Science with ALMA
Ast735: Submillimeter Astronomy IfA, University of Hawaii
1 Science with ALMA
Contents
1 Introduction 2 1.1 Millimeter Astronomy ...... 2 1.2 A Major New Telescope for Millimeter Astronomy ...... 2 1.3 The Atacama Large Millimeter Array ...... 3 1.4 Science with the Atacama Large Millimeter Array ...... 7
2 Galaxies and Cosmology 9 2.1 The High-Redshift Universe ...... 9 2.2 Gravitational lenses ...... 18 2.3 Quasar absorption lines ...... 20 2.4 Sunyaev-Zel’dovich observations with ALMA ...... 23 2.5 Gas in Galactic Nuclei ...... 24 2.6 The Active Galactic Nucleus engine ...... 28 2.7 Galaxies in the Local Universe ...... 31 2.8 ALMA and the Magellanic Clouds ...... 32
3 Star and Planet Formation 34 3.1 The initial conditions of star formation ...... 34 3.2 Young Stellar Objects ...... 38 3.3 Chemistry of star-forming regions ...... 48 3.4 Protoplanetary disks ...... 52
4 Stars and their Evolution 59 4.1Excellent (but una ributed) review paper found on the web The Sun ...... 59 h p://www.ifa.hawaii.edu/users/jpw/classes/alma/lectures/ALMA-science.pdf4.2 Millimeter continuum emission from stars ...... 62 4.3 Circumstellar envelopes ...... 64 4.4 Post-AGBSources ...... 66 2 4.5 Supernovae ...... 69 4.6 Gamma ray bursts ...... 69
5 The Solar System 71 5.1 Planetary atmospheres ...... 71 5.2 Asteroids and Comets ...... 73 5.3 Extrasolar planets ...... 76
6 Synergy with other Major Astronomical Facilities 78
1 1 INTRODUCTION 3
Figure 1: Left Sensitivity of ALMA, compared with some of the world’s other major astronomical facilities, for typical integration times of several hours. Right Angular resolution of ALMA, compared with otherSensi vity requirements major telescopes. The➞ collec ng area top of the band➞ 50 x 12m dishes shown for ALMA corresponds to the compact 150Resolu on requirements m configuration, and the➞ bottom 10 km baselines correspondsAtacama to the large array with 12 km baselines. For theWavelength requirements VLT, the solid line corresponds➞ very dry site to the seeing} limited case, and the dashed line to the diãraction limited case with adaptive optics. 3 a site that is high, dry, large and flat is required, and a plateau at high elevation (5000 m) in the Atacama region of northern Chile is ideal - hence the Atacama Large Millimeter Array (ALMA). Figure 1 show the sensitivities and resolving powers of some of the world’s major astronomical facilities. These define the “front line” of astronomical research, and any new large telescope opening up a new part of the electromagnetic spectrum should have comparable performance. ALMA will be far more sensitive than either existing millimeter arrays or single telescopes with bolometer arrays. Furthermore, single-dish telescopes can only survey as deep as the confusion limit imposed by the finite beamsize. ALMA represents a giant step for millimeter/submillimeter astronomy, placing it in a unique position at the front line of astronomical research.
1.3 The Atacama Large Millimeter Array The scientific objectives presented in the sections below lead to the high-level science requirements and technical specifications of the array. It is clear from HST observations of high-redshift galaxies that an angular resolution of at least 0.1 arcsec is required to image the features of star-forming regions in the early Universe, and an angular resolution of order 10 milli-arcsec is required to study the details of nearby protoplanetary systems. At the same time, high surface brightness sensitivity is required in order to image faint extended star-formation regions in our Galaxy and the total emission over nearby galaxies. Thus, a “zoom-lens” capability is called for, with movable antennas and the longest baselines extending to 10 km or more. The sensitivity requirement is given by the need to detect the most distant star-forming 1 INTRODUCTION 4
Figure 2: Zenith transmission as a function of frequency under typical atmospheric conditions at the ALMA site. The three curves correspond to the 25%, 50% and 75% percentiles. 4
galaxies, and to achieve adequate surface brightness sensitivity at high angular resolution. For constant system temperature, the noise in brightness temperature increases as the baseline length squared, and can rapidly dominate the signal, which is limited to a few tens of degrees by the emission physics. In the case of spectroscopic observations, the bandwidths are governed by the linewidths, not the receivers. Furthermore, in some of the millimeter/submillimeter windows, receiver performance is close to the atmospheric noise limits. Thus, for much of the spectral line research, the only way to increase sensitivity is to increase the collecting area. In particular, an angular resolution of < 0.100 can only be reached for thermal lines if the collecting area is increased by an order of magnitude over current values. The high sensitivity is also required for calibration purposes, and for image quality. The collecting area of an array can be enhanced by increasing the number of antennas, their size, or both. There are clearly several trade-oãs to be considered. Small antennas have higher precision, larger field of view, and their large number gives better image quality. The use of large antennas maximizes the collecting area, and reduces the number (and therefore cost) of receivers and the demands on the correlator. For ALMA, 64 antennas of 12-meter diameter gives a good solution. Excellent imaging capabilities are important to achieve many of the scientific objectives. ALMA should provide instantaneous imaging capability, with the large number of antenna pairs allowing complete uv coverage in snapshot mode. It should provide high fidelity imaging, and wide field imaging capability through the use of mosaicing techniques. The receiver bands should ultimately cover all the millimeter and submillimeter atmo- spheric windows. The atmospheric transparency at the high altitude of 5000m considered for ALMA is shown in Fig.2. In total, ten bands would cover the available spectral range, and the ALMA dewars will be built to accommodate all of them. Initially the four highest- priority bands will be provided. The system must be capable of high spectral resolution Science highlights (so far)
• Submillimeter galaxies
• Nearby galaxies
• Star forma on
• Protoplanetary disks
• Evolved stars
ADS list of ALMA papers at h ps://science.nrao.edu/facili es/alma/pubs
5 Lensed SMGs
1 minute integra ons! + spectral sweeping to get lines
Vieira et al. 2013, Nature
6 Figure 3: Figure 3: The cumulative redshift distribution of luminous, dusty starburst galaxies, as measured with different techniques. The SPT millimetre-selected sample, with redshifts directly determined from spectroscopic observations of the molecular gas in the galaxies, is shown in black. The existing samples of radio-identified starbursts2,14,17,18,27 with redshifts determined from rest-frame ultraviolet spectroscopy, are compiled in the blue distribution. The redshift distribution24 of millimetre-identified starburst galaxies in the COSMOS survey is shown in red/orange, though the majority of redshifts in this sample are derived from optical/IR photometry of the sources rather than spectroscopy, and therefore less certain. Sources at z<1 were removed from the previous samples of starburst galaxiestobettercomparetotheselectioneffect imposed on the SPT sample due to gravitational lensing. The distribution of redshifts for radio-identified sources is incompatible with the distribution for the samplepresentedinthiswork.Thismeasurement demonstrates that the fraction of dusty starburst galaxies at high-redshift is greater than previously derived and that radio-identified samples were biased to lower redshift than the underlying population. Figure 2: Figure 2: ALMA 3 mm spectra of 26 SPT sources. The vertical axis is observed flux density in units of mJy, with 30 mJy offsets between sources for clarity. Spectra are continuum-subtracted. The 7 strong CO lines are indicative of dust-enshrouded active star formation. The spectra are labeled by source and redshift. Black labels indicate unambiguous redshifts (18), with the subset in bold (12) having been derived from the ALMA data alone. Five sources labeled in blue(5)areplottedatthemostlikelyredshift of multiple options, based on the dust temperature derived from extensive far-infrared photometry. Three sources with no lines detected are placed at z =1.85,inthemiddleoftheredshiftrangeforwhichwe expect no strong lines, and labeled in red. Total integrationtimesforeachsourcewereroughlytenminutes. The synthesized beam size ranges from 7!!×5!! to 5!!×3!! over the frequency range of the search, which is inadequate to spatially resolve the velocity structure of the lensed sources.
13
12 A closer starburst
Herrera et al. 2012, A&A 8 Protostars are sweet! Glycoaldehyde (HCOCH2OH) in IRAS16293-2422
9Jorgensen et al. 2012, ApJ Le ers –4–
10 Fig. 1.— Spectra in the central beams toward the continuum peaks of IRAS16293A (upper) and IRAS16293B (lower). Fits from LTE models of the methyl formate (blue) and glycolaldehyde (red) emis- sion are overplotted. The purple line indicates the model fit to the possible ethylene glycol transition. The X-axis represents the frequencies in the rest frame of the system (i.e., corrected for the system VLSR of 3kms−1). The green line is an indication of the RMS level (13 mJy beam−1)representedbyaspectrum extracted from an off source position. Note the much narrower lines toward IRAS16293B which facilitate identification of individual features. Streamers in a planet forming disk?
Casassus et al. 2013, Nature
11 Fomalhaut’s debris disk
12 Boley et al. 2012, ApJ Le ers Fomalhaut’s debris disk
1998 2012
2003
13 1998 Δr/a ~ 0.1
14 AU Mic’s debris disk The Astrophysical Journal Letters,762:L21(5pp),2013January10 The Astrophysical Journal Letters,762:L21(5pp),2013January10MacGregor et al. MacGregor et al.
Figure 2. Left: the observed 1.3 mm emission from AU Mic. Center: the best-fit model (see Section 3.3). Right: the imaged residuals. Contours are drawn at 4σ 1 (120 µJy beam− )intervals. (A color version of this figure is available in the online journal.)
Table 2 positionally coincident with the outer belt center: ∆rcen ! Model Parameters 1.9AU(3σ ). Regarding the outer belt, the most notable result Figure 2. Left: the observed 1.3 mm emission from AU Mic. Center: the best-fit model (see SectionParameter3.3). Right: Description the imaged residuals. Best-Fit Contours 68% Confidence are drawn Interval at 4σis that the models strongly favor rising emission profiles with large, positive gradients: x 2.3 0.3. Models with the 1 Fbelt Belt flux density (mJy) 7.14 +0.12, 0.25 (120 µJy beam− )intervals. − ≈ ± x Belt radial power law index 2.32 +0.21, 0.31 standard assumption of x<0producesignificantresiduals, − (A color version of this figure is available in the online journal.) ri Belt inner radius (AU) 8.8 +11.0, 1.0 underpredicting the intensities at 1##–2## from the belt center. − ± ro Belt outer radius (AU) 40.3 +0.4, 0.4 Because of the steep increase in the emission profile, there is − P.A. Beltpositionangle( ) 128.41 +0.12, 0.13 only a weak constraint on the inner edge of the outer belt. The ◦ − Table 2 αbelt Belt spectral index 0.15 +0.40, 0.58 best-fit ri deviates from 0 at the 2σ level: the 3σ limit is positionally coincident with the outer− belt center:− ∆rcen ! ∼ Model Parameters Fcen Gaussian flux density (mJy) 0.32 +0.06, 0.06 ri ! 21 AU. 1.9AU(3σ ). Regarding the outer belt, the most notable− result ∆rcen Gaussian offset (AU) 0.71 +0.35, 0.51 2 2 − Parameter Description Best-Fit 68% Confidence Interval is that the modelsσcen Gaussian strongly variance favor (AU ) rising!5.9 emission (3σ limit) profiles with 4. DISCUSSION αcen Gaussian spectral index 0.35 +2.1, 4.5 large, positive gradients: x 2.3 − 0.3. Models− with the Fbelt Belt flux density (mJy) 7.14 +0.12, 0.25 We have presented new, subarcsecond resolution ALMA ∆α R.A. offset of belt center≈ (##)0.61+0.02,± 0.02 − x< − observations of 1.3 mm emission from the AU Mic debris disk x Belt radial power law index 2.32 +0.21, 0.31 standard assumption∆δ Decl. ofoffset of belt center0producesignificantresiduals, ( ) 0.03 +0.02, 0.02 − ## − − and analyzed the data with a simple parametric model. This ri Belt inner radius (AU) 8.8 +11.0, 1.0 underpredicting the intensities at 1##–2## from the belt center. − ± emission is resolved into two distinct components: (1) an edge- ro Belt outer radius (AU) 40.3 +0.4, 0.4 Because of the steep increase in the emission profile, there ison outer belt with an emission profile that rises with radius
Liu 2004 − 3.3. Results of Model Fits 15P.A. Beltpositionangle( ◦) 128.41 +0.12, 0.13 only a weak constraintMacGregor et al. 2013, ApJ Le ers on the inner edge of the outer belt. Theout to 40 AU, and (2) an unresolved peak at the center of the − outer belt. This distribution is more complex than the single, αbelt Belt spectral index 0.15 +0.40, 0.58 best-fit r deviatesThe best-fit from parameter 0 at thevalues and2σ theirlevel: 68% the uncertainties 3σ limit is − − i narrow ring often assumed for debris disks. However, it has determined from the marginalized∼ posterior probability distri- F Gaussian flux density (mJy) 0.32 +0.06, 0.06 ri 21 AU. some similarities to other nearby resolved systems, such as % Eri cen ! butions are listed in Table 2.Thedataandbest-fitmodelare − (Backman et al. 2009)orHR8799(Suetal.2009), that show ∆rcen Gaussian offset (AU) 0.71 +0.35, 0.51 compared in the image plane in Figure 2;therearenosignifi- − an inner component inferred from excess infrared emission, σ 2 Gaussian variance (AU2) 5.9 (3σ limit) cant residuals. The4. best-fitDISCUSSION model has a reduced χ 2 1.37 (905, cen ! = separate from an extended and colder outer belt. αcen Gaussian spectral index 0.35 +2.1, 4.5 920 independent data points, 12 free parameters). The model- − − We haveing presented procedure was new, performed subarcsecond on each SB individually resolution and ALMAthe 4.1. The Central Emission Peak ∆α R.A. offset of belt center ( )0.61+0.02,0.02 full data set (all four SBs together). The results were entirely ## − observations of 1.3 mm emission from the AU Mic debris disk ∆δ Decl. offset of belt center (##) 0.03 +0.02, 0.02 consistent, although the parameter uncertainties were notably The stellar photosphere is much fainter than the central peak − − and analyzedsmaller the from data thewith superior a SB-4 simple data set parametric alone, and we model. focus on Thisnoted in Figure 1.ANextGenstellarmodel(Hauschildtetal. those results. 1999)withTeff 3720 K, L 0.11 L ,andM 0.6 M emission is resolved into two distinct components: (1) an edge- = ∗ = ' ∗ = ' on outer beltMost with parameters an emission are determined profile with highthat precision.rises with We find radius(e.g., Metchev et al. 2005;Chenetal.2005)thatmatches good agreement of the outer belt parameters Fbelt, ri, ro with the AU Mic photometry from 0.4 to 25 µmcontributesonly 3.3. Results of Model Fits out to 40 AU,the less and well-constrained (2) an unresolved fits of Wilner peak et at al.{ the (2012 center), and} on of theF 52 µJy at 1.3 mm, 6 times fainter than observed. outer belt.the This disk distribution P.A. from measurements is more ofcomplex scattered thanstarlight the (e.g., single,However,∗ = AU Mic is an active∼ star that exhibits radio-wave The best-fit parameter values and their 68% uncertainties Krist et al. 2005). We measure a flat spectrum for the outer belt bursts. In quiescence, observations find <120 µJy at 3.6 cm narrow ring often assumed for debris disks. However, it has determined from the marginalized posterior probability distri- (αbelt 0) across the four basebands, which corresponds to the (White et al. 1994), and the contribution at 1.3 mm from hot some similaritiesdifference≈ tobetween other nearby the spectral resolved slopes of systems, AU Mic and such Neptune as % Ericoronal plasma seen in X-rays is unlikely to be significant butions are listed in Table 2.Thedataandbest-fitmodelare (Backman et(αNeptune al. 20092.1),)orHR8799(Suetal. consistent with data from 3502009µm to), 1.3 that mm show(though better spectral constraints are desirable; see Leto et al. compared in the image plane in Figure 2;therearenosignifi- (Wilner et≈ al. 2012). 2000). Flares are detected from AU Mic at 200–1200 µJy at an inner component inferred from excess infrared emission, ∼ cant residuals. The best-fit model has a reduced χ 2 1.37 (905, The central emission peak is detected with high confidence 6cm(Boweretal.2009), but this non-thermal emission is much separate fromat F ancen extended320 µJy (> and10σ colderbrighter outer than the belt. outer belt at that weaker at 1.3 mm. While the unknown variability makes any = = 920 independent data points, 12 free parameters). The model- location). It is unresolved, with cen 3.0 AU (3σ ), and extrapolation to 1.3 mm problematic, the temporal properties R ! ing procedure was performed on each SB individually and the 4.1. The Central Emission Peak full data set (all four SBs together). The results were entirely 3 consistent, although the parameter uncertainties were notably The stellar photosphere is much fainter than the central peak smaller from the superior SB-4 data set alone, and we focus on noted in Figure 1.ANextGenstellarmodel(Hauschildtetal. those results. 1999)withTeff 3720 K, L 0.11 L ,andM 0.6 M Most parameters are determined with high precision. We find (e.g., Metchev= et al. 2005;Chenetal.∗ = ' 2005)thatmatches∗ = ' good agreement of the outer belt parameters Fbelt, ri, ro with the AU Mic photometry from 0.4 to 25 µmcontributesonly the less well-constrained fits of Wilner et al.{ (2012), and} on F 52 µJy at 1.3 mm, 6 times fainter than observed. the disk P.A. from measurements of scattered starlight (e.g., However,∗ = AU Mic is an active∼ star that exhibits radio-wave Krist et al. 2005). We measure a flat spectrum for the outer belt bursts. In quiescence, observations find <120 µJy at 3.6 cm (αbelt 0) across the four basebands, which corresponds to the (White et al. 1994), and the contribution at 1.3 mm from hot difference≈ between the spectral slopes of AU Mic and Neptune coronal plasma seen in X-rays is unlikely to be significant (αNeptune 2.1), consistent with data from 350 µm to 1.3 mm (though better spectral constraints are desirable; see Leto et al. (Wilner et≈ al. 2012). 2000). Flares are detected from AU Mic at 200–1200 µJy at The central emission peak is detected with high confidence 6cm(Boweretal.2009), but this non-thermal∼ emission is much at Fcen 320 µJy (>10σ brighter than the outer belt at that weaker at 1.3 mm. While the unknown variability makes any = location). It is unresolved, with cen 3.0 AU (3σ ), and extrapolation to 1.3 mm problematic, the temporal properties R ! 3 Mass loss from an AGB star
Maercker et al. 2012, Nature 16 −20.0 −19.0 −18.0 −17.0 −16.0 −15.0 CO 3-2 channel maps 0 15 − −14.0 −13.0 −12.0 −11.0 −10.0 −9.0 R Sculptoris binary AGB star 0 thermal pulse; 15 15 −
−8.0 −7.0 −6.0 −5.0 −4.0 −3.0 t = 1800yr, Δt = 200 yr -3 M = 3x10 Msun 0 vej = 14.3 km/s 15 15 −
−2.0 −1.0 0.0 1.0 2.0 3.0 0 20 0 Offset (") 15 15 − 10 4.0 5.0 1006.0 7.0 8.0 9.0
0 0 200 15 15 Offset (") − 0.5 10.0 11.0 p.a. [degrees] 12.0 13.0 14.0 15.0 −10 0.4
0 300 0.3
15 15 −20 − 0.2 16.0 17.0 18.0 0 19.0 10 20.0 20 21.0 20 10 0 −10 −20 15 Radius (") Offset (") 0.1 0 0
15 15 Figure 2: The CO(J = 3 ! 2) emission at the stellar v of R Sculptoris. Left: Position − LSR −0.1 15 0 −15 15 0 −15 15 0 −15 15 0 −15 15 0 −15 15 0 −15 angle (p.a.) vs. radius based on the stellar vLSR image. The p.a. starts in North and Offset (") 17 increases counter-clockwise. Right: The stellar vLSR image. The color scale is given in Jy/beam. The green lines show linear fits to the emission peaks (white circles) in the p.a. vs. radius diagram. The first 2.5 windings are nearly! parallel, with a constant Figure 1: ALMA Early Science observationsseparation of ofthe 2.6” CO(±0.07”,J = 3 !indicating 2) emission that fromthe expansion the velocity has been constant (on AGB star R Sculptoris. The panels in averagethe figur) efor show the last the 2.5 emission binary periin theods. dif Theferent present -day expansion velocity is !1 velocity channels. The color scale is givenestimated in Jy to/beam. be 10.5 The km stellar s , giving velocity a binary is at period of 350 years. The linear fits can v =!19 km s!1. The numbers in the tophence-right be corntranslateders indicate directly the into velocity a velocity in evolutionkm s!1 (Fig. 3). The corresponding LSR spiral is plotted on top of the stellar v image (right). Deviations from a perfect with respect to the stellar velocity. The spherical detached shell appears asLSR a ring in spiral are on the order of ±1.5 km s!1, indicating small velocity variations over times the individual velocity channels, with its largest extent at the stellar velocity. The of " 50 years. Partial spiral arms and arcs connecting the third and fourth winding shell is clearly visible at 18.5” at the stellarshow av LSRlarger, as variationwell as a in spiral the wind structure velocity during these orbital periods on the connecting the central star with the detachedtimescale shell. of " 100The years.structure can be traced through all velocity channels.
16 0 500 1000 1500 −4.5 log(dM/dt) (solar masses/year)
14 −5
12 −5.5
10 −6 v(exp) (km/s) −6.5 8
−7 6 0 500 1000 1500 Time (yr)
Figure 3: The velocity and mass-loss rate evolution of the stellar wind around R Sculptoris. The solid and dashed lines show the velocity and mass-loss rate as a function of time, respectively. The points correspond to the expansion velocities of the emission peaks in the p.a. vs. radius diagram (Fig. 2, left), assuming a binary Fic on
18 Fact
19 What are you wai ng for?
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