − ). Also, CN detections are difficult;Crutcher (2012) described only 8 CN Zeeman detections towards 14 positions observed withsignificant sensitivity.ALMA Band 1 receivers provide the opportunity to detect the Zeeman effect fromthe CCS 4 –3 line at 45.37903 GHz and hence greatly advance our understanding in starformation. CCS has been widely recognized as being present only very early in the star-forming process through chemical models (Aikawa et al. 2001, 2005) and observations (Suzukiet al. 1992; Lai & Crutcher 2000). Therefore the mass-to-flux ratio derived from the CCSZeeman measurements will be very close to the initial values before the onset of gravitationalcollapse. CCS 4 –3 also has a relatively large Zeeman splitting factor ( ∼ µ G;Shinnaga & Yamamoto 2000) compared to most molecules. ALMA’s antennas and site willbe excellent at these “long” wavelengths, providing the stability and accuracy needed forsuch sensitive polarization work. The linearly polarized detectors on ALMA’s antennas willalso be ideally suited to measurement of Stokes V signatures from CCS.Using the BIMA survey results from Lai & Crutcher (2000), Figure 13 demonstratesthat detections of CCS Zeeman effects can be achieved if the ALMA specifications for Band 31 –1 receivers are met. Zeeman effect detection depends on two factors: the magnetic fieldstrength and the line intensity. The two lines in Fig. 13 show the 3 σ detection limits forStokes V spectrum with channel width of 0.024 km s − and 1 hr or 10 hr integration time fora range of magnetic field strengths and line intensities. The channel width is chosen to haveat least 6 channels across the FWHM of the total intensity spectrum (Stokes I). If we scalethe line intensity from Lai & Crutcher (2000) assuming the intensity distribution is uniformwithin the 30 (cid:48)(cid:48) BIMA beam, the expected line intensity would be around 0.1-0.4 Jy forALMA observations with 10 (cid:48)(cid:48) beam. Therefore, Fig. 13 shows that for the magnetic fields of0.2-1 mG (typical values estimated from the application of the Chandrasehkar-Fermi methodto dust polarimetry in dense cores), we can detect the CCS Zeeman effect with reasonableon-source integration time (less than 10 hr).Note that the SiO v=1, J=1–0 transition at 43.12 GHz could be also used to probemagnetic fields using the Zeeman effect, under certain circumstances. Though its Zeemansplitting factor is lower than that of the above CCS lines, the Zeeman effect may be detectiblein situations where the SiO line is extraordinarily bright, e.g., as a maser (see McIntosh,Predmore & Patel 1994). (Non-Zeeman interpretations of circularly polarized SiO emissionhave also been advanced; see Weibe & Watson 1998).In summary, ALMA Band 1 receivers will provide the opportunity to measure the initialmass-to-flux ratio of molecular cores through the detection of the Zeeman effect, which cannotbe done with any other instruments in the foreseeable future. The results will allow us totest realistically the expectations from theoretical and numerical models for the first time.
Radio continuum emission is observed from the jets and winds of young stellar objectsand is due to the interaction of free electrons, i.e., “free-free emission.” The radio imagesappear elongated and jet-like and are usually located near the base of large optical Herbig-Haro flows (Reipurth & Bally 2000). These regions usually have only sub-arcsecond sizes,indicating the youth of the emitting material and the short dynamical times involved. Theemitted flux is usually weak, with a flat to positive spectral index with increasing frequency,and it can be obscured by the stronger thermal emission from dust grains at higher frequen-cies (e.g., Anglada 1995). Multi-wavelength studies of the brightest radio jets at centimeterwavelengths trace either earlier and stronger sources or more massive systems. The triplesystem L1551-IRS 5, one of the most studied low-mass systems (Rodriguez et al. 1998, 2003;Lim & Takakuwa 2006), is illustrative of the sub-arcsecond scales required (Figure 14). 32 –Ground-based, interferometric studies of radio jets provide the best opportunity to re-solve the finest scales of the underlying source, comparable or better than optical studiesof jets by the Hubble Space Telescope. Such fine-detail images can provide the ability todifferentiate between theoretical ideas about the nature of these jets; i.e., the launch region,the collimation process, and the structure of the inner disks. Modeling efforts with the radiocontinuum emission presented in Shang et al. (2004) demonstrate one such possibility inconstraining theoretical parameters using earlier millimeter and centimeter interferometers(Figure 14). Band 1 observations will discriminate between competing jet launch theoriestied to the disk location of the launch point by achieving better than 0.1 (cid:48)(cid:48) angular resolution.The high sensitivity of Band 1 on ALMA will also allow detection of radio emission fromless luminous sources. ALMA will thus have the potential to discover a significant number ofnew radio jets, providing a catalog from which evolutionary changes in the physical propertiescan be deduced. As well, multi-epoch surveys will be able to follow the evolution of the freshlyejected material down to a few AU from the driving sources through movies. The 35-52 GHzfrequency range of Band 1 will show contributions to the observed emission from both theionized component of the jet and the thermal emission from the dust. Together with lowerfrequency JVLA observations and detailed theoretical modelling, a complete understandingof properties of the spectral energy distribution (SED) from the ionized inner regions ofyoung stellar jets will be uncovered.
The Sub-millimeter Array (SMA) has proven to be a successful instrument for the studyof the youngest molecular outflows and jets from the most deeply embedded sources (e.g.Hirano et al. 2006; Palau et al. 2006; Lee et al. 2007a,b, 2008, 2009). The detection ofexcitation from rotational transitions of the molecule SiO up to levels J = 8–7 and CO upto J = 3–2 have uniquely identified a molecular high-velocity jet-like component locatedwithin the outflow shell. This component displays similarities to the optical forbidden linejets observed in T-Tauri stars (Hirano et al. 2006; Palau et al. 2006; Codella et al. 2007;Cabrit et al. 2007). These observations have provided a new probe of how jets are launchedand collimated during the earliest protostellar phase.One unique opportunity offered by Band 1 is observation of the J = 1–0 transition ofthe SiO molecule at 43.424 GHz. This transition has not yet been detected nor surveyedaround even the brightest molecular outflows, except using single-dish telescopes (Haschick& Ho 1990). One feature of this line that may be potentially distinct from the higher-Jtransitions of SiO is that it may be tracing the outer and more diffuse gas located on the 33 –outskirts of outflow shells that can be easily excited by shocks. Potential morphological andkinematic studies of the regions where the outflows interact with their own pre-natal cloudscould be contrasted with other transitions using knowledge of their excitation conditions. Roughly half of the high-redshift objects detected in CO line emission are believed tohost an active galactic nucleus (AGN). Although they are selected based on their AGNproperties, optically luminous high-redshift quasars exhibit many characteristics indicativeof ongoing star-formation, e.g., thermal emission from warm dust (Wang et al. 2008) orextended UV continuum emission. Indeed, galaxies with AGNs in the local Universe reveala strong correlation between the mass ( m ) in their supermassive black hole (SMBH) andthat in their stellar bulge (measured from the stellar velocity dispersion ( σ ); e.g., Kormendy& Richstone 1995; Magorrian et al. 1998; Gebhardt et al. 2000). Such a correlation can beexplained if the SMBH formed coevally with the stellar bulge, implying that the luminousquasar activity signaling the formation of a sub-arcsecond SMBH at high-redshift should beaccompanied by starburst activity. High spatial resolution observations of CO line emissionin high-redshift quasars can be used to infer the dynamical masses, which are found to becomparable to the derived molecular gas + black hole masses, meaning that their stellarcomponent cannot contribute a large fraction of the total mass. There is mounting evidencethat quasar host galaxies at redshifts z = 4–6 have SMBH masses up to an order of magnitudelarger than those expected from their bulge masses and the local relation (Walter et al. 2004;Riechers et al. in prep.), suggesting that the SMBH may have formed first. The possible timeevolution of the m − σ relation is of fundamental importance in studies of galaxy evolution,and this new finding needs to be made more statistically robust. Future observations of high-redshift AGN with the Band 1 receivers on ALMA would allow us to address this questionthrough the study of low-J CO line emission in galaxies beyond redshifts z ≈ . § z ∼ z ≈ .
3) to those which existed when the Universe wasreionized sometime before z > ∼ (cid:29)
100 M (cid:12) yr − ), significant masses of molecular gas ( > M (cid:12) ) have been discoveredin more modest star-forming galaxies at z = 1 . − . z -K colour diagram (Daddi et al. 2004) andhave star-formation rates of ∼
100 M (cid:12) yr − (Daddi et al. 2007), while their number densityis roughly a factor of 30 larger than that of the more extreme SMGs at similar redshifts.Observations of CO J=2–1 line emission in these BzK galaxies reveal comparable masses ofmolecular gas to that of the SMGs, so their star-formation efficiencies appear lower. Theexcitation conditions of their molecular gas (temperature and density) are similar to thoseof the Milky Way (Dannerbauer et al. 2008), as indicated by the “turnover” in the CO linespectral energy distribution occuring at the J=3–2 transition, i.e., lower than that of theSMGs which typically occurs at the J=6–5 or J=5–4 transition (Weiss et al. 2005). Todevelop a full spectral energy distribution for the CO line excitation, observations of thesegalaxies in the J=1-0 transition are needed with Band 1 receivers on ALMA. Such data willalso provide a more robust estimate of the total molecular gas mass, along with the spatialresolution needed to constrain the gas kinematics, as has been done for the SMGs (Tacconiet al. 2006). Indeed, recent high-resolution studies of CO J =1–0 from lensed Lyman Breakgalaxies (Riechers et al. 2010) and unlensed BzK galaxies (Aravena et al., in prep.) havebeen made with the JVLA. Also, CO J =1–0 emission has been detected with the JVLA orGBT towards SMGs Ivison et al. 2010, 2011; Frayer et al. 2011; Riechers et al. 2011a,b).
6. Other Considerations6.1. Weather Considerations at the ALMA Site
The ALMA site is exceptionally well-suited for Band 1 observing. Even during the worstoctile of weather, the typical optical depth through the Band 1 Receiver range is less than0.1. Though other frequency ranges like Band 3 can still use such weather, the addition ofcloud cover and water droplets in the air will make still lower frequency observations moreattractive. 35 –As shown in the previous sections, the top science cases for Band 1 can stand shoulder-to-shoulder with the primary Level 0 goals of ALMA. Thus, while Band 1 observations willclearly benefit from a larger fraction of available observing conditions, the primary motivationfor the enhancement is not as a “poor weather” back-up receiver but rather the excellentscience that can be achieved.
The Jansky Very Large Array (JVLA) currently has observing capability over the nom-inal Band 1 frequency range of 35–50 GHz, through its receivers in the K a -band (26.5–40GHz) and Q-band (40–50 GHz). Note, however, that the JVLA cannot observe at 50-52GHz, the extension proposed for ALMA Band 1. Moreover, it is important to note thatJVLA and ALMA are located at complementary latitudes, the former at +34 ◦ and the lat-ter at − ◦ . With both telecsopes, the entire sky will be observable at high resolution atBand 1 frequencies. (The southern hemisphere-based Australia Telescope Compact Array(ATCA) can also observe some Band 1 frequencies but at much lower relative sensitivitythan ALMA or JVLA, and hence we do not consider it further.) Here, we compare therelative capabilities of ALMA and JVLA in Band 1.Relative to ALMA, JVLA has fewer antennas (27 vs. 50) but these have larger surfaceareas (25-m diameter vs. 12-m), lower pointing accuracies (2-3 (cid:48)(cid:48) vs. 0.6 (cid:48)(cid:48) ) and lower apertureefficiencies at Band 1 (0.34–0.39 vs. 0.78). Combining these numbers (except pointing ac-curacy), the effective surface area of JVLA is a factor of 1.02–1.17 times that of ALMA.(Adding Band 1 to the ACA antennas would minimize even this small difference.) Due to itsWIDAR correlator, JVLA will have the same 8 GHz maximum bandwidth as ALMA. TheWIDAR correlator can provide users with much higher spectral resolution than ALMA’scorrelator can, i.e., a maximum of below 1 Hz vs. 3.82 kHz.The JVLA is located at a lower altitude (2124 m) than ALMA (5000 m) and its localatmosphere has correspondingly larger typical precipitable water vapor. Weather statisticsare needed for both the Plains of San Augustin and the Llano de Chajnantor to compareproperly the availability of Band 1 observing conditions at both sites. An important factorthat also must be quantified is the stability of the atmosphere above both sites.Given differing antenna numbers, sizes, and baselines, the two observatories differ invarious imaging metrics: • Comparing the face value “single-field sensitivity” (
N D ; where D is the antennadiameter and N is the number of antennas), JVLA is over twice as “sensitive” as 36 –ALMA (17000 vs. 7200). The ALMA Band 1 receiver specifications require receivertemperatures of 40–80 K over the Band, the same as the specification for the JVLAK a /Q-band receivers. Assuming similar system temperatures for each set of receiverswhenever local weather conditions are appropriate for Band 1 observing and factor-ing aperture efficiencies into account, the continuum sensitivities for Band 1 are sim-ilar for both observatories { footnoteNote also that for all sensitivity calculations inthis section we assume the original ALMA specifications for Band 1 receiver per-formance.. For example, a 1 σ rms of ∼ µ Jy beam − is expected at 35 GHz after1 hour of integration at both observatories. At higher frequencies (e.g., 45 GHz),however, the point source sensitivity of ALMA is better than that of JVLA by afactor of ∼
2. Note also that the JVLA sensitivities require the best weather whilea relatively high PWV was chosen for ALMA in these calculations. Table 6 pro-vides comparison of JVLA and ALMA sensitivities for point sources across the pro-posed Band 1 frequency coverage estimated using the JVLA and ALMA sensitiv-ity calculators ( https://science.nrao.edu/facilities/evla/calibration-and-tools/exposure and http://almascience.eso.org/call-for-proposals/sensitivity-calculator ). In addition,Figure 16 shows simulated “blank-sky” observations at 45 GHz carried out with CASA,giving another perspective on this comparison. (Note that additional parameters suchas pointing accuracies and phase stabilities have not been fully incorporated in thesecalculations.) • Comparing “mosaic image sensitivity” (
N D ), again on face value, JVLA and ALMAare similar (680 vs. 600). JVLA’s smaller number of baselines, however, yield a lower“image fidelity” ( N ) by a factor of > • At present, JVLA has maximum baselines that are a factor of 2 larger than ALMA’s(36.4 km vs. 15-18 km), suggesting the JVLA can produce images of resolution up toa factor of 2 higher than ALMA can at the same frequency.In summary, at 35–50 GHz, the JVLA and ALMA are effectively the same instrument atvery complementary latitudes and with somewhat different strengths. JVLA can go to lowerfrequencies than 35 GHz but ALMA Band 1 may include 50–52 GHz. Weather conditionsat both observatories suggest Band 1 observing conditions are relatively rare for JVLA butshould be common for ALMA, but the relative fractions of available “Band 1 time” need tobe determined. Also, the fraction of time at ALMA where no other frequency but Band 1can be observed needs to be determined. 37 –Table 6: Comparison of Point-Source Sensitivity between JVLA and ALMAJVLA ALMAno. of antennas 25 50polarization dual dualweather winter auto (5.2mm) PWVsource position zenith zenithweighting natural naturalon-source time 60 s 1 hrs 60 s 1 hrsbandwidth 1MHz 1MHzfreq. 35 GHz 3.2 mJy 0.41 mJy 3.0 mJy 0.38 mJy40 GHz 3.6 mJy 0.47 mJy 3.1 mJy 0.40 mJy45 GHz 5.1 mJy 0.68 mJy 3.6 mJy 0.47 mJy50 GHz 25.5 mJy 3.29 mJy (not avail.) (not avail.)bandwidth 8GHz 8GHzfreq. 40 GHz 50 µ Jy 5.4 µ Jy 35 µ Jy 4.5 µ Jyfreq. 45 GHz 78 µ Jy 10 µ Jy 41 µ Jy 5.3 µ Jy
7. Summary
The Band 1 receiver suite has been considered an essential part of ALMA from theearliest planning days. Even through the re-baselining exercise in 2001, the importanceof Band 1 was emphasized. With the ALMA Development Plan taking shape, we haveundertaken an updated review of the scientific opportunity at these longer wavelengths. Thisdocument presents a set of compelling science cases over this frequency range. The sciencecases reflect the new proposed range of Band 1, 35-50 GHz (nominal) with an extension upto 52 GHz, which was in fact chosen to optimize the science return from Band 1. The sciencecases range from nearby stars and galaxies to the re-ionization edge of the Universe. Twoprovide additional leverage on the present ALMA Level One Science Goals and are seen asparticularly powerful motivations for building the Band 1 receiver suite: (1) detailing theevolution of grains in protoplanetary disks, as a complement to the gas kinematics, requirescontinuum observations out to ∼
35 GHz ( ∼ < z <
10 alsorequires Band 1 receiver coverage. 38 –Table 7: Comparison of angular scale coverage between JVLA and ALMA at 45 GHzJVLA ALMAConfiguration A D most extended most compactB min (km) 0.68 0.035 0.04 0.015B max (km) 36.4 1.03 16 0.15 θ P RIMARY
60 60 135 135 θ HF BW θ LAS
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M from 1.3mm
Taurus Ophiuchus M d i sk f r o m mm d a t a ( M s un ) Mdisk from 1.3mm data (Msun)
TaurusOphiuchus
Fig. 1.— Disk masses measured from 9 mm continuum emission compared to those measuredfrom 1.3 mm continuum emission in the regions of Taurus and Ophiuchus. Many disks showhigher mass measurements at the longer wavelength, indicating the presence of larger grainsthan those detected at 1.3 mm measurements. (Greaves et al., in prep.) 48 –Fig. 2.— Spectral energy distribution plot showing the differences between three disk modelshaving different maximum grain sizes. The solid curve is the model with a max = 1 cm, whichkeeps declining with roughly constant slope all the way to 1 cm. The two dashed curvesare for a max = 10 µ m and 1 mm. The top one, which breaks around 5 mm is the modelwith a max = 1 mm. It’s interesting to note how the fluxes are very much the same fora max = 1 mm or 1 cm, except precisely towards ALMA’s Band 1. There is at least an orderof magnitude difference in power at 1 cm between the max size = 1 mm versus the max size = 1cm disks. These models indicate that observations at the ALMA Band 1 regime are crucialfor determining whether grain-growth to cm-sizes is indeed occurring. 49 –Fig. 3.— Spectral energy distribution of TW Hya, showing the fit to the SED for an irradi-ated accretion disk model with a maximum particle size of 1 cm (Calvet et al. 2002). 50 –Fig. 4.— (Left) Image from an SPH simulation showing the surface density structure of a0.3 M (cid:12) disk around a 0.5 M (cid:12) star. A single dense clump has formed in the disk, at a radiusof 75 AU and with a mass of ∼ Jup . (Right) VLA 1.3 cm images toward HL Tau. Themain image shows natural weighting with a beam of 0.11 (cid:48)(cid:48)
FWHM. The arrow indicates thejet axis. Upper inset: compact central peak subtracted. Lower inset: uniform weighting,with a beam of 0.08 (cid:48)(cid:48)
FWHM. The compact object lies to the upper right hand side. Thiscondensation was also detected at 1.4 mm with the BIMA array (Welch et al. 2004). 51 –Fig. 5.— VLA redshifted CO J =3–2 map of the quasar J1148+5251 using the combined B-and C-array data sets (covering the total bandwidth, 37.5 MHz or 240 km s − ), from Walteret al. (2004). Contours are shown at –2, –1.4, 1.4, 2, 2.8, and 4 × σ (1 σ = 43 µ Jy beam − ).The beam size (0.35 (cid:48)(cid:48) × (cid:48)(cid:48) ) is shown in the bottom left corner; the plus sign indicates theSDSS position (and positional accuracy) of J1148+5251. 52 –Fig. 6.— Observable frequencies of CO rotational transitions and [CII] P / – P / as afunction of redshift. The frequency ranges of the ALMA Bands are also shown. Note thatthe range for Band 1 reflects the new nominal range of 35-50 GHz. 53 –Fig. 7.— Three-colour image of the ρ Oph W photo-dissociation region (Casassus et al.2008).
Red : MIPS 24 µ m continuum Green : IRAC 8 µ m continuum, dominated by the7.7 µ m PAH Band Blue : 2MASS K s -band image. The x − and y − axes show offset in RAand Dec from ρ Oph W, in degrees. The contours follow the 31 GHz emission, with levelsat 0.067, 0.107, 0.140, 0.170, and 0.197 MJy sr − . 54 –Fig. 8.— Two-colour VLBI image of SN 1986J highlighting the emergence of a centralcomponent. The red colour and the contours represent the 5.0 GHz radio brightness. Thecontours are drawn at 11.3, 16. 22.690.5% of the peak brightness of 0.55 mJy/bm. The blueto white colours show the 15 GHz brightness of the compact, central component. The scaleis given by the width of the picture of 9 mas. North is up and east to the left. For moreinformation on the emergence of the compact source, see Bietenholz et al. (2004). 55 – d d c b h1 h2 Sgr A*
Fig. 9.— (a) Left
A 22 GHz image of 0 . (cid:48)(cid:48) × . (cid:48)(cid:48) resolution (PA=2 ◦ ) constructed bycombining JVLA A- and B- array data. 56 –Fig. 10.— Theoretical light curves of Stokes I for optically thick synchrotron emission atfive different bands corresponding ALMA Bands 3, 6, 7 and 9 as a function of expandingblob radius. These light curves assume an energy power law index p=1 where n(E) ∝ E − p . 57 – a Fig. 11.— The light curves of Sgr A* on 2007 April 4 obtained with XMM in X-rays(top), VLT and HST in NIR (middle), and IRAM-30m and VLA at 240 GHz and 43 GHz,respectively (bottom). The NIR light curves in the middle panel are represented as H (1.66 µ m) in red, K s and K s -polarization mode (2.12 µ m) in green and light blue, respectively,L’ (3.8 µ m) in black (Dodds-Eden et al. 2009), and NICMOS of HST in blue at 1.70 µ m.In the bottom panel, red and black colors represent the 240 GHz and 43 GHz light curves,respectively. 58 –Fig. 12.— Radio emission as a function of frequency expected from G2 cloud (red) whencompared to quiescent emission from Sgr A*, as shown in blue (Narayan, Ozel, & Sironi2012). Left and right panels show predictions based on different assumptions on the energyspectrum of nonthermal particles (p). 59 –Fig. 13.— The expected detection limits (3 σ ) with integration time of 1 hr and 10 hr for arange of magnetic field strengths and CCS line intensity. 60 –Fig. 14.— Background : The VLA+Pie Town continuum image of L1551 IRS 5 at 3.5 cmobtained by Rodriguez et al. (2003) in their Figure 1. The size of the beam (0.18 X 0.12 (cid:48)(cid:48) ;P.A. = 35 ◦ ) is shown in the bottom left-hand corner. Black rectangles mark the positionsand deconvolved dimensions of the 7 mm compact protoplanetary disks. The dashed linesindicate the position angles of the jet cores. Inset : map of the south jet from the X-Windmodel convolved with the beam and plotted with the same contour levels from Figure 4 ofShang et al. (2004). 61 –Fig. 15.— Simulated 1 . (cid:48) × (cid:48) . The shock is represented as a Gaussian component 5 (cid:48)(cid:48) × (cid:48)(cid:48) in extent with a peak SZE of y = 10 − , considerably weaker than the amplitude observed inRXJ1347-1145 by Mason et al. (2010). The Band 3 data were tapered to the innate resolutionof the Band 1 map, ∼ (cid:48)(cid:48) (FWHM). ACA baselines were not included in this simulationbut the overplotted contours show the ACA Band 1 image (using a 45 (cid:48)(cid:48) taper) of the bulkICM in this system in a simulated 12 hr integration after subtraction of the shock signal.The bulk ICM is modeled as an elliptical isothermal β model with R core = (150 , β = 0 .
7, and y o = 3 × − at z = 0 .
7, characteristic of disturbed, merging systems. 62 –
JVLA (rms = 9.6 µ Jy)
ALMA (rms = 4.5 µ Jy) arcsecondarcsecond a r c s e c o n d Fig. 16.— Images from JVLA and ALMA observations simulated with CASA. The obser-vations were set toward a “blank” sky at 45 GHz with 8 GHz (continuum) bandwidth, withJVLA in its D-configuration while ALMA in its ∼ . (cid:48)(cid:48) σ rms noiselevels after 2 hours of on-source integration are 9.6 µ Jy and 4.5 µµ