Doc 6.1, ver. 1 | March, 2018

Observing with ALMA – A Primer(Cycle 6)

ALMA, an international astronomy facility, is a partnership of ESO (representing its member states), NSF (USA) and NINS (Japan),
together with NRC (Canada), NSC and ASIAA (Taiwan), and KASI (Republic of Korea), in cooperation with the Republic of Chile.
The Joint ALMA Observatory is operated by ESO, AUI/NRAO and NAOJ.

For further information or to comment on this document, please contact your regional Helpdesk through the ALMA Science Portal atwww.almascience.org. Helpdesk tickets will be directed to the nearest ALMA Regional Center at ESO, NAOJ or NRAO.

Version

Date

Editors

1

22 March 2018

G. Schieven

This document was produced by the National Research Council of Canada (Herzberg Astronomy & Astrophysics) with the National Radio
Astronomy Observatory, with contributions from staff of the European and East Asian ALMA Regional Centers (ARCs) and the Joint ALMA
Observatory (JAO).

In publications, please refer to this document as:
Schieven, G., ed., 2018, Observing with ALMA – A Primer, ALMA Doc. 6.1, ver. 1!2

Table of Contents
Purpose of this Document
Some Acronyms Used in this Document22

What is ALMA?
ALMA Regional Centers (ARCs)

3
4

What is Interferometry?
Science with ALMA

5
6

ALMA Full Operations Specifications

6

ALMA Cycle 6 Capabilities

7

Receivers and Sensitivities in Cycle 6
ALMA Correlators in Cycle 6
Did You Know? In Cycle 6 ALMA can…

8
8
9

Science During Cycle 6

10Before You Propose

10

Examples of Cycle 6 Observing With ALMA

12

Molecular Absorption Lines at z=0.9
Mapping a Lensed, High Redshift, Gas-Rich Galaxy
A Survey of Submillimeter Galaxies
Observing a GRB Afterglow (A Target of Opportunity)
Mosaicing the Nearby Spiral Galaxy M100
Multi-wavelength Continuum Survey of Protostellar Disks in Ophiuchus
Dust Polarization and Magnetic Fields in Star Forming Clouds
Observing Molecular Gas in a Planetary NebulaContinuum High Resolution Imaging of the Asteroid 3 Juno
Continuum Mapping of the Sun at Millimeter Wavelengths

Proposals, Observations and Data Reduction

12
13
14
15
16
17
18
19
20
21

22

Proposal Submission and Observing Process
Observing Considerations
Simulating ALMA Data
Using Single-Dish Data to Estimate ALMA Sensitivity Requirements
Creating Images From Your Data

22
23
27
28
29

Interferometry Concepts for ALMA: A Glossary of Terms
A Few Useful Equations

31
36

A Summary of “Quick Links”

37!1

Purpose of this Document
This document is designed to provide basic introductory information on the Atacama Large Millimeter/submillimeter Array (ALMA) and its capabilities, plus basic terminology and concepts related to radio interferometry.
Our goal is that, with all the basic information in one place, and a few examples of how to plan a science observation, this document can help all astronomers become familiar with ALMA’s capabilities and to start planning theirown ALMA observations.

Figure 1: A panoramic view of the ALMA antennas under the Milky Way. (Credit: Y. Beletsky (LCO)/ESO)

Some Acronyms Used in this Document
ALMA

Atacama Large Millimeter/submillimeter ArrayLAS

Largest Angular Structure

ACA

Atacama Compact Array (Morita Array)

NAOJ

National Astronomical Observatory of Japan

AOS

Array Operations Site (at 5000 m elevation)

NRAO

National Radio Astronomy Observatory

ARC

ALMA Regional Center

MRS

Maximum Recoverable Scale

CASA

Common Astronomy Software Applications

OSF

Operations Support Facility (at 2900m elevation)

CfP

Call for Proposals

OT

Observing Tool

DDT

Director’s Discretionary Time

SCO

Santiago Central Office, headquarters of the JAO

DSO

ALMA Division of Science Operations

SB

Scheduling Block

Early Science

SG

Science Goal

ESO

European Southern Observatory

S/N

Signal to Noise Ratio

FDM

Frequency Division Mode (spectral correlator mode)

SV

Science Verification

FOV

Field of View (or Primary Beam)

JAO

Joint ALMA Observatory

ES

TDM
TP Array!2

Time Division Mode (continuum correlator mode)
Total Power Array (part of the ACA)

Figure 2: An aerial photo of the ALMA
array, taken with a camera mounted on ahexacopter. Note the ACA 7-m Array inthe center foreground. (Credit: EFE/
Ariel Marinkovic)

What is ALMA?
The Atacama Large Millimeter/submillimeter Array (ALMA) is one of the largest multi-national science projects
to date. Located on the Chajnantor plain of the Chilean Andes at an elevation of about 5000 m and at a latitude of-23°, ALMA consists of the 12-m Array, made up of fifty 12 m diameter antennas, plus the Atacama Compact Array (ACA), also known as the Morita Array, made up of twelve 7 m antennas packed closely together (the 7-m Array) and four 12 m antennas (the Total Power or TP Array).
ALMA is a complete imaging and spectroscopic instrument operating at millimeter/submillimeter wavelengths,
providing scientists with capabilities and wavelength coverage which complement those of other research facilities of its era, such as the Jansky Very Large Array (JVLA), James Webb Space Telescope (JWST), and planned extremely-large-aperture optical and radio telescopes. ALMA enables transformational research into the physics ofthe cold Universe: regions that are optically dark but shine brightly in the millimeter portion of the electromagnetic spectrum. Providing astronomers a new window on celestial origins, ALMA probes the first stars and galaxies
and can directly image the disks in which planets are forming.
Unlike most radio telescopes, the ALMA antennas are located at a very highaltitude on the Llano de Chajnantor in northern Chile (see Figure 3), one of thedriest locations on earth. Decade-long monitoring studies of the sky above thissite have shown it to have the transparency and stability at submm wavelengths essential for ALMA (Figure 4). The site is large and open, allowing easy
re-positioning of the antennas over a region at least 16 km in extent (Figure 7).
Array operations are the responsibility of the Joint ALMA Observatory (JAO).
The telescope array itself is located at the Array Operations Site (AOS). Due tothe limited oxygen at this 5000 m elevation, the array is operated from the Operations Support Facility (OSF) at an elevation of 2900 m, with trips to the AOS
only to install, retrieve, or move equipment and antennas. OSF site facilities
include the array control room, offices, labs, staff residences, and a contractor
camp. The OSF is where ongoing operations, maintenance, and repairs of
ALMA antennas and receivers take place. The JAO has a central office in Santiago. The interface between the observatory and the global astronomical community is through the ALMA Regional Centers (ARCs; Figure 5).

Quick Links
Find a virtual tour of the ALMA site and vicinity at:
http://www.almaobservatory.org/en/multimedia-en/virtual-tour/
The current status of the ALMA observatory can be found at:
http://almascience.org/observing/alma-status-pag
Figure 3: Map of Chile, showing
the location of ALMA (red star).

Policies for using ALMA by the astronomical community are detailed at:
http://www.almascience.org/documents-and-tools/latest/alma-user-policies!3Figure 4: Curves showing the transparency of the atmosphereabove the ALMA site as a function of frequency. Plotted in blueand black are the transparency values for the 50th and 12.5th
percentile conditions respectively averaged over the year. Thismeans that 50% of the time the sky transparency is better thanshown in the blue line (corresponding to a precipitable water
vapor (PWV) of 1.3 mm), and 1/8 (12.5%) of the time is better
than the black line (corresponding to 0.5 mm of PWV). The
horizontal lines represent the frequency coverage of the ALMA
receiver bands. Bands 3 through 10 are available on all antennas in Cycle 6.

Figure 5: ARCs are located in North America, Europe,and East Asia. The ALMA headquarters are located in
Santiago, Chile. (Credit: ALMA (ESO/NAOJ/NRAO))

ALMA Regional Centers (ARCs)
Each of the three ALMA partners (Executives) maintains an ALMA Regional Center (ARC) within its respective region. The ARCs provide the interface between the JAO and their respective communities, either through the ALMA
Helpdesk or face-to-face at the ARC. In addition, the ARCs provide operational support to the JAO, and for researchand development activities in support of future upgrades to ALMA.
The North American ARC is part of the North American ALMA Science Center (NAASC) based at NRAO headquarters in Charlottesville, VA, USA. With the assistance of the National Research Council of Canada (NRC) and theAcademia Sinica Institute for Astronomy and Astrophysics (ASIAA) of Taiwan, the NAASC is responsible for supporting the science use of ALMA by the North American and Taiwan astronomical communities.
Researchers within the European Southern Observatory consortium are supported by the EU-ARC, based at the ESO
headquarters in Garching, Germany, along with regional nodes based in Germany, Italy, Sweden, France, the Netherlands, the United Kingdom, Portugal, and the Czech Republic.
The East Asian ARC (EA-ARC) is based at the NAOJ headquarters in Tokyo, in collaboration with ASIAA and the
Korea Astronomy and Space Science
Institute (KASI), and supports the
Quick Links
astronomy communities of Japan,
The three ARCs can be reached through the Science Portal or via their web sites:
Taiwan and the Republic of Korea.
Chilean astronomers may be supported by any of the three ARCs.
Similarly, astronomers from nonpartner countries may choose any of
the three ARCs for their support.

NAASC http://science.nrao.edu/facilities/alma/
EU-ARC http://www.eso.org/sci/facilities/alma/arc.html
EA-ARC http://alma.mtk.nao.ac.jp/e/ea-arc/!4

What is Interferometry?
In contrast to direct imaging, e.g.
with a CCD camera on an opticaltelescope, an interferometer samples
the power spectrum of the sky
ACA 7-m Array
Technical Building
brightness distribution; this is equivalent to measuring the Fourier transform of the sky. In a single integration (typically a few seconds or less),
each pair of antennas, called a baseline, samples a single point in thispower spectrum, at a position in
Fourier space related to the distance
between the pair of antennas and theposition angle of the baseline vector.
Antennas which are close together(short baselines) sample large-scale
angular structure, while long baselines sample very small-scale angularstructure. By combining these data,
called visibilities, over a large number
Figure 6: An areal view of ALMA in a relatively compact configuration. Theof baselines (the uv-coverage), themaximum distance between the 12-m Array antennas in the most compact conFourier plane is sampled, which canfiguration is about 150 m, and the minimum distance a mere 15 m. Indicatedthen be inverted (a Fourier transalso are the AOS Technical Building, and the ACA 7-m Array. (Credit: C.
form) to reconstruct an image. The
reconstructed image quality is very
Padilla/ALMA (ESO/NAOJ/NRAO))
sensitive to the uv-coverage — howcompletely the raw visibility data
covers the range of real angular scales on the sky. Even a few minutes of observations with the 43 12 m antennas (903 baselines) available in Cycle 6 provides good coverage. During longer observations, more of the uvplane is filled in by the rotation and foreshortening of baselines as the Earth rotates on its axis. Antennas are
periodically moved (reconfigured) to provide a wide range of array configurations with different baseline
lengths; observations with different configurations may be combined to improve the uv-coverage. Furthermore,
to recover very large-scale structure, the short spacings gap in the uv-coverage can be filled in by adding ACA
observations. For more detailed descriptions of these terms, see the Glossary starting on page 31.
Figure 7: In the most extended configuration of the ALMA array, antennas are scattered across the Chajnantor plateau as muchas 16 km apart. Here we see the array during the November 2015 long baseline campaign, when the maximum distance between
antennas was 10 km. Chajnantor means "place of departure" in the Kunza language of the Atacameño people that lived there andnamed this plateau. (Credit: S. Dougherty)!5Science with ALMA
Level One Science Aims for ALMA
While ALMA will revolutionize many areas of astronomy, the technical requirements of ALMA were driven by
three Level One Science Aims:
I.
The ability to detect spectral line emission from C+ in a normal galaxy like the Milky Way at a redshift of z =3, in less than 24 hours of observation.
The ability to image gas kinematics in a solar-mass protostellar/protoplanetary disk at a distance of 150 pc(roughly, the distance of the star-forming clouds in Ophiuchus or Corona Australis), enabling one to study
the physical, chemical, and magnetic field structure of the disk and detect the tidal gaps created by planets
undergoing formation.
III. The ability to provide precise images at an angular resolution of 0.1″. Here the term "precise image" means an
accurate representation of the sky brightness at all points where the brightness is greater than 0.1% of thepeak image brightness. This requirement applies to all sources visible to ALMA that transit at an elevation
greater than 20 degrees.
II.

ALMA Full Operations Specifications
Specification
Number of Antennas
Maximum Baseline Lengths
Angular Resolution (“)
12 m Primary beam (“)
7 m Primary beam (“)
Number of Baselines
Frequency Coverage
Correlator: Total Bandwidth
Correlator: Spectral Resolution
Polarimetry

50×12 m (12-m Array), plus 12×7 m & 4×12 m (ACA)0.16 - 16 km~0.2” × (300/ν GHz) × ( 1 km / max. baseline )~20.6” × (300/ν GHz)~35” × (300/ν GHz)
Up to 1225 (ALMA correlators can handle up to 64 antennas)
All atmospheric windows from 84 GHz - 950 GHz(with extension to ~30 GHz when Band 1 is deployed)
16 GHz (2 polarizations × 4 basebands × 2 GHz/baseband)
As narrow as 0.008 × (300/ν GHz) km/s
Full Stokes parameters

ALMA’s Breadth of Science
Other fields which ALMA will transform include:• Imaging the redshifted dust continuum emission from evolving galaxies at epochs of formation as early as z
= 10. The inverse K-correction, where the rising flux density on the Rayleigh-Jeans side of the spectral energy distribution of a dusty galaxy compensates for dimming at high redshift, makes ALMA the ideal instrument for investigating the origins of galaxies in the early universe, with confusion minimized by the high angular resolution.
• Using the emission from CO to measure the redshift of star-forming galaxies throughout the universe. The
frequency spacing between successive transitions of CO shrinks with redshift as 1/(1 + z), and the large instantaneous total bandwidth of ALMA makes it possible to carry out blind surveys in order to establish the star-forming
history of the universe, without the uncertainties inherent in optical and UV studies caused by dust extinction.
• Probing the cold dust and molecular gas of nearby galaxies, allowing detailed studies of the interstellar
medium in different galactic environments, the effect of the physical conditions on the local star formation history,!6

Figure 8: Nighttime panorama of ALMA (Credit: ESO/B. Tafreshi (twanight.org)

and galactic structure. The resolution of ALMA will reveal the kinematics of obscured active galactic nuclei andquasars on spatial scales of 10-100 pc and will be able to test unification models of Seyfert galaxies.
• Revealing the details of how stars form from the gravitational collapse of dense cores in molecular clouds. The
angular resolution of ALMA will enable the accretion of cloud material onto an accretion disk to be imaged and will
trace the formation and evolution of disks and jets in young protostellar systems. For older protostars as well pre-main
sequence stars, ALMA will show how planets and proto-planets sweep gaps in protoplanetary and debris disks.
• Uncovering the chemical composition of the molecular gas surrounding young stars. For example, establishing the role of the freeze-out of gas-phase species onto grains and the re-release of these species back into thegas phase in the warm inner regions of circumstellar disks. ALMA will have the large total bandwidth, high spectral resolution, and sensitivity needed to detect the myriad lines associated with the heavy, pre-biotic moleculesthat may have been present in the young Solar System.
• Imaging the formation of molecules and dust grains in the circumstellar shells and envelopes of evolvedstars, novae, and supernovae. ALMA will resolve the crucial isotopic and chemical gradients within these circumstellar shells, which reflect the chronology of the invisible stellar nuclear processing and early seeding of the ISM.
• Studying the physics of the Sun; refine dynamical and chemical models of the atmospheres of planets in ourown Solar System and provide unobscured images of cometary nuclei and hundreds of other Solar System objects.
• Plus countless other science goals, including unforeseen discoveries which always occur when exploring
new wavelength/sensitivity/resolution regimes.

ALMA Cycle 6 Capabilities
Capabilities offered for Cycle 6 include:• At least 43 antennas in the 12-m Array; plus ten 7 m and three 12 m antennas in the ACA• Complete spectral coverage at all frequency bands available from the ground from 84 to 950
GHz (3.6 mm to 300 µm), using receiver Bands 3 through 10 (see Figure 4)• Baselines up to 16 km for Bands 3 - 6; up to 8.5 km for Band 7; and up to 3.6 km for Bands 8, 9 & 10
• Single field interferometry at all receiver bands and mosaics at Bands 3 through 9• New! Full polarization capability, including circular polarization, at Bands 3 through 7 in both
continuum and spectral line mode
• Large programs requiring more than 50 hours of observations with the 12-m Array or morethan 150 hours with the ACA in stand-alone mode
• New! Stand-alone ACA observations may be requested in Bands 3 through 8• Solar observations and millimeter-wavelength VLBI at Bands 3 and 6• New! An extended Band 6 IF bandwidth will allow simultaneous 12CO, 13CO and C18O J=2-1
observations with broader spectral windows for galactic sources and nearby galaxies• See the Proposer’s Guide on the Science Portal for a full list of Cycle 6 capabilities!7

Receivers and Sensitivities in Cycle 6
Table 1 shows the receiver bands which are available for Cycle 6, and various properties. The sensitivities assume an integration time of 60 seconds, a continuum bandwidth of 7.5 GHz or a spectral resolution of 0.976 MHz, the “standard”
frequency for the Band (usually mid-band), and adopting the default weather conditions from the OT ALMA Sensitivity
Calculator, for a source at the zenith. See pp. 25 & 27 for a discussion of resolution and maximum scales with the ACA.

Table 1: Receiver Bands and Selected Properties
Cycle 6 Receiver Bands

Most Compact

Most Extended

Band

Frequency(GHz)

Wavelength(mm)

Primary
Beam(FOV; “)Continuum
Sensitivity
(mJy/
beam)

AngularResolution(“)

Approx.
Max.
Scale (“)
(see P.24)

Spectral
Sens.
∆Tline(K)

AngularResolution(mas)Approx.
Max.
Scale (“)
(see P.24)

Spectral
Sens.
∆Tline(K)

384-116

2.6-3.6

73-53

0.088

4.0-2.9

34-25

0.083

50-36

0.59-0.43537

4125-163

1.8-2.4

49-38

0.12

2.7-2.1

23-18

0.11

34-26

0.40-0.30726

5163-211

1.9-1.1

38-22

0.12

2.1-1.6

17.9-13.9

0.12

26-20

0.31-0.24727

6211-275

1.1-1.4

29-22

0.12

1.6-1.2

14-10

0.12

20-15

0.24-0.18754

7275-373

0.8-1.1

22-16

0.22

1.23-0.91

10.4-7.6

0.21

35-26

0.52-0.38260

8385-500

0.78-0.6

16-12

0.42

0.88-0.68

7.4-5.7

0.40

55-42

0.67-0.521029602-720

0.5-0.42

10-8.52.0

0.56-0.47

4.7-4.0

1.9

35-29

0.43-0.36477

10787-950

0.38-0.32

7.8-6.5

4.6

0.43-0.36

3.6-3.0

4.4

27-22

0.33-0.271144

Note: These sensitivities were calculated using the expected receiver temperatures at the time of writing, and may not represent
the values that are currently available. For the most up-to-date values, use the ALMA Sensitivity Calculator. To convert sensitivity in K to sensitivity in Jy/beam, see pages 28 & 36. Quoted angular resolutions are for sources which transit at the zenith.

ALMA Correlators in Cycle 6
The ALMA Correlators combine the signals of the individual antennas and are immensely powerful and flexible instruments.

Table 2: Spectral Capabilities per baseband forobservations in dual polarization

Each receiver outputs four 2 GHz-wide basebands in each
Polarpolarization, which are fed into the correlator. (See the Glos- Mode ization✤sary starting on page 31 for an explanation of unfamiliar
terms.) These basebands can be tuned independently of FDM
Dualeach other, so for example one might place the centers of
two basebands 1.8 GHz apart so as to get a contiguous
FDM
Dual
spectrum 3.6 GHz wide. Alternatively, one might tune
two basebands to the same frequency so as to completely
Dualoverlap, and use one baseband to focus on small sections FDM✤Note:

Single polarization modes are available for all bandwidths, which yield double the number of channels and half
the channel spacing. Full Stokes polarization mode yields
half the number of channels.
†Note: Resolution is 2 × the spacing due to a Hanning filter
applied to the data. Quoted resolution is at 300 GHz (1 mm).
‡Note: Because of filtering, the useful (effective) bandwidth of
this mode is 1875 MHz.
!8

Band
width
(MHz)

Nchan

Chan.
Spacing(MHz)

Spectral
Resolution†
300 GHz (km/s)

1875

3840

0.488

0.98

938

3840

0.244

0.49

469

3840

0.122

0.24

FDM

Dual

234

3840

0.061

0.12

FDM

Dual

117

3840

0.0305

0.061

FDM

Dual

58.6

3840

0.0153

0.031

TDM

Dual

2000‡

128

15.62531.2

of the baseband at very fine spectral resolution (small channel spacing) and the other baseband to look at a large velocity range with very coarse spectral resolution (large channel spacing).
Each baseband is sampled by the correlator according to a given correlator mode, defining the total bandwidth, number of channels, and spectral resolution. Currently a subset of the correlator modes is available for use. See Table 2(previous page) for the modes available in Cycle 6.
The channels within each baseband may be set up as one contiguous spectral window, or split up into two to four narrower windows. For example, one could choose to observe a contiguous 234 MHz range within a baseband, with 3840
channels of width 61 kHz (fourth entry in Table 2 above). If there were two lines of interest in that baseband, one
might choose to have two spectral windows in that baseband, each with 1920 channels covering 117 MHz, or one mightchoose to have four spectral windows in that baseband, each with 960 channels covering 58.5 MHz (see Figure 36).

Another feature of the correlators is the ability to apply spectral averaging prior to writing the data. Spectral
averaging allows one to reduce the number of channels in order to improve the S/N per channel, to lower thedata rate, and to reduce the sheer volume of data to manage. It is worth noting that, because the correlator
channels are Hanning smoothed before averaging, there is only a marginal decrease in spectral resolution
when an averaging factor of 2 is chosen, but with a factor of 2 decrease in the data rate and volume.

Did You Know? In Cycle 6 ALMA can…
…resolve molecular structures in M83:
create an HCN(J=1-0) mosaic of the full M83 bar with 30 pc resolution

in 2.17 hours…detect the ISM in high-redshift galaxies:
dust emission in a 1012 L⊙ luminous infrared galaxy out to z=10

in 21 minutes
HC…reveal the behavior of solar system objects:
trace the atmospheric water content of Venus using HDO lines

in 8.1 minutes

detect volatiles (HCN, Ch3OH, H2CO, CS and HNC) on active comets

in 4.7 hours…survey Galactic Clouds and star forming regions:
measure the polarization of dust in 30 protostars in a single star-forming region

in 1.2 hours…reveal the nature of planetary disks around nearby stars:
detect a dust disk gap induced by a Jupiter mass planet at 120 pc

in 2.1 hours…measure stellar activity from low-mass to high-mass stars:
investigate heating mechanisms of red giant stars

in 2 minutes…study black holes and their environments, near and far:measure black hole mass of NGC 4526 from molecular gas kinematics

in 47 minutes…trace the formation of galaxy clusters, cosmic structure:
characterize merger shocks in cluster gas with the Sunyaev-Zel’dovich Effect

Quick Links
To see more quick facts like these, as well as the calculations that went into them, check out
http://science.nrao.edu/facilities/alma/didyouknow!9

in 1.3 hours

Science During Cycle 6
From the start of science observing in lateQuick Links
September 2011 (Cycle 0), ALMA was alreadya powerful millimeter/submillimeter interFind a list of published ALMA papers at
ferometer, and is now approaching full operahttp://almaobservatory.org/en/science-at-jao/articles-a-publications
tion mode. By the call for Cycle 6 proposals,
nearly 1000 papers featuring ALMA data
have appeared in refereed journals. Some exciting recent results are shown in figures throughout this document. A
few example projects suitable for Cycle 6 have been compiled by the world-wide ALMA science staff and can befound on the previous page, or worked out in more detail on pages 12 to 21.

Before You Propose
In order to be able to submit proposals using the Observing Tool (OT), or request help from the Helpdesk, youwill need to be registered through the ALMA Science Portal.
When putting together an ALMA proposal using the OT, you should have at hand:• The “Proposing Guidance” page. This page (at http://almascience.org/proposing/learn-more/) provides a
guide for proposing to observe with ALMA, with links to a wealth of useful documents, including this Primer.
• A Science Case. The case must demonstrate how your proposed observations will address key scientific questions.
• Source coordinates, radial velocities, proper motions (for nearby sources), ephemerides (for solar system objects).
• Observing frequency, bandwidth, and spectral resolution. ALMA’s wavelength coverage is excellent, and you
can often get several molecular lines of interest in a single setup. However, setups targeting several specific lines
can sometimes be challenging, requiring careful positioning of the spectral windows in more than one Science
Goal, which strongly affects the time required. (See also spectral averaging, page 34.)
• Required angular resolution and largest angular structure (LAS). The angular resolution determines the largest
baselines needed in the array, while the largest angular structure of interest in the source determines the shortest
baselines needed (see Glossary, p. 31). Note that if you want very high angular resolution, you may end up resolving out larger structures of interest unless you use a combination of array configurations, which can be costly oftime. See page 23-24 for a discussion of angular resolution and maximum recoverable scale. The OT automatically determines the configurations you will need, as shown in Figure 28. The schedule of array configurations for Cycle 6
can be found in the Proposer’s Guide, available
from the Science Portal.
• Required mapping area. The field of view (FOVor primary beam; see Table 1 and page 33) is thearea of the sky that the array maps at a single
pointing. Maps larger than the primary beam
will require a mosaic, where the field centers arespaced by approximately half the primary beam
size. Note that the sensitivity of the primarybeam decreases with distance from the field center. It is strongly recommended that if the map
area of interest is larger than about 1/3 the primary beam, that you consider making a small (3Figure 9: (above) 360 degree panorama of the room containing the ALMA correlators. The correlators arehoused in the ALMA AOS Technical Building (below).
In the background, Chajnantor Volcano looms over theplateau bearing its name. (Credit: ALMA (ESO/NAOJ/
NRAO))!10

Figure 10: ALMA has produced astonishing, ultra-highangular resolution images of circumstellar disks aroundprotostars, showing gaps suggestive of being cleared out
by planets (HL Tau, upper left; TW Hydrae, upper right,
with the inset image showing the inner 10 AU of thedisk); multiple stars forming in a fragmenting disk (L1448
IRS3B, lower left); sweeping spiral arms indicating density waves in the outer circumstellar disk (Elias 2-27, lower
right), and (center) a dust ring (red) and gaseous spiralstructure (blue) in AB Aurigae.
(Credits: (upper left) ALMA Partnership et al., 2015, ApJ,
808, L3; (upper right) Andrews et al., 2016, ApJ, 820,
L40; (lower left) Tobin et al., 2016, Nature, 538, 483;
(lower right) Perez et al., 2016, Science, 353, 1519; (center) Tang et al., 2017, ApJ, 840, 32)

point) mosaic instead of a single pointing.
• Required sensitivity. Care must be taken when estimating the sensitivity needed per synthesized beam, particularly if estimating the source brightness from single-dish millimeter/submillimeter observations. A source
which is bright in, say, a 30” beam, may be difficult to detect in a 1” beam if the emission is spread out over a
few arcseconds. (Indeed, if smoothly distributed over the 30” beam, it may have a flux density only (1/30)2 as
bright.) See p. 28 for a discussion on using single-dish data toestimate required ALMA sensitivities.
Quick Links• Dynamic range needed. The dynamic range (ratio of brightest to
the faintest emission) in an image (or ratio of strongest to weakest
Resources for proposers, including the OT,
detectable features in a spectral line image (spectral dynamic
Technical Handbook, and other links can berange)) improves with the length of the observation and the number
found on the Science Portal atof baselines (which goes nearly with the square of the number of
http://almascience.org/proposing/antennas). Observations of fields where there is a bright nearby
source but where the faint emission is of interest (e.g. faint sources
near OMC1) will require great care, and may not achieve the theoretical noise level. (See p. 26.)
During Cycle 6 with at least 43 antennas, ALMA will have 903 or more baselines. Thus the array should be able toimage most “simple” fields with good fidelity. Imaging dynamic ranges are typically expected to reach up toaround 50-100 depending on band and configuration (extended configurations, or Bands 9 and 10 are likely to becloser to 50). In many cases it may be possible to achieve higher dynamic ranges using special calibration techniques or self-calibration. However, if your science aim involves imaging a complex field or one with a nearby brightsource, then the uv-coverage can be as important to consider as sensitivity. If you have a reasonable model for yoursource structure and brightness, an ALMA simulator (e.g. CASA simalma or the web based Observation Support
Tool) can be used to test and demonstrate the uv-coverage needed to achieve your imaging requirements. As an
example of what the CASA simalma or the Observation Support Tool can do, see page 27.
• Calibration and overhead. ALMA’s sensitivity is exquisite compared to earlier submm/mm instruments. However it is important to keep in mind that calibration and overhead can take up a significant amount of time, so thatsurveys of objects scattered across the sky are much less efficient than surveys of objects nearby in the sky which
can share calibration. The Time Estimate button in the ALMA OT can tell you the overhead time for your project.
!11

Examples of Cycle 6 Observing With ALMA
In the following sections we provide a few examples of observations that could be done with ALMA during Cycle 6.
Note that these examples use the sensitivities and capabilities of ALMA as they were known at the time of the Cycle 6 Announcement. Any astronomer proposing observations with ALMA should carefully check the publishedcapabilities and sensitivities at the time of the Call for Proposals.
For each example below, we start with a brief science aim and discuss the required receiver band at which the observations should be undertaken. Next, we determine the angular resolution and LAS needed. The necessary spectral resolution is discussed to provide the appropriate correlator settings (see Table 2 page 8). In addition, the continuum or channel
sensitivity is quantified so the on-source observing time can be calculated. Note that the examples may not include
overheads (calibration, telescope movement, etc.) which, for short observations in particular, can add a significant
amount of time to an observation. Users who are unfamiliar with these terms, or who would like a “refresher” on radio interferometry terms and concepts, should first read “Interferometry Concepts for ALMA” starting on page 31.

Molecular Absorption Lines at z=0.9
Science Aim: To study high-redshift absorption lines toward a bright background quasar
Observing gas in absorption against a bright background continuum source can provide detailed information on themolecular interstellar medium of the foreground galaxy. Observations of absorption lines are very powerful because
the detectability of the intervening gas depends only on the brightness of the background source. Many molecularspecies have already been detected in the few known intermediate-redshift absorbers, but ALMA's sensitivity andbandwidth will allow an unbiased survey for absorption lines in selected distant galaxies.
Here we prepare a spectral survey in Band 3 of PKS1830-211, where a spiral galaxy at z~0.9 is detected in front of thebright background quasar at z~2.5. The background source is sufficiently bright (~2 Jy in Band 3) that short observations with 12-m Array will result in very good optical depth limits of about 1%. By
covering a wide range of frequencies with several spectral settings, we expect to detectmany molecular lines, enabling detailed comparison with the interstellar chemistry of
the Milky Way interstellar medium.
Receiver(s): Band 3. The molecules HCN, HNC, HCO+, HOC+, CS, HC3N, for example, have transitions redshifted into this band.
Angular resolution: In principle, angular resolution is not important for this experiment because we are looking for absorption lines against a bright background point
source. However, an angular resolution of at least ~1” is needed so that the two lensed
components of the quasar can be separated. In the OT, we will select a range of angular resolutions, from 0.3” to 1”, and LAS of 0.5”.
Spectral resolution: We want to cover a large bandwidth but with adequate spectral
resolution to resolve the absorption lines of a few km/s width. The 1875 MHz bandwidth spectral (FDM; see Table 2) correlator mode gives a spectral resolution of ~3
km/s, but to decrease the data rate (and data volume), we will use a Spectral averaging
factor of 2, which still yields a resolution of ~3.4 km/s.
Line sensitivity: In order to reach 1% optical depth limits at 5 σ significance, an rms
noise level of 4 mJy per channel is required in Band 3. Additional spectral smoothingcan be applied to search for broader, weaker absorption features.
Continuum sensitivity: N/A: The continuum sensitivity of the spectral line observations
is sufficient for the present purposes given that a single spectral line channel has only a 1%
error in the continuum level. Continuum imaging would also be useful for self-calibration.
Band 3 observing time: Using the spectral scan mode, setting up such spectral surveys isvery simple. In the OT in the Spectral Setup frame, we request a start and end sky freFigure 11: Examples of molecuquency of 86.0 and 115 GHz respectively, which is nearly the entire frequency range oflar absorption lines detected tothe Band 3 receiver (84-116 GHz). The OT automatically sets up a set of five tunings toward PKS1830-211. These obcover the requested interval. To reach our target line sensitivity, a total integration time
servations were made by ALMAin late 2010 when there wereof 18 minutes (plus overheads of approximately 90 minutes) will be required to cover
only a few antennas.
this frequency range.
!12Mapping a Lensed, High Redshift, Gas-Rich Galaxy
Science Aim: To resolve the continuum andmolecular gas in a distant lensed starburst galaxy
At high redshift there is a population of gas-richstarburst galaxies that are relatively bright in thesubmillimeter, but extremely faint in the opticaldue to dust obscuration. The few that have been
observed with mm/submm interferometers areunresolved, but increasing numbers are now being discovered that are gravitationally lensed
and therefore brighter and larger than would
Figure 12: (a) SMA (very extended configuration) 870 !m image ofotherwise be the case. A pre-eminent example of
the Cosmic Eyelash. The red line marks the division between the twothis is the so-called "Cosmic Eyelash". This starimages of the background source. (b) Spectrum of CO(3-2) taken withburst galaxy at z=2.3 has been gravitationallylensed into two “images” (Figure 12) that have a
the Plateau de Bure interferometer. Both figures are from Swinbank et
combined extent of ~5". Each image has beenal. (2010, Nature, 464, 733).
resolved into at least four components (Figure12a) and large amounts of molecular gas havebeen detected (Figure 12b). Although resolved, the source is small and dominated by relatively compact structures and therefore well-suited to ALMA. As an example project, we will attempt to map both the continuum
and the molecular gas at high angular resolution where the ALMA array may be able to resolve the source
structure.
Receiver(s): Band 7 (spectral line and continuum, 312 GHz).
Angular Resolution and LAS: 0.1" (spectral line [CO(9-8)] and continuum, Band 7). This resolution is sufficient to spatially resolve the components revealed by the Harvard-Smithsonian Submillimeter Array (SMA),
and perhaps resolve individual lenses. Though the total extent of emission is ~5”, only structures of order 1”
are of interest so the LAS selected is 1”. Hence, only one 12-m Array configuration and no ACA observations
will be required.
Spectral Resolution: For Band 7, we use the continuum (TDM; see Table 2) correlator mode to provide 14 km/
s channels (i.e., 30 km/s resolution) and a total bandwidth of 1.875 GHz (~1800 km/s) in one baseband. The
remaining 3 can be used for mapping continuum.
Spectral (Band 7) Sensitivity: The peak flux density of the redshifted CO J=9-8 line is expected to be ~10 mJy.
Assuming the clumps are 0.2” in extent, the expected flux density at 0.1” resolution would be 10 mJy * (0.1”/0.2”)2 = 2.5 mJy/beam. To detect the line at a sufficient S/N across the entire width of the line, we aim toachieve a S/N~10 at the peak in a 100 km/s channel. This requires a 1 σ noise level of 0.25 mJy/beam. We use
the TDM correlator mode and average the channels to achieve the 100 km/s channel width. The ALMA sensitivity calculator with 43 12 m antennas predicts 43 minutes of integration to reach 1 σ = 0.25 mJy/beam (notincluding overheads). This amount of time will also produce a sensitivity of 29 µJy/beam in a map of the remaining 5.625 GHz of line-free continuum providing a S/N ~30 for the continuum signal (~1 mJy).

Figure 13: (left) ALMA Band 6 continuum
image of the Cosmic Eyelash showing showing
the main lensed components of the z=2.3 galaxy. (right) Continuum subtracted emission/
absorption spectrum of CH+ (1-0) at a rest
frequency of 835 GHz. These observations confirmed the existence of huge reservoirs of highly
turbulent gas in this and other high-redshift
galaxies. (Credit: Falgarone et al., 2017, Nature, 548, 430)
!13

A Survey of Submillimeter Galaxies
Science Aim: To measure accurately the positions of SMGs
A large fraction of the star formation activity at the epoch of galaxy evolution (1 < z < 3) is traced by submillimeter galaxies(SMGs). SMGs are typically detected with single-dish telescopes
at coarse resolution; identification of a counterpart has required
deep radio (centimeter) observations followed by deep opticalor near-infrared spectroscopy. With ALMA, we can precisely
locate SMGs very rapidly. In this example we lay out a strategy
to pinpoint a large number of sources with Band 7 continuum.
Receiver(s): Band 7 (Continuum, 345 GHz)
Angular Resolution and LAS: ~0.2” at Band 7. A good rule ofthumb is that 1” corresponds to ~ 8 kpc for z ~ 1, so spatiallyresolved observations can be made with ALMA. The ACA
and more compact 12-m Array configurations are not needed
for these observations since the sources are small. An LAS of
1” is selected.
Spectral Resolution: These are purely continuum detections,
so only the TDM mode (15.6 MHz channels) is required.
Continuum (Band 7) Sensitivity: In Cycle 6, up to 150 sources(see Proposers Guide for restrictions) may be observed ineach science goal. As an example, there are ~150 sources in the
COSMOS-AzTEC catalogue (Aretxaga et al. 2011) with (deboosted) flux densities > 3.3 mJy. Furthermore the sources are
Figure 14: The inset of this HST image galaxy
all within a few degrees of each other, making the observations
cluster Abell 2744 shows an ALMA image of thehighly efficient since the calibration can be shared among
z=8.38 gravitationally-lensed galaxy A2744_YD4.
many sources. Given the exceptional atmospheric conditions
At this distance, the universe was a mere 600 milat ALMA, we choose to pinpoint these sources at a higher frelion years old, and yet the ALMA observations show
quency (344 GHz or 0.8 mm) where they are significantly
that the galaxy is rich in dust. (Credit: ALMA
brighter (S α νβ, typically β~2 at z~2). Therefore, these sources(ESO/NAOJ/NRAO); NASA; ESA; ESO; D. Coeshould have S0.8mm > 5 mJy. It is possible that these sources(STScI)/J. Merten (Heidelberg/Bologna); Laporte et
may beal., 2017, ApJ, 837, L21)extendedor may resolve into more than one source, so we aim to get
a S/N ~ 40, adequate to identify the counterparts and obtain excellent relative astrometric accuracy, which is usuallyestimated as ~ θ/(S/N). We thus request a sensitivity of
0.125 mJy/beam.
Band 7 Observing Time: For Band 7, the ALMA sensitivity calculator, assuming 43 12 m antennas and an effective
7.5 GHz continuum bandwidth per polarization, predicts
2.7 minutes integration per source to reach a 1 σ = 0.125
mJy. Each science goal of 150 sources would require about
6.75 hours of on-source integration, not including overheads and calibration.
Figure 15: Four Milky-Way-like progenitor galaxies at redshifts
1.2-1.3, shown at left, have been found to be incredibly rich inmolecular gas. CO J=3-2 emission detected using ALMA isshown in red, superposed on images taken with HST. (Credit:
Papovich et al., 2016, Nature Astronomy, 1, Article no. 3)!14

Observing a GRB Afterglow (A Target of Opportunity)
Science Aim: Detect & monitor the mm/submm afterglow of a GRB from its burst to one month laterLong-duration gamma-ray bursts (GRBs) are among the most energetic phenomena in the universe. The broad-band
afterglow emission accompanying GRBs can be described quite robustly as non-thermal emission from electrons accelerated in relativistic shocks that occur as the GRB-generated outflow is decelerated by the ambient medium. At mm
and submm wavelengths, we can observe the synchrotron emission directly, where it is free from interstellar scintillation. Continuum monitoring at mm and submm wavelengths provides information on GRB physics, not only fornormal GRBs but also dark GRBs, and it sets constraints on theoretical models of GRB phenomena as well.
Only a few GRB afterglows have been detected at mm and submm wavelengths, because instruments at thesewavelengths have had limited sensitivity. GRB 030329, the second nearest (z=0.1685) GRB to date, was intensively monitored at mm and submm wavelengths. The GRB 030329 afterglow light curve at 100, 250, and 345 GHz
showed a “plateau” for about one week after the burst, followed by a steep decline (Sheth et al. 2003, ApJ, 595,
L33; Kuno et al. 2004, PASJ, 56, L1; Kohno et al. 2005, PASJ, 57, 147; Smith et al. 2005, A&A, 439, 981). These data
support a model with a two-component, jet-like outflow (Berger et al. 2003, Nature, 426, 145).
Receiver(s): Bands 3, 6, 7: (100 GHz, 230 GHz, and 345 GHz)
Angular Resolution and LAS: The sources will be unresolved, since they are about the size of a star. We selectan angular resolution of “Any” using the OT so that any configuration can be used.
Spectral Resolution: N/A. Each measurement will use the largest bandwidth possible in dual polarization.
Sensitivity and Observing Time: With 43 antennas and with 7.5 GHz of bandwidth in two polarizations, extraordinary sensitivities are possible in only a few minutes of integration time. For this project, however, observation cadence is required over a four week period. For proper sampling, the target must be observed every twoor three days for the first two weeks, followed by two additional epochs over the final two weeks. These can be
specified in the OT by selecting “Yes” for time constraints in the Control and Performance window, and specifying the required time constraints.
Triggering Target of Opportunity (ToO) Observations: As an observation target that can be anticipated but not
specified in detail, this would be an archetypal ToO project. The proposal must specify in detail the observing
modes, sensitivities, the observing cadence, and the trigger needed to initiate the observations. The trigger must
include an accurate position (within ~10”). During Cycle 6, there is no limit on the size of the time window, butwhether such observations are technically feasible will be judged on a case-by-case basis. See the Proposer’s
Guide, Appendix A.10, for details. Once the proposal is approved, full details on how to trigger the ToO observations will be sent to the PI.

Figure 16: Observations of the host galaxy of GRB 020819B. The cross marks the location of the GRB. (left panel) CO emissionand (center panel) dust continuum emission both of which were observed with ALMA; (right panel) visible light image captured
by Gemini North. Less gas was observed than expected in the vicinity of the GRB, and much more dust. (Credit: B. Hatsukade
et al., 2014, Nature, 510, 247)
!15

Mosaicing the Nearby Spiral Galaxy M100
Science Aim: Map the distribution and kinematics of molecular gas in a
nearby spiral galaxy
M100 (NGC 4321) is a bright, fairly face-on spiral galaxy in the Virgo Cluster.
As one of the nearest (d ~ 16 Mpc) spiral galaxies with well-defined arms
and active star formation, M100 has been studied at virtually every wavelength, including the millimeter. These studies reveal a rich molecular interstellar medium (ISM) fuelling the star formation visible at optical and IR
wavelengths. A bar funnels gas to the center of M100, leading to a nuclear
concentration of molecular gas and star formation while fainter molecularemission traces the spiral arms. ALMA imaged M100 as a science verification target and the data are available from the Science Portal. ALMA
casaguides provide a tutorial on combining data from the 12-m Array, the 7-m
Array, and the TP Array (see https://casaguides.nrao.edu/index.php/
M100_Band3).
ALMA’s excellent imaging fidelity, mosaicing capabilities, and sensitivity tolarge angular scale (via the ACA) make a wide-field, high-resolution imageof a galaxy like M100 a viable project. Such data would probe the dynamical
Figure 17: Integrated intensity
effects of spiral arms and the nuclear bar on molecular gas, map the density
map of ALMA 12CO (1–0) feathered
distribution of the molecular ISM, and allow comparisons of molecular ma(12 m+ACA) data of M100. (Credit:
terial to the distributions of dust, recent star formation, stellar populations,
Pan & Kuno, 2017, ApJ, 839, 133.)
This publication made use of Scienceand other phases of the ISM.
Verification data publicly available
Receiver(s): Band 3 (115 GHz): We consider a mosaic to observe CO (1-0)
from the ALMA archive.
emission (~115 GHz, Band 3) across most of M100.
Spectral Sensitivity: Large molecular clouds like the Orion complex havemasses ~5 × 105 M⨀ or more. We design our survey of M100 to detect suchclouds. For a CO-to-H2 conversion factor of ~2 × 1020 cm-2 (K km/s)-1 and assuming a CO line width of 10 km/s,
a 1σ sensitivity of 3 mJy/beam corresponds to a 1σ molecular mass sensitivity of 8 × 104 M⨀ per beam per 10km/s channel, enough to detect Orion-like clouds at S/N ~ 5 or more.
Angular Resolution: We target a resolution of 2”, or about 150 pc at M100. This is suitable for comparison to awide variety of multi-wavelength data and sufficient to resolve spiral arms, the bar, and large molecular complexes. At this resolution in Band 3 ALMA can achieve surface brightness sensitivity well matched to the brightness of M100 in a reasonable amount of time.
LAS: M100 exhibits structure over a large range of angular scales. We set the LAS to the extent of the map (270”)
to recover structure at all scales down to the resolution of the observations. All three arrays (12-m Array, 7-m Array, and TP Array) will be required.
Coverage: We know the overall extent of CO from previous observations, including wide field single dish maps.
A single ALMA beam covers only a small part of the galaxy, but a rectangular grid of 137 pointings spaced by~half of the primary beam does manage to encompass most CO emission. Mosaics of up to 150 pointings per
Science Goal (SG) are allowed by the Cycle 6 call.
Sensitivity and Mosaicing: Entering the mosaicing parameters and a required sensitivity of 3 mJy/beam directly into the OT, the integration time
calculator (which takes into account the overlapping of the mosaic fields)
estimates an observing time of ~2.3 minutes per mosaic position. With 137
fields we expect that the program will need about 5.5 hours on source, orjust over 9 hours with overheads. The OT estimates that 52 overlapping
pointings will be needed for the 7-m Array requiring about 75 hours total,
and a similar amount of time for the total power observations.
Figure 18: Moment 1 map of 12CO (1–0) data after combining 12-m, 7-m and TP Array data, showing the velocity structure of the nearby spiral galaxy, M100. (Credit:
ALMA (ESO/NAOJ/NRAO); casaguides (https://casaguides.nrao.edu/index.php/
M100_Band3))!16

Multi-wavelength Continuum Survey of Protostellar Disks in Ophiuchus
Science Aim: To investigate the evolutionary states of protostellar disks in the nearby Ophiuchus molecular cloud by measuring
the global variation in their dust spectral energy distribution (SED)
to infer dust properties.
The characteristics of dust in a disk around a protostar are expected
to evolve over time as dust grains settle to the disk mid-plane, accumulate onto larger solid bodies, and eventually, perhaps, form
planets within the disk. The evolutionary state of the disk may be
traced by determining the continuum SED of a coherent sample ofprotostellar disks. Wide frequency coverage is necessary to measure the SED of the disk. The SED is determined by a combinationof the dust temperature and density distribution, and the opticalproperties of the dust grains. Although in Cycle 6 ALMA will beable to resolve nearby protostellar disks in exquisite detail (see theimages in Figure 10), for this survey we will only measure global
properties but at great sensitivity in relatively short integrations.
Figure 19: ALMA observations showing theHere we discuss a potential project to observe a set of six Class II
rotation of the bipolar outflow emanating fromprotostellar disks in Ophiuchus, selected from the catalogue prethe disk surrounding a massive protostar in
sented by Evans et al. (2009, ApJS, 181, 321).
Orion, KL Source I. Green indicates the dust
Receiver(s): Bands 3, 5, 7, & 9: (98, 203, 344, & 679 GHz respectivecontinuum emission from the disk; blue andly)
red indicate dense gas (Si18O) emission movAngular Resolution & LAS: The typical size of a protostellar disk is ~ing toward/away from the line-of-sight respec100 AU. At the distance to the nearby Ophiuchus molecular cloud
tively. (Credit: ALMA (ESO/NAOJ/NRAO);(125 pc), 100 AU subtends 0.8”. In this experiment we aim to measure
Hirota et al., 2017, NatAstr, 1, 146)
global properties of the disks, including the SED. We therefore seek
observations at each frequency at 0.4”. Using the OT, we simply request a resolution of 0.4” for each frequency band. The observations would then be executed in more compact configurations for the higher frequencies, while Band 3 data would be observed in a more extended configuration. If necessary, we can fine-tune the angular resolution achieved in each band by applying a uv taper during the data analysis. It
should be noted however that such a taper suppresses the contribution from the longest baselines and changes thesensitivity. For this experiment, a configuration which gives an angular resolution of 0.4” (see Technical Handbook
Chapter 7) will have recoverable angular scales of about 3”, so most likely all emission will be recovered.
Spectral Resolution: These are continuum observations, so this is not relevant.
Sensitivity and Observing Time: We calculate the expected brightness of a typical disk assuming a disk mass of 0.01
M⨀, an average dust temperature of 20 K, a plausible dust emissivity, and a distance of 125 pc. The flux densities we
expect are 5.2, 20, 65, and 255 mJy for Bands 3, 5, 7, & 9, respectively. We plan our observations with the considerationthat the angular size of the disks is ~3 times the size of the synthesized beam at Band 3 (and thus the average flux density per beam is 1/9 the values quoted above) and that the edges of the disk
~10% of the average flux density. For a 3σ detection of the disk edges, we aim
for continuum sensitivities of 0.019, 0.074, 0.24, and 0.94 mJy/beam for theseven bands respectively. Using the ALMA sensitivity calculator with 43
main array antennas (12 m diameter) and 7.5 GHz of bandwidth, we find thatthese sensitivities can be achieved in 23 minutes per disk (Band 3), 2.5 minutes per disk (Band 5), 41 seconds per disk (Band 7), and 2.4 minutes (Band 9)
for each science target, not including overheads.
Figure 20: An ALMA image of the center of the Milky Way galaxy revealing 11 young
protostars within about 3 light-years of our galaxy’s supermassive black hole. The lines
indicate the direction of the bipolar lobes created by high-velocity jets from the protostars.
The star indicates the location of Sagittarius A*, the 4 million solar mass supermassive
black hole at the center of our galaxy. (Credit: ALMA (ESO/NAOJ/NRAO), Yusef-Zadeh
et al. 2017, ApJ, 850, L30; B. Saxton (NRAO/AUI/NSF))!17

Dust Polarization and Magnetic Fields in Star Forming Clouds
Science Aim: To detect magnetic fields in a star forming
core at the thermal Jeans length scale through dust polarization.
Stars form in giant molecular clouds under the influence of turbulence and large-scale magnetic (B) fields.
Theoretically, the significance of the B field influences
how structures are formed, such as the density contrast
within structures, the star formation rate, and the suppression of fragmentation. One method to probe the B
field is through observations of dust polarization. Dust
grains are known to align with their shorter axes parallel to the field lines in most circumstances. The planeof-sky projected B field integrated along the line ofsight can be traced by rotating the detected polarizationof the thermal dust emission by 90°. ALMA observations toward the star forming core W51 e2 can provide
information on B field orientations within the core with
unprecedented sensitivity and angular resolution. This
will provide information of how the B field influences
the formation of structures at the size-scale of the thermal Jeans length.
Receiver(s): Band 7 (343 GHz) For this experiment theBand 7 receiver (343 GHz) provides the highest sensitiviFigure 21: Magnetic field of the W51 e2 core at 0.9 mm
ty to polarized dust emission.
obtained at the Submillimeter Array. The total intensity
Angular Resolution & LAS: 0.2” & 0.8” To resolve theis shown in contours. The polarization intensity is
W51 e2 dense core at the thermal Jeans length scale, weshown in color scale. The star denotes the ultracompact
will need an angular resolution of 0.2” (1400 AU at 7
HII region. (Credit: Ya-Wen Tang, 2009 ApJ, 700, 251)kpc). The size of the core is 0.8”.
Continuum Sensitivity: 100µJy/beam The flux densityof the W51 e2 core is 9.3 Jy over the 0.8” core. At 0.2”
resolution, the flux would be spread over (0.8”/0.2”)2=16 beams. Hence the estimated ALMA flux density is
9.3/16=0.6 Jy/beam. Assuming that the polarization percentage is 1%, the expected polarization intensity willbe ~6 mJy/beam. We request a sensitivity of 100 µJy/beam to achieve a 60σ detection of the polarization overthe entire core.
Observing Time: It will take 5.0 minutes onsource with a 7.5 GHz bandwidth to achieve
the sensitivity. However, the OT imposes aminimum time of three hours for full-polarization observations in order to get sufficientparallactic angle coverage for calibration.

Figure 22: Multi-scale view of the magnetic
field around the protostar Ser-emb 8. The grey
scale shows the 870 µm (Band 7) continuum dust
emission from the JCMT (upper left), CARMA(lower left) and ALMA (right). These data showthat the magnetic field is chaotic at all scales
down to the resolution limit, and may only haveplayed a minor role in star formation in this region. (Credit: Hull et al., 2017, ApJ, 842, L9)!18

Observing Molecular Gas in a Planetary Nebula
Science Aim: To map the structure of molecular gas (CO) in a Planetary Nebula
In Planetary Nebulae (PNe) molecular gas is often observed in a torus
surrounding a core of ionized gas. The detailed structure of molecular
gas in PNe, however, is of great interest since it contains information on
the physical processes that created the nebulae. High resolution observations of a few PNe show that this molecular gas is characterized by ahigh degree of fragmentation. For example, the Helix Nebula has beenfound to be made of thousands of small (Diameter < 1"), dense (n ~ 105cm-3), quiescent (∆V < 1 km/s FWHM), and faint (peak TA* < 5 K)
clumps that are slowly evaporating in the radiation field of the central
white dwarf. The origin of these tiny clumps is still debated and until
ALMA began science observing the highest angular resolution millimeter molecular line observations had beam sizes greater than 3" (see Huggins et al. 2002, ApJ, 573, L55).
Receiver(s): Band 6 (230.538 GHz) For this project we choose to observe
Figure 23: Composite image of thethe CO (2-1) line.
pre-planetary Boomerang Nebula.
Angular Resolution & LAS: 0.3" Taking the Helix Nebula results as aBand 7 ALMA observations (orange)
starting point, the angular resolution desired should be below the fragshow the hourglass-shaped outflow,which is embedded inside a roughly
mentation scale of ~1" (which we set as the LAS). We attempt to gain around ultra-cold outflow. The hourfactor of one hundred in angular resolution over previous observations.
glass outflow stretches more than
Mosaic Required: The Helix is quite large (diameter ~ 25’) and highly
21,000 au, and is the result of a jet
fragmented. However, the diameter of the primary beam at 1.3 mm isthat is being fired by the central star,
only about 27". It would take an enormous mosaic of pointings to map
sweeping up the inner regions of thethe entire Helix and thus in this proposal one pointing each toward theultra-cold outflow like a snow-plow.
SE and NW portion of the nebula are chosen.
The ultra-cold outflow is about 10
Spectral Resolution: The spectral resolution is chosen to match the extimes bigger. The ALMA data arepected line profiles (~0.8 km/s fwhm) within the Helix Nebula. We
shown on top of an image from thechoose the 117 MHz bandwidth spectral mode and spectral averaging
Hubble Space Telescope (blue). (Credit:
factor 4 to give a resolution of 0.158 km/s.
ALMA (ESO/NAOJ/NRAO); NASA/
Channel (Line) Sensitivity: 0.5K. The fragments observed in the Helix
ESA Hubble; NRAO/AUI/NSF; Sahai
Nebula are quite faint, with peak < 5 K). In this scenario, a moderate
et al., 2017, ApJ, 841, 110)
sensitivity is desired (S/N ~ 5), which would detect the brighter Helix
Nebula fragments.
Observing Time: For Band 6, the ALMA sensitivity calculator,
assuming 43 antennas, 0.3” angular resolution, and an effective 0.16 km/s spectral resolution predicts 3.9 hours to reach
0.5 K. The two separate pointings will require about 7.8 hours
of ALMA observing time, plus overheads.
Figure 24: (left) HST image and (right) ALMA image of an expanding spherical shell and spiral wind emanating from the evolved
star R Sculptoris. Recent modelling suggests that the bifurcation of
the spiral-shell pattern (indicated by a white box) results from a highly
elliptical binary system. (Credit: ESA/NASA & R. Sahai. (Right)
ALMA image of LL Pegasi. Credit: ALMA (ESO/NAOJ/NRAO) /Kim et al., 2016, NatAs, 1, 60)!19

Continuum High Resolution Imaging of the Asteroid 3 Juno
Science Aim: To observe the 1.3 mm continuum emission on the Asteroid 3 Juno:
Submillimeter wavelengths are well-matched as probes of the thermal response of asteroid surfaces.
ALMA can resolve thermal emission from an asteroid's surface to measure the temperature distribution
across the surface and estimate the regolith thickness and composition. Tracking the asteroid over time
will reveal its rotational period and 3D shape.
Receiver(s): Bands 6: (233 GHz or 1.3 mm)
Angular Resolution & LAS: Near-infrared adaptive optics imaging combined with modelling puts thesize of Juno at ~240 km. To obtain at least eight resolution elements across the asteroid, a resolution of ~29
km is required which translates to 20 mas at the distance to Juno of 1.97 AU. Multiple resolution elements
will reveal variations in brightness temperature over the surface of the asteroid. With an angular size of~0.17”, Juno is smaller than the Maximum Recoverable Scale of 0.21” with the most extended configuration available in Cycle 6 (see ALMA Technical Handbook, Chapter 7), and so all relevant size-scales will berecovered.
Continuum Sensitivity: Juno has a single-dish flux density of 240 mJy at 250 GHz. At a resolution of 20
mas, there are ~64 synthesized ALMA beams across the 240 km (i.e. 168 mas) wide asteroid. Thus the flux
density per beam will be
240 mJy / 64 beams = 3.8
mJy/beam. The sensitivity required to reach asignal-to-noise ratio of
100 is then 3.8 mJy/beam/100 ~ 38 µJy/beam.
With a bandwidth of 7.5
GHz, dual polarization,and 43 antennas, this canbe achieved in 12 minutes of on-source integration time (but with approximately 33 minutesof overhead).
Observing Time: Juno’s
known rotation period isapproximately 7.2 hours.
Sampling at least 10
epochs of the orbital period should allow a detailed model of the 3D
shape of the asteroid.
These 10 epochs do notnecessarily need to beobserved
contiguously,
and indeed contiguous
observations longer than
2-3 hours are not permitted in Cycle 6 (see the
Figure 25: Asteroid 3 Juno was observed in October 2014 during the commissioning peProposer’s Guide
riod of the ALMA long baseline campaign. Over the course of the observations, theAppendix A.10), and canachieved resolution varied from 32 x 24 mas to 42 x 36 mas, differing mostly due tobe scheduled over severchanges in phase stability as the observations spanned the transition from night throughal days using the Timedawn and into daytime. The baselines used to achieve this resolution ranged from 26 m toConstraint feature in the
13 km. The observations spanned 10 epochs, covering 60% of the 7.2h rotation period.
OT.
(Credit: ALMA Partnership et al., 2015, ApJ, 808, L2)!20

Continuum Mapping of the Sun at Millimeter Wavelengths
Science Aim: To map the brightness distribution of the solar chromosphere at millimeter wavelengths:
Radiation from the Sun at millimeter and submillimeter wavelengths probes a fascinating and dynamic
layer of the solar atmosphere where non-radiative heating becomes manifest. Since the source function is
Planckian and these wavelengths are in the Rayleigh-Jeans regime, observations in these wavelengths offer a convenient linear thermometer with which to probe the thermal, spatial, and dynamical structure of
the chromosphere in a variety of physical regimes: the quiet Sun, coronal holes, active regions andsunspots, prominences and filaments. The example shown here was obtained during the December 2015
solar commissioning campaign.
Receiver(s): Bands 3 and Band 6: (100 and 239 GHz)
Angular Resolution & LAS: The ALMA primary beam is much less than the angular size of the Sun (e.g.,
24” vs ~1920” in Band 6). The Sun therefore fills the ALMA primary beam and its sidelobes with complex
emission. Dense snapshot uv-coverage is essential and, for accurate recovery of “absolute” brightness, total power mapping observations are also needed. To this end, Cycle 6 offers observations of the Sun thatuse a hybrid array comprised of both the fixed-station ACA 7 m antennas and the 12-m Array in configurations C43-1, -2, and -3, corresponding to angular resolutions of approximately 3.4”, 1.8”, and 1.2” inBand 3 and 1.5”, 0.8”, and 0.5” in Band 6, respectively. An angular resolution of 1” corresponds to ~730km on the Sun. Fields of several arcminutes are typically mapped through mosaicking techniques, withup to 150 offset pointings relative to a PI-specified solar ephemeris that provides offset pointing coordinates and optionally corrects for differential solar rotation. The largest angular scales are recovered using
fast-scan mapping techniques with the total power (TP) antennas to map the entire Sun with a total fieldof view of 2400”.
Spectral Resolution: These are continuum observations and so the observations were made in TDM
mode and default channelization is used.
Sensitivity and Observing Time: The source dominates the system temperature although, for solar mode
observing in the “active Sun mode”, the receiver temperature is ~1000 K. For imaging programs requiring“large” fields, mosaicking is required with each pointing receiving a few seconds integration time on
source. The sensitivity is limited by various systematic effects: time variability of the source, details of the
uv-sampling, as well as integration time and bandwidth. However, test data using 30 antennas yielded anrms of approximately 26 K (GHz-s)-1/2 for Band 3 and 68 K (GHz-s)-1/2 for Band 6. With Cycle 6 data, forwhich 43 antennas in the 12-m Array and 10 antennas in the 7 m array are offered, the sensitivity will besignificantly better.
Figure 26: Images obtained by ALMA of a large sunspot on 18 December 2015 at 3mm (Band 3; left) and 1.25 mm (Band 6; center).
Observations at shorter wavelengths (center) probe deeper into the solar chromosphere than longer wavelengths (left). Hence Band 6
observations map a deeper layer of the chromosphere that is closer to the visible surface (photosphere) of the Sun than Band 3 observations. The full map of the Sun (right) was taken at Band 6 using a single antenna of the TP array using a “fast-scanning” technique.
(Credit: ALMA (ESO/NAOJ/NRAO)!21

Proposals, Observations and Data Reduction*
Proposal Submission and Observing Process
Call for Observing Proposals: The general procedure is that the Joint ALMA Observatory (JAO) prepares the Call for
Proposals (CfP), which includes the anticipated capabilities of the observatory (available observing bands, correlator
modes, observing modes, configurations, etc.) for the upcoming observing cycle.
The CfP is broadcast to the regional and worldwide communities by the ARCs using standard broadcasting means(e.g. society and observatory newsletters and mailing lists), and is posted to the ALMA science websites. If you wish toreceive these notifications automatically, you can subscribe to the email distribution list of your ARC by submitting a request through the Helpdesk.
Registering as an ALMA User: All users who wish to be part of any ALMA proposal (either as Principal Investigatoror Co-Investigator), submit tickets to the ALMA Helpdesk, track a project, or retrieve proprietary data from the ALMA
Science Archive, must register as an ALMA user via the ALMA Science Portal (see link page 37). Non-registered users may
still access ALMA user tools, software, or archived data, or browse the Helpdesk Knowledgebase archive.
Users will be associated with one of the ALMA partners (EU, NA, EA, or Chile) or as being outside the partnershipbased on their institutional affiliation(s). This affiliation factors into time allocation and specifies the ARC that users will bedirected to for data retrieval and helpdesk support. Users from Chile or from non-ALMA member regions may select anyof the three ARCs for support.
Proposal Preparation & Submission (Phase 1):
Quick Links
After the CfP is issued, users have some period (1-2months) to prepare their Phase 1 materials using the
The Observing Tool (OT) is accessible from
ALMA OT. Phase 1 consists of a detailed observing
https://almascience.org/proposing/observing-tool
proposal with a scientific and technical justification
Instructional videos on using the OT can be found at
submitted to the Observatory through the OT. The OT
includes a sensitivity calculator and viewers for assisthttp://almascience.org/proposing/observing-tool/video-tutorials
ing with correlator setups and mapping parameters
while preparing Science Goals. Users can use the
ALMA Helpdesk, available from the ALMA Science Portal, to get assistance from ARC staff at any stage of the preparationand submission process (Phase 1 Support). Users should be very careful to input the correct information (source coordinates, frequencies, bandwidths, etc.) into the OT at Phase 1. Only very minor changes are permitted during scheduling blockpreparation (Phase 2), and significant changes require a major change request which can cause delay or may not be granted.
If desired, one can simulate ALMA observations using the simalma tasks of the CASA software package, or using the webbased ALMA Observation Support Tool. These tools take a model image as input and simulate the resulting ALMA image, accounting for the array configuration, instrumental noise, atmospheric phase delay, as well as the data reduction process. One
can also use the compilation of molecular spectral line databases provided by the Splatalogue on-line catalogue to help plan spectral line observations. (See page 39 for links to these tools.)
Proposal Review Process: Phase 1 submissions are peer-reviewed by a single international committee that is divided
into a number of science-themed review panels (Scientific Review). Priority status is awarded based on the proposal’s scientific ranking and available time. The time available for projects will depend on the PI’s institutional affiliation, with33.75%, 33.75%, 22.5%, and 10% made available toprojects associated with the North American, Europe,
Quick Links
East Asian, and Chilean partners, respectively. (Large
programs, i.e. those requesting more than 50 hours of
The spectral line catalogue tool Splatalogue is available
12-m Array time or 150 hours of ACA time (see Proat http://www.splatalogue.net/
poser’s Guide section 4.3) will have their time split
Information on observing with ALMA, including the Project
among the partner regions of the co-PIs.) Time will
Tracker (SnooPI) and allocated projects, can be found at:
also be available to projects from non-ALMA regions
http://almascience.org/observing/
through the Open Skies policy. Users should consult*For the non-expert, a list of Interferometry Concepts is included, starting on page 31. *!22

Figure 27: The image shows the measurement of the SZ effect in the galaxy cluster RX
J1347.5-1145 taken with ALMA (blue). The background image was taken by the Hubble Space
Telescope. A “hole” caused by the SZ effect is seen in the ALMA observations. (Credit: ALMA
(ESO/NAOJ/NRAO); Kitayama et al., 2016, PASJ, 68, 88; NASA/ESA Hubble Space Telescope)

the Proposer’s Guide available from the Science Portal for details. Users will receive
notification of the proposal review process via email. (See details on the proposal
review process at http://almascience.org/documents-and-tools/latest/alma-proposal-review-process.)
Scheduling Block Preparation (Phase 2): The highest-ranked projects (Grades A
and B, and some Grade C as required) then proceed to Scheduling
Block Preparation, also called Phase
Quick Links
2. During Phase 2, successful proThe Proposer’s Guide lists capabilities and restrictions for Cycle 6:
posers will have the opportunity to
http://almascience.org/documents-and-tools/latest/alma-proposers-guidereview their SGs and make minorchanges if necessary (but see below).
After this their Phase 1 Science Goals
will be converted into Scheduling Blocks (SBs), which contain all the instructions necessary to carry out the observations
specified in the SG, including calibrations, etc. SBs are then submitted to a scheduling queue so they are available to thearray operators when conditions are appropriate at the ALMA site, according to the ranked list of proposals, operations
schedule and weather conditions. Once SBs have been submitted, users can track the status of their project through the
ALMA Project Tracker (SnooPI), a user application available from the Science Portal (https://almascience.org/observing/
snoopi). Note that only very minor changes (e.g. source position change less than half the primary beam size) to the SBs
are permitted during Phase 2. Any significant change requires a major change request through the Helpdesk, and may
result in delays to implementing the SBs and/or may not be granted. See the Proposer’s Guide section 7.2 for details.
Archive & Data Delivery: After each SB has been successfully observed, the data will be calibrated, reduced and quality assured by ALMA/ARC staff and deposited into the ALMA Science Archive, available from the Science Portal, where
they may be retrieved by observers. ALMA data have a one-year proprietary period from the date when they are placed in
the ALMA Science Archive and made available to the PI. Archived data products include the raw and calibrated visibilities, telescope logs, relevant data reduction scripts, and reference images and cubes.

Observing Considerations
While considering a possible ALMA project, it is important to understand that ALMA is a very flexible instrument.
Data can be obtained over a wide
range of observational parameters: angular resolution,
Quick Links
field-of-view, spectral resolution, and sensitivity. These
A fuller description of ALMA observing considerations can be found inquantities must be specificallythe Technical Handbook
defined and justified for a
http://almascience.org/documents-and-tools/latest/alma-technical-handbook
given project in a proposal,and careful choices are required
to ensure that the project's scientific aims can be met. These quantities are also used during Phase 2, to guide in planningthe execution of the project. Depending on the nature of a given project, the observational parameters may be interrelated.
In the following, we describe the basis for choosing these parameters.
Angular resolution (or “synthesized beam”) is the minimum angular separation whereby adjacent spatial features canbe distinguished. Angular resolution fundamentally varies as the inverse of the product of the observing frequency anddistances between the antennas used to make the image; higher frequencies or longer antenna baselines result in data offiner angular resolution. An important concept to remember about interferometers is that they can only observe emission
on a discrete set of angular scales (i.e., spatial frequencies), as measured by the antenna pairs making up an array (see “uv!23

coverage”, page 35). Since the number of angular scales measured is finite, the resulting image is spatially "filtered" and onlyreflects the emission on the observed angular scales. Even for a given baseline distribution, however, the observer has some
control over the effective resolution of the image during post-processing. By using different weighting schemes to reconstruct an image, it is possible to make moderate tradeoffs between effective resolution and surface brightness sensitivity.
Maximum Recoverable Scale (MRS) is the largest angular scale structure that can be recoverable from observations by
an interferometer, and is defined to be 0.6 × (wavelength/minimum baseline) in radians (or ~124” × (1m/Dmin) × (300
GHz/ν)), where Dmin is the minimum
distance between antennas in metersand ν is the observing frequency in
Quick Links
GHz). MRS is a guideline for thelargest angular structure on which
A video explaining Maximum Recoverable Scales can be viewed atsome of the flux of a smooth struchttps://science.nrao.edu/opportunities/courses/alma-instructional-videos
ture can be reasonably recovered by
the interferometer. This rule-ofthumb applies to the size scale ofsmoothly varying structures in both dimensions. Smooth structures larger than ~1.0 × (wavelength/minimum baseline)
will be "resolved out" by the interferometer. This is the well known "missing flux" problem intrinsic to interferometry. The
minimum baseline depends both on the array configuration (i.e. compactness) and source elevation. To recover emissionthat has been "resolved out," additional observations are needed, including observations with more compact configurations (such as compact configurations of the 12-m Array and/or the 7-m Array) or large single-dish telescopes (e.g. the TP
Array). One can explore with CASA simalma or the Observation Support Tool whether the ACA will be required for a particular project. The OT shows the MRS for the most compact and most extended 12-m Array configurations in the Control
and Performance tab, and based on the requested angular resolution and largest angular structure (LAS), may recommend theuse of multiple 12-m Array configurations and/or the ACA (see Figure 28).
Note that if you are only interested in small details, faint extended emission of no interest can be ignored if it is of too low
surface brightness (a few times the rms) at the expected resolution. However, emission which is so bright that it would bestrongly detected were it not smooth, can cause artefacts in high resolution images if low-resolution data are not also used.
Field-of-view (FOV) is the area on the sky over which an interferometric image is obtained. The instantaneous FOV isformally the angular size of the half-power width of the Gaussian beam (FWHM) of the individual antennas and is also
called the width of the "primary beam". The size of the FOV depends on the inverse of the product of the frequency of theobservation and the diameter of the individual antennas used; larger antennas or higher frequencies result in smaller FOVs.
Larger FOVs and flatter map sensitivities across images can be attained by observing in series many adjacent locations on thesky (best separated by λ/2D where λ is the observed wavelength, and D is the diameter of the antennas, to achieve Nyquist
sampling) and using the resulting data to create a "mosaic" map. For a single pointing, the sensitivity of the observation is not
uniform across the FOV; it declines with angular separation from the center position with the approximately Gaussian responsivity of the main antenna beam. For this reason, if the map region of interest extends beyond about 1/3 of the primarybeam from the field center, it is strongly recommended that a small mosaic is made instead of a single pointing.
To have constant sensitivity across the mosaic, each pointing must be observed to the same relative sensitivity. Thus,
mosaics can be quite costly in terms of observing time. Deciding whether a mosaic or a single pointing should be observed requires an understanding of the expected source structure and size, i.e., whether or not the observed emission will
be extended, based on previous data from other telescopes. Furthermore, if multi-band images over the same FOV are
needed for a given project, mosaics may be required with higher frequency bands in order to match the area coverage of asingle pointing with lower frequency bands. Mosaics can also aid in recovering some emission on scales larger than thosethat are sampled by single pointings, though they cannot compensate for emission that has been “resolved out”. (SeeMaximum Recoverable Scale above.)
Spectral resolution is the minimum separation in frequency whereby adjacent independent features can be distinguished. The digitized data from ALMA allows for an incredible range in spectral resolution. (See Table 2 page 8.) Spectral resolution depends on how the correlator has been configured. ALMA's correlator can be configured to provide data
cubes with up to 8192 independent spectral channels (though see Data Rate on page 32 for a caveat). The width of thesechannels can be defined from 3.8 kHz to 31.25 MHz. (In Cycle 6, the smallest available channel spacing is 7.6 kHz (in single
polarization), with a spectral resolution of 15.3 kHz due to Hanning smoothing.) For continuum observations or for observations of very broad spectral lines (e.g. high-redshift galaxies), wide bandwidths and low spectral resolution channels!24Figure 28: This figure links Desired Angular Resolution to the Maximum Recoverable Scale (MRS). The desired angular resolution
determines the largest baselines (and hence the configuration) of the 12-m Array. The observations from that configuration can recover
structure over a range of angular scales; for example, an observation at 1” resolution at 100 GHz can recover structure over size-scales
from ~1” to ~11” (cyan region). Structures larger than the “Maximum Recoverable Scale” (MRS; see p. 24) are “resolved out” by theinterferometer. Structure larger than the MRS may be recovered either by adding observations using a smaller 12-m Array configuration, by adding observations using the 7-m Array, or by adding both another configuration and the 7-m Array. Still larger scale structure may be recovered for Bands 3-8 spectral observations by adding data from the Total Power (TP) array. To use this figure, you firstneed to scale your desired angular resolution and LAS by the factor (ν/100GHz), where ν is your desired frequency in GHz. For example,
suppose you wanted to observe at 300 GHz (Band 7) with angular resolution 0.4” and LAS of 10”. These become 1.2” and 30” respectively when scaled to 100 GHz. Reading from the figure above, two 12-m Array configurations plus the ACA 7-m Array would be needed to recover all angular scales. (The OT will automatically determine the required combination of arrays/configurations to achieve theresolution/LAS you require.) Note that adding configurations or the ACA will significantly increase the required observing time. See
Chapter 7 of the Technical Handbook and the Proposer’s Guide in the Call for Proposals for details and restrictions. E.g. angular resolutions finer than (0.21” x 100 GHz/ν) are not permitted for Bands 8-10 (Limit Bands 8-10 on figure above), and angular resolutions finer
than (0.096” x 100 GHz/ν) are not permitted at Band 7.
!25

are used to achieve high sensitivity; the total bandwidth of all correlator settings used cannot exceed 7.5 GHz. For observations of spectral lines, narrower bandwidths with higher spectral resolution channels may be required. There is, however, a cost to sensitivity in using small bandwidth channels. Sensitivity can be improved by averaging channels together,i.e., by the inverse square root of the number of channels averaged, but at the expense of the spectral resolution. Channel
averaging can be set up during Phase 1 in the OT, which has the added benefit of reducing the data rate, and later reducing
the sheer volume of data during reduction (see spectral resolution, page 24). (It should be noted that, because the correlator
channels are Hanning smoothed before averaging, there is only a marginal (~14%) decrease in spectral resolution when an
averaging factor of 2 is chosen, but with a factor of 2 decrease in the data rate and volume. In Cycle 6, the OT will select aspectral averaging factor of 2 by default.) The ALMA correlators are highly complex and extremely flexible and can be
configured to observe simultaneously several spectral lines within the 7.5 GHz band at high spectral resolution while additional correlator channels can be simultaneously used to observe continuum emission at low spectral resolution. In addition, a combination of high and low resolution correlator windows can be chosen over the same bandwidth to determine
how emission from lines at these frequencies is contributing to the emission observed at low spectral resolution. (See Baseband and Spectral Windows, pp. 31 and 34.)
Sensitivity is usually defined as the 1 sigma rms variation of noise in the data (∆S) and so serves as a threshold for the detection of emission. For ALMA, basic sensitivity depends on: the number of antennas; receiver performance; atmosphericconditions (i.e., water vapor content and other atmospheric gases with strong spectral lines in the submillimeter (e.g. ozone),
atmospheric turbulence, and target elevation); and, of course, integration time (see Useful Equations, page 36). Receiver andatmospheric conditions are quantified by one parameter called "system temperature" (Tsys). High Tsys values (in K) indicate
low sensitivity and vice versa. Note that atmospheric opacity and stability are strongly frequency dependent (see Figure 4),
and thus the ability to observe with any particular receiver will usually depend strongly on the weather conditions. These
conditions include the water content of the atmosphere which attenuates astronomical emission, and atmospheric turbulence
which results in phase instability. The magnitude of these problems generally increases with observing frequency.
Two other aspects of the observational set-up strongly affect sensitivity: spectral resolution and angular resolution.
Continuum intensities are often given in units of Janskys per beam where 1 Jansky (Jy) = 10-26 W m-2 Hz-1, while line intensities, particularly from single-dish telescopes, are sometimes given in units of Kelvin (K). (See Using Single-Dish Data to
Estimate Sensitivity Requirements, p. 28, or watch a 6-minute video available at https://science.nrao.edu/opportunities/
courses/alma-instructional-videos.) Converting from one unit to another requires knowledge about the angular resolution of the data, where the sensitivity in K is proportional to the sensitivity in Jy divided by the angular size of the beam
when the source is resolved. For a given ∆S, the corresponding ∆K increases with decreasing beam size; it is harder to detect extended line emission at high angular resolution. The quantity ∆S itself varies as the inverse-square root of the product of total integration time and the total bandwidth of the observation. (See page 36.) How data are weighted during
imaging also affects sensitivity. The total bandwidth of the observation is determined by the correlator settings and how
many spectral channels, i.e, resolution elements, are averaged together. For continuum data, a bandwidth of up to 7.5
GHz in each of two polarizations (effectively 15 GHz) can be used. Sensitivity also depends on the inverse square root of
the number of observed polarizations; all ALMA bands have two polarization channels. There is an online sensitivity calculator available (see tools p. 37), also built into the OT.
Dynamic range is the ratio between the peak flux in an image and the required rms. When making an image from interferometric data, the incomplete uv-coverage causes every source to be convolved with the array's point-spread function(dirty beam). This spreads sidelobes throughout the map and thus makes detecting faint sources challenging. The dirty
beam must therefore be deconvolved from each source (usually using the "clean" algorithm) but this is difficult for very
bright sources, depending on the uv-coverage and the accuracy of the calibration. Imperfectly deconvolved sources leave
residual sidelobe patterns in the cleaned image and prevent the theoretical image rms being reached and faint sources being detected. The ratio of the peak source flux in the field to the image rms is referred to as the imaging dynamic range andlarge values can be difficult to achieve. In Cycle 6 imaging dynamic ranges are typically expected to reach up to around 50100 depending on band and configuration (extended configurations, or Bands 9 and 10 are likely to be closer to 50). In
some cases it may be possible to achieve higher dynamic ranges using special calibration techniques or self-calibration.
Similarly, spectral dynamic range is the ratio of the brightnesses of the strongest and weakest detectable features in one
channel in a spectral line image, with obvious generalizations to smoothed or integrated data. Since the weakest detectablefeature is some multiple of the single-channel rms, the spectral dynamic range is often quoted as the ratio of the strongest signal to the channel-to-channel rms, much like the continuum case. In some cases the brightest signals will actually be continuum and the presence of very strong continuum emission limits the detectability of weak line signals for the same reasons that!26

are discussed in the case of the imaging dynamic range, though bright continuum also means that self-calibration can improve
the S/N. The spectral dynamic range is ultimately limited by the quality of the bandpass calibration, whose own fractional
noise (rms/continuum) is added in quadrature with that of the imaging data. Most spectral line emission experiments have
low S/N (5:1) and the spectral dynamic range is not limited by the bandpass calibration, but when the continuum is strong
the fractional noise in the bandpass calibrator can produce noise in a spectrum which overwhelms weak line signals. To reduce the fractional noise of the bandpass calibration it is often possible to smooth the bandpass calibrator spectrum.
Calibration and overhead are important considerations when observing. Interferometric observations at submm/mm
wavelengths are extraordinarily sensitive to atmospheric conditions which can vary over timescales of minutes to seconds
depending on frequency and other factors. Each execution of a scheduling block (see page 23 for a discussion of calibration and scheduling blocks) requires a significant amount of time, at least several minutes up to several tens of minutes, forcalibration alone. This can make surveys of bright objects, which may only require very short integration times, muchless efficient than surveys of objects nearby in the sky which can share calibration.

Simulating ALMA Data
Observing sources with a range of spatial scales (such as powerlaw envelopes) requires careful consideration of the effects of finite baselines and the respective uv-coverage as shown by the following analysis of reconstructed maps of complex regions. Thesesimulations were generated using the CASA task simalma (the online ALMA Observation Support Tool (OST) is also available). The
largest angular structure (LAS) is shown by the arrows in themodel image (top left image in Figures 29(a) and (b).
Figure 29(a) [above]: Demonstrated here is a simulated observation of asource with structure on very large angular scales. This theoretical model ofa protostellar disk has a power-law envelope, so in principle has power on all
spatial scales (top left). In this case, for a high-sensitivity observation, the
LAS is clearly as large as the field of view.
When observed at 345GHz with a desired resolution of 0.75", the 12-m Arrayonly ALMA observation (top center) shows that significant large-scale structure has been resolved out; this cleaned image has negative bowls, and a significant amount of restored flux is missing. The OT will recommend use of the
ACA in this case, and indeed the combined 12-m Array + ACA 7-m Array
image (top right) shows that more, but not all, of the large-scale structure isrecovered (incorporating the TP Array 12 m data will add in the missing flux
density). The bottom row shows the u-v coverage of the ACA (red) and
ALMA-array (blue) (left) and the amplitude versus u-v spacing for the ACA
and 12-m Array baselines (right). The correlated flux densities for the ACA
baselines are significantly larger than that for the shortest ALMA baselines.
Figure 29(b) [below]: Here we show simulated observations of a source
which has structure only on small spatial scales, such as a cluster of compact
galaxies or protostars (top left). The LAS is well within the range of theshorter 12-m Array baselines.
When observed at a desired resolution of 0.75" with just the 12-m Array(top center), we see no negative bowls in the cleaned image and the recovered
flux is within 95% of the total flux density. Adding the ACA (top right)
makes no significant difference.

Quick Links
The CASA simulator and
ALMA Observation Support Toolcan be used to simulate ALMA data
and are available at
http://almascience.org/tools/

Images!using!12)m!C2!array!with!a!resolu=on!of!!0.8”x0.7”!in!pa!80d!
!!!!!!!!!!!!!!!!!!!MODEL!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!12)m!image!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!12m+7m!Image!

LAS!
Restored!ﬂux!!!!!!!11000!Jy!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!7000!Jy!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!9000!Jy!
Primary!beam!corrected:!20%!cutoﬀ:!Contours:!!)20,20,50,100,200,300,400,600,800,1000,1200,1600,2000!!

U)V!coverage!!!!!!!!!!!!!!(red=ACA,!blue=ALMA12m)!!!!!!!!!!!!!!Amplitude!vs!uv)distance!!!!!!!!!

Images!using!12)m!C2!array!with!a!resolu=on!of!!0.8”x0.7”!in!pa!80d!
!!!!!!!!!!!!!!!!!!!MODEL!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!12)m!image!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!12m+7m!Image!

LAS!

Restored!ﬂux!!!!!!230!Jy!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!225!Jy!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!229!Jy!
Primary!beam!corrected:!20%!cutoﬀ:!Contours:!!0.03x)1,1,2,4,8,16,32,64,100,200,400!!

The bottom row shows the u-v coverage of the ACA and ALMA-array (left)
and the amplitude versus u-v spacing for the ACA and ALMA-array baselines (right).
In a case such as this, the OT will recommend not using the ACA.

U)V!coverage!!!!!!!!!!!!!!(red=ACA,!blue=ALMA12m)!!!!!!!!!!!!!!Amplitude!vs!uv)distance!!!!!!!!!

!27

Using Single-Dish Data to Estimate ALMA Sensitivity Requirements
Results from single-dish telescopes like Herschel, the JCMT, Nobeyama, IRAM, and so on, can greatly strengthen an ALMA
proposal, but converting those observations into a sensitivity estimation for ALMA is not always straightforward. Fortunately, if
we just need to use these single-dish results to estimate ALMA sensitivities, we can use some simple approximations.
Suppose your observation was presented in Jy/beam, which is the flux density of the source measured within the telescope’s beam. To estimate the flux density within ALMA’s usually much smaller beam, you’ll need an estimate of the angular size of the source. If the source is small compared to the ALMA beam, for example, then all of the flux will fit into thebeam. You now know the expected flux density for an ALMA observation, and you just need to decide on the signal to noise
ratio (SNR) you need.
If you know, or have reason to believe, that the source is bigger than the ALMA synthesized beam, then you will need tocorrect the measured flux by the ratio of the area of the source to the area of the ALMA beam. For example, suppose you
have a JCMT observation of a submm galaxy, and have measured a flux density of 32 mJy within the JCMT beam. From other evidence you believe that the angular size of the galaxy is about 2”, and you would like to observe it with ALMA at an angular resolution of 0.5”. That 2” wide galaxy would be spread over (2”/0.5”)2=16 ALMA beams in area. The expected flux
density in the ALMA beam would thus be only 32 mJy/16 = 2 mJy/beam. If you wanted a 10 sigma detection, you’d need asensitivity of 0.2 mJy/beam.
Observations from single-dish telescopes, particularly spectral data, are often quoted in units of Kelvins or K km/s,
which is the integrated intensity over a spectral line. The simplest, if not most rigorous, way to use an antenna temperature
to estimate an ALMA observation is first to convert the antenna temperature into a flux density using the point source approximation, usually written as:

2kTA = S⌫ Ae
where TA is the intensity in Kelvins, Sν is the flux density, k is the Boltzmann constant, and Ae is the effective area of the antenna surface of the single dish telescope, equal to the area of the antenna times the overall aperture efficiency (ηA), which is usually available from on-line documentation for the single dish telescope in question.
Armed with this approximation, we’re now ready to convert our integrated intensity into a flux density. Plugging in theappropriate values and constants, the conversion factor simplifies to:

S⌫ =

3614
TA Jy/K⌘A D 2where D is the diameter of the single-dish telescope in meters, and ηA is the aperture efficiency.
Let’s try an example. Suppose we had an IRAM observation of a protostellar envelope in rho Ophiuchus, and detected
the 1.6 km/s wide N2H+ line with an integrated intensity of 0.75 K km/s. IRAM has an antenna diameter of 30 m and,checking the IRAM website, an aperture efficiency of 0.6 at 93 GHz. Plugging these values into the point source approximation, we find that the measured flux density is 4.9 Jy km/s per beam. Let’s observe it with ALMA at an angular resolution of1”. What flux density would we expect in a 1” ALMA beam? As we found before, that depends on how large the source is.
In the worst case, the emission is spread evenly over the 26” IRAM beam, or over (26/1)2=676 ALMA beams. In this case theexpected flux density would be about 7.2 mJy km/s per beam over the 1.6 km/s line. Dividing by the line width, that’s
7.2/1.6 = 4.5 mJy/beam. If you just want to detect the line, you might choose to get a 5-σ detection with sensitivity 0.9 mJy/
beam. If you want to measure the line shape you’ll need narrower channels necessitating longer observing times.
Conversely if the source were only 1” or smaller in size, the expected flux density would be 4.9 Jy km/s per beam or
4.9/1.6=3.1 Jy/beam. For a 5-σ detection, you’d need a sensitivity of 0.61
Jy/beam. As you can see, some a
Quick Links
priori knowledge of source size
A video about estimating ALMA sensitivities can be viewed at
can be quite important when estimating the required sensitivity ofhttps://science.nrao.edu/opportunities/courses/alma-instructional-videos
an ALMA observation.
!28

Creating Images From Your Data
Once the data are taken, ALMA data will be reduced by staff from the JAO and ARCs using a pipeline employing the Common Astronomy Software Applications (CASA) package. After the data have been reduced and quality assured, the reduced data are ingested into the JAO archive and transferred to the ARC archives, where they are
made available to the project teams. The project teams are provided with the calibrated data, the reference images(i.e. the pipeline-reduced calibrated data cube), and the scripts used by the pipeline.
Once the pipeline-reduced data are released, the PI may still want to optimize the reduction and imaging to get
the best possible data for the project, most likely using CASA. In the following, we describe the basic concepts ofreducing interferometer data. This
process can be distilled down to
Quick Links
two stages, calibration and imaging, and we discuss these below
The CASA homepage
in turn. Casaguides are available
with links to documentation, downloads and tools is available at
http://casa.nrao.edu/
which will guide you step-bystep in reducing real ALMA SciCasaguides
ence Verification data.
are step-by-step guides for learning data reduction and imaging with CASA
http://casaguides.nrao.edu/
Calibration: ALMA observing is heavily constrained byweather conditions on the Llanode Chajnantor. Therefore, ALMA projects are divided up into blocks of time (Scheduling Blocks or SBs) that can be
executed dynamically by the on-site array operators when appropriate conditions are available. These blocks contain observations of well-characterized, typically bright objects (calibrators) either before, during or after the target
source observations are made. The calibrator data will be used to calibrate the target data during post processing.
Target data require calibration of their amplitudes and phases as a function of frequency and time. Flux calibration requires the observation of at least one or two sources of known flux density and angular extent. The brightnessof these objects should vary only relatively slowly, so that an accurate estimate of their flux densities can be determined. Typically, bright Solar System objects (planets, moons, asteroids) which are unresolved or only partially resolved for the required array configuration, and which have accurate models, are used as flux standards. The observed data from these objects can be used to scale accurately the intensities recorded from the target. Bandpass calibration also requires observations of bright sources with the same correlator setup as the target. This is used to correct for frequency-dependent variations in amplitude and phase. Phase (or gain) calibration corrects for phase errors
due to differential atmospheric changes above
each antenna on timescales of seconds tominutes, and for amplitude fluctuations as
well. Atmospheric water vapor can causevery rapid fluctuations over timescales ofseconds. Water vapor radiometers (WVRs)
are installed on all antennas on the 12-m Array to monitor these variations, which areroutinely applied to the data during pipeline
processing. Longer timescale phase drifts aremonitored using periodic observations of amoderately bright, very compact source at
Figure 30: Using ALMA, a forest of molecularlines from at least 19 different molecular species
was detected in the nucleus of the starburst galaxy, NGC 253. (Credit: ESO/J. Emerson/VISTA;ALMA (ESO/NAOJ/NRAO); Ando et al., 2017,
ApJ, 849, 81; Acknowledgment: Cambridge Astronomical Survey Unit)!29

relatively small anguQuick Links
lar distance from theLearn the principles of interferometry.
target.
The best
sources for phase cali- Lectures and tutorials from the 2014 Synthesis Imaging Workshop (NRAO) are found at:
bration are unresolved
https://science.nrao.edu/science/meetings/2014/14th-synthesis-imaging-workshop/at the angular scales
probed by the array;since such objects are point sources, their data have intrinsically zero phase (no emission at any angular offsets), andany phase changes recorded in the data are due only to changes in the system and/or atmosphere. The cadence atwhich the phase calibrator will be observed will depend on the stability of the atmosphere, the observational frequency, and the maximum baseline length. Atmospheric phase varies more rapidly at longer baselines, while athigher frequencies the variations have larger magnitude. The ALMA antennas will need to "fast-switch" between thetarget and gain calibrators from every few minutes to every few seconds (depending on frequency and baseline
length) to capture these variations.
If a science observation has a relatively bright and reasonably compact source in it, it may be possible to apply
self-calibration to the data, correcting for phase and amplitude variations during the observations. This powerful
technique must be applied carefully, but can result in an increase in effective sensitivity by large factors. It works
because the large number of baselines provides redundancies in the solutions which can be used to improve thetemporal gain calibration, leading to better source images. (For more information, see the casaguides.)Imaging: ALMA datasets will be processed through a reduction pipeline so that project teams will be able to see
preliminary results quickly. Non-standard and more complex projects may have their data calibrated by members of
the DSO and the ARCs using a combination of the pipeline and by hand using CASA. There are many approaches toreducing and imaging interferometer data but here are the basics.
The heart of imaging is a Fourier Transform (FT) of the interferometer data (termed "visibilities") into
images. The reduction process itself is threefold: first, the data are calibrated, then poor quality data must be
flagged and removed from the ensemble before the FT, and finally, the image is made from the inverse FT. Flagging is used to remove poor quality data, which might affect image quality, or atmospheric spectral lines in calibrator data which may skew the calibration solutions. Flagged data can be ignored by the reduction software andare then effectively removed from the data ensemble.
Data that have gone through flagging and calibration are ready to be imaged through an FT of the ensemble.
Images need to be large enough to cover the field of view of ALMA, which varies with frequency, and sampled
finely enough such that the structures observed at the high angular resolution of the data can be accurately represented. Various angular and spectral frequency weights can also be applied to the data during the FT to emphasize certain characteristics. For example, resolution and sensitivity can be traded-off by weighting the data in various ways; ”natural" weighting, where data are weighted relative to the number of angular scales observed in theensemble, typically provides the highest sensitivity (but at a slightly reduced resolution). In addition, spectral
channels can be averaged prior to the FT to improve sensitivity. The resulting image may include significant artifacts, depending on the complexity and brightness of the target region, the uv-coverage, and the amount and quality of data obtained; such images are sometimes called "dirty images". Since interferometers cannot measure all
angular frequencies, there will be gaps in the data ensemble that will translate into image artifacts after the FT.
Even dirty images of point sources have these artifacts. The workaround to deal with these artifacts has been tomodel the data through various deconvolution techniques. A common algorithm is called "CLEAN". It works byiteratively subtracting low-amplitude versions of the "dirty beam" from the dirty image, starting at the brightest
part of the dirty image and working down in intensity until only a residual image is left. The dirty beam is an image of a theoretical point source observed with the same uv-coverage as the actual data normalized to one. Cleaning typically continues until the flux density in the residual image is a small multiple of the noise in the dirty image but other thresholds are possible. The sky locations of the beam subtractions, called "clean components," aresaved. The clean components are placed on a blank image, and these are all convolved with a Gaussian of size
equal to the inner part of the dirty beam, i.e., a "clean beam". Finally, the residual image is added to the convolved
component image to produce a "clean image". There are many approaches to deconvolving images; even Clean
has many variations, but this is the basic idea. Of course, data will need to be deconvolved one spectral channel ata time, and this can be quite time consuming if the images are large or if there are many channels with emission.
!30

Interferometry Concepts for ALMA: A Glossary of Terms
Aliasing According to the Nyquist principle, aliasing occurs when a signal (the uv-plane) is under-sampled (uvcoverage), shifting higher angular frequency components to lower angular frequency. These “aliased” components introduce false large-scale structure into the resultant image, i.e the “dirty” image. Aliasing artifacts canalso be introduced by strong sources (including natural and human-generated) outside of the primary beam.
Angular Frequency See Spatial Frequency.
Angular Resolution (or Synthesized Beam) The effective angular resolving power (equivalent to a point spreadfunction) provided by the ensemble of transformed visibilities given its range of spatial frequency coverage. Essentially this is roughly proportional to λ/Lmax, where λ is the observed wavelength and Lmax is the size of the array’s largest baselines projected in the plane of the sky. (Note that this means that an observation at low elevation
has a larger synthesized beam than an observation at high elevation.) The angular resolution can also be adjusted
through the choice of an imaging weighting function (see Imaging page 29). A typical observation will result in asynthesized beam with a primary feature that can be approximated by a Gaussian whose FWHM is typically given as the achieved high resolution of the image or cube.
Array An ensemble of antennas where signals measured by each antenna are cross-correlated with signals fromall others to obtain data of high angular resolution. A homogenous array consists of antennas of the same diameter, like the 50 x 12 m antennas of the ALMA 12-m Array. A heterogeneous array consists of antennas of different
diameters, like the ALMA 12-m Array plus 7-m Array.
Band The emission frequency/wavelength range over which a given receiver is able to detect astronomical signals. For example, ALMA Band 3 is sensitive to astronomical emission over the range of 84-116 GHz (i.e., 2.6-3.4mm). See Table 1 on page 8 for the full list of receiver bands that ALMA will have in Cycle 6.
Bandwidth The subrange of frequencies in a given band over which data are obtained in a given observation.
For example, a 3.75 GHz bandwidth can be sampled in each sideband over the 84-116 GHz range of Band 3.
Baseband A baseband is a 2 GHz wide portion of the available signal (effectively 1.875 GHz because of filters between the receiver and correlator) which isdigitized at the antenna. Up to four 2 GHz
wide basebands are delivered to the ALMA
correlators (see Figure 31). For the dual polarization receivers (e.g. Bands 3-8), up to
two basebands can be placed in each sideband, or all four in one sideband. (Double
sideband receivers such as Bands 9 and 10
have the same basebands in both sidebands
simultaneously, though in Cycle 6 in limited
circumstances it will be possible to separate
out the signal from the unwanted image
sideband. (See Sideband)) The user-selected Figure 31: A graphical view of basebands and sidebands. Basebands may becorrelator configuration determines how tuned to overlap if the user wishes, or may be located so as to maximize themany basebands are ultimately used, where total bandwidth (as shown). Each baseband may be further subdivided into as
they are placed in the available IF range, and many as 8 spectral windows. Up to four spectral windows per baseband will
which correlation products are produced be available during Cycle 6.
(single, dual, or full polarization).
(See
Spectral Window; Sideband)
Baseline A pair of any two antennas in the array. The angular frequency that a given baseline measures is related
to the instantaneous foreshortened distance between the two antennas relative to the source and to the wavelength of the observed emission. An array of N antennas will have N(N-1)/2 baselines, so the 50-antenna 12-m
Array will have 1225 baselines. In Cycle 6, at least 43 antennas provide 903 or more baselines in the 12-m Array.
Clean Image or Cube A deconvolved image (or cube of images), where the emission in each has been modelledin some manner so that distortions induced by secondary features to the synthesized beam are minimized. The!31

Figure 32: The ALMA transporter carrying one of the 7m Array antennas to the OSF for servicing. (Credit G.
Schieven)

optimal method of deconvolution depends on thescience goals of the observation. (See page 30.)
Configurations There are 192 antenna foundations (pads) distributed over the Chajnantor
plateau. The pattern of the antennas on a subset ofthese pads is called a configuration, and is designed
to provide as uniform a uv-coverage as possible fora given angular resolution. An extended configuration provides a very small synthesized beam
though at the expense of larger scale structure, andvice versa (see Figure 28). See the Proposer’s
Guide for the configuration schedule for Cycle 6.
Correlator A powerful computer which cross-correlates the amplified, down-converted signals from each antenna
pair to produce the interference measurement (i.e, the “visibility”) from that pair. A user-selected correlator mode(see Table 2 page 8 for the available modes in Cycle 6) defines the bandwidth and resolution of the spectral window.
Data Rate The rate (in MB/sec) that data are fed from the correlator into the archive. With the large number ofbaselines and the huge number of channels (up to 3840 in each of four basebands), the data rate can reach unsustainable levels. Unless the science requires such high data rates, proposers should look for ways to lower the
rates, for example by using Spectral Averaging.
Dirty Image or Cube The dirty image is produced by the appropriate Fourier transform of the measured visibilities. A single image is produced from a given window if all channels are combined (e.g., through averaging,summing, etc.). A cube is the ensemble of images, typically ordered in velocity or frequency, where visibilities
from each channel have been Fourier transformed independently of those from other channels. The image or
cube is considered "dirty" because the secondary sensitivity features of the synthesized beam have distorted thelocation and brightness of the true emission distribution, producing unphysical artifacts. Essentially, the dirty image is the convolution of the true brightness distribution with the synthesized or dirty beam.
Dynamic Range This is the ratio between the brightest emission in an image and the rms noise of the image.
(Spectral dynamic range is the ratio between the brightest flux density and the rms noise within a spectral channel.) Large dynamic ranges (≳100 for Bands 3-8, ≳50 for Bands 9 & 10 in Cycle 6 can be difficult to achieve becauseof limitations of interferometry (see p. 26), but may be enhanced in some cases using self-calibration.
Field of View (FOV) See Primary Beam.
Fringes See Visibilities.
Image Sideband See Sideband.
Largest Angular Structure (LAS) The largest scale structure of interest in
the source to be observed. If the LAS is larger than the maximum recoverable scale which the array can recover, then a more compact array configuration, e.g. with the ACA, may be required. A 4-minute video explaining LAS can be viewed at https://science.nrao.edu/opportunities/
courses/alma-instructional-videos.
Figure 33: An open cryostat being shown off to guests at the ALMA Inaugurationin March 2013. The demonstrator is pointing to the feedhorn of a Band 3 receiver.
During Cycle 6, all antennas will be equipped with receivers covering Bands 3
through 10. (Credit: J. Hesser)!32

Figure 34: Some of the residents of the OSF
and environs. (top row) Vicuña; gecko;
mother and young burro; Andean condor;
flamingos; (bottom row) culpeo (Andeanfox); viscacha (member of the chinchilla family); viscacha. (Credit: G. Schieven)

Maximum Recoverable Scale (MRS) The maximum angular scale structure that may be recoverable from observations with an interferometer, and is dependent on the minimum separation between antennas; MRS~0.6 λ/Dmin.
Larger structures are “resolved out” and cannot be recovered. See Total Power, the discussion on page 23-24, andFigures 28 & 29.
Nyquist Sampling This is the minimum sampling interval needed to preserve the signal content without introducing aliasing errors. For ALMA, the Nyquist sampling rate for mosaicing fields is of order λ/2D, where λ is theobservational wavelength and D is the antenna diameter.
Polarization The ALMA receivers measure the two directions of linear polarization separately. When later processed by the correlator, the polarizations may be combined to give greater sensitivity, effectively “doubling” thebandwidth (dual polarization). Alternatively, one polarization may be discarded in order to double the number ofchannels over the requested bandwidth, yielding much greater spectral resolution. A third option is to use the polarization information to obtain the full Stokes parameters in order to measure the polarization of the source. A
limited full Stokes capability, including circular polarization, will be available in Cycle 6; see the Proposer’s Guidein the Call for Proposals for details (link p. 37).
Point Source Sources which are much smaller than the synthesized beam are often called point sources. Because
the source is not resolved, the visibilities will be the same whether the antennas are close together or far apart;hence point sources can be observed with any configuration of the antennas.
Primary Beam The angular sensitivity pattern on the sky of each individual antenna in the array, i.e., the sensitivity to emission relatively close to their pointing direction. The primary beam is typically approximated by aGaussian of FHWM equal to ~1.2(λ/D), where λ is the observational wavelength and D is the antenna diameter.
Parabolic radio antennas can have significant secondary angular sensitivities called sidelobes or the error beam,
but these can be minimized by careful design and construction. The primary beam sets the field of view (FOV) foran observation with the array, unless a larger mosaic is made.

Figure 35: Twilight panorama of ALMA in a compact configuration. (Credit: Y. Beletsky/ESO)!33

Receiver The instrument at each antenna in the array where astronomical signals are collected. The signals are
combined with a highly accurate frequency signal at each antenna (the local oscillator) to produce a lower frequency (downconverted) signal that can be handled more effectively by array system electronics (e.g., amplification or transmission).
Self-Calibration Self-calibration is the use of a bright source to solve for the relative gains of the individual antennas in phase (and, optionally, amplitude). A minimum of three antennas is required to self-calibrate phase; fourantennas are required to self-calibrate amplitude. In effect, the data are compared to an input model and the observed phases are corrected to reproduce the model as well as possible. For self-calibration to work, however, thedata themselves must be fairly well characterized, i.e., they must have high S/N over a wide range of angular frequencies.
Sideband At any given tuning, each receiver is sensitive to two separate ranges of sky frequency of equal width
called sidebands (see Figure 31). The four available basebands can be placed in one sideband or distributed between
the two sidebands; however, baseband numbers 0 and 1 must be placed in the same sideband as must basebands 2
and 3. In Bands 9 and 10, the receivers provide no inherent separation of the sidebands, so each baseband contains
signal from both sidebands simultaneously. By default, the signal from one sideband is rejected by a techniquecalled LO offsetting. Alternatively, the signals from both sidebands can be separated and recorded using a techniquecalled 90º Walsh switching, which will be available in some restricted cases in Cycle 6. (see Proposer’s Guide).
Shadowing Partial eclipsing of one antenna by another. When observations are made of sources at very low elevation, there is a potential for antennas which are close together to “shadow” one another, i.e. one antenna is attempting to look partially through another. This is particularly a problem for the ACA when observing at highnorth or south declinations (low elevations). Visibilities involving the shadowed antenna may need to be flagged
as bad when reducing the data, resulting in a lower sensitivity than expected.
Snapshot A short-duration set of integrations of an astronomical source using all baselines. Since only a limited
number of angular frequencies is sampled (see uv-coverage) the resulting image quality can be relatively poor, unless the number of baselines is large.
Spatial Frequency The inverse of an angular distance scale on the sky. In Fourier analysis, any distribution ofemission can be decomposed into information over a set of such spatial frequencies. Low spatial frequencies
equate to large angular scales and high spatial frequencies equate to small angular scales. The uv-coverage is the
sampling of the spatial frequencies by the interferometer.
Spectral Averaging Averaging of spectral channels in order to reduce the number of channels, to improve the S/
N per channel, and to lower the data rate. Within the OT’s Spectral Setup frame, one can choose a spectral averaging factor of 1 (no averaging), 2, 4, 8, or 16, which reduces the number of channels by averaging that number ofchannels together. Note that, because the correlator applies a Hanning filter to the data
before averaging, choosing a spectral averaging factor of 2 only marginally changes the
spectral resolution, while halving the data
rate and subsequent volume of data duringcalibration and imaging. The OT default is touse an averaging factor of 2.
Spectral Window A spectral window is afrequency subrange of a baseband. Eachbaseband may be divided into one or more
spectral windows by allocating a fraction of
the correlator resources to each window (seeTable 2 page 8). The properties of the spectral
window depend on the fraction of the correlator resources allocated to it. A single window will have the maximum bandwidth andminimum channel spacing, while multiple
spectral windows provide reduced band-

Figure 36: Four spectral windows were set up to observe these four lines
simultaneously in one baseband using the OT. The windows are each 234
MHz wide (with 0.43 km/s channels), and are centered at 345.80, 345.45,
346.11 and 345.09 GHz. The spectral windows must fit within the 2
GHz wide baseband.
!34

Figure 37: ALMA Band 6 observations of the remnant of an explosive event some 500 years ago in Orion OMC1, showing streamers ofhigh velocity gas shooting from a central point. The colours represent the relative velocities of the gas, from most blue-shifted (-150km/s) to most red-shifted (+145 km/s). (Credit: ALMA (ESO/
NAOJ/NRAO); Bally et al., 2017, ApJ, 837, 60)

width per spectral window and/or larger channel spacings.
Spectral windows can be placed anywhere within the ~2
GHz wide baseband. In Cycle 6, each baseband may containup to four spectral windows (with the same correlator modein each window within a baseband).
Synthesized Beam See Angular Resolution
Total Power, or Zero-Spacing Flux The large-scale emission which the array cannot detect (see Maximum Recoverable
Scale). A pair of antennas cannot be physically separated bya distance less than the antenna diameter. Hence in any observation there is a range of low spatial frequencies (from 0
to the lowest spatial frequency sampled by the array) where emission has not been sampled, or has been "resolved
out" by the array. Emission at large scales (low spatial frequencies) can be restored to images by combining arraydata with those from single-dish telescopes and/or a more compact configuration of antennas. For example, datafrom an extended 12-m Array configuration can be combined with data from smaller 12-m Array configurations,
and/or with the ACA 7-m Array and/or the 4-antenna TP Array (Fig. 28).
Track A long-duration set of integrations of an astronomical source using all baselines. As the Earth rotates, the instantaneous foreshortened distances between antennas change. Obtaining integrations over different hour angles, i.e., "tracking
the source," thus allows visibilities over a larger number of spatial frequencies to be measured (uv-coverage) and the resulting images more accurately reflect the actual emission distribution (assuming zero noise and perfect calibration).
uv-Coverage The breadth of spatial frequencies sampled during an interferometric observation, so named because"u" and "v" are the spatial frequency counterparts to angular distances in Right Ascension ("x") and declination ("y")
respectively (see visibility). Since an interferometer can only sample a finite amount of spatial frequencies, the ability
to reconstruct the true sky brightness from an interferometric observation increases with the uv-coverage. Images
made from data of low uv-coverage (snapshots) tend to have more secondary features due to aliasing, whereas images made from data of high uv-coverage (tracks) tend to have less aliasing. The number of points in the uv-coverage of a given frequency channel is the number of integrations times N(N1), where N is the number of antennas.
Visibility An interferometric observation of a source made at a specificspatial frequency. The ensemble of (calibrated) visibilities is what is
Fourier transformed to produce an image. Correspondingly, visibilities
are complex numbers with amplitudes and phases that are related to thebrightness and position of the emission relative to the position where theantennas are pointed. These amplitudes and phases need to be calibrated during observations by observing bright sources of known flux density and position. Visibilities are sometimes referred to as fringes. Eachcorrelator channel produces its own visibilities. The ensemble of visibilities is the uv-coverage.
Walsh Switching See Sideband.
Zero Spacings See Total Power.
Figure 38: This view shows several of the ALMA antennas and the central regionsof the Milky Way above. The image was taken during the ESO Ultra HD (UHD)
Expedition. (Credit: ESO/B. Tafreshi (twanight.org))!35

A Few Useful Equations
Converting Units: In radio astronomy, one is often converting between different units of measurement and computing the required integration time for varying spectral resolutions. Here we provide a few important reference
equations. (For further details, see "Tools of Radio Astronomy" by Rohlfs and Wilson.)
To convert between frequency and wavelength, a handy rule-of-thumb is to remember that the wavelength is ~1 mm(to three decimal places) when the frequency is 300 GHz. Thus, to convert frequency (in GHz) to wavelength (in mm),
λ (mm) = (300 GHz)/(ν GHz)To achieve a particular spectral resolution in velocity units ∆v at a given observing frequency ν requires a spectral
resolution in frequency units ∆ν of!
For example, a 1 km/s resolution at 300 GHz would require a resolution of 1 MHz. Similarly a resolution of
0.0153 MHz (see Table 2 page 8) would correspond to a resolution of 0.0153 km/s at 300 GHz, or 0.0051 km/s at
900 GHz.
For a gaussian source, the conversion from Rayleigh-Jeans temperature T to flux density Sν with synthesized
beam solid angle ,s is!
An alternate formulae that is often useful is
T
1K
!

=

⇣

S
1 Jy beam

⌘h
13.6
1

300 GHz 2
⇥

⇣

100
max

⌘⇣

100
min

⌘i

The noise ∆Sν, in an integration time ∆t, varies with system temperature Tsys, spectral resolution ∆ν, number ofantennas used N, diameter of the antennas D, and number of polarization measurements obtained np, in the following manner:!
As described on p. 28, in simple cases one can estimate a flux density from an antenna temperature measurementfrom a single-dish telescope with the following:

S⌫ =

3614
TA Jy/K⌘A D 2where TA is the source intensity in K, D is the diameter of the single-dish telescope in meters, and ηA is the aperture
efficiency.

!36

A Summary of “Quick Links”
ALMA Science Portal

http://almascience.org/

Call for Proposals

http://almascience.org/proposing/call-for-proposals/

Proposer’s Guide

http://almascience.org/documents-and-tools/latest/alma-proposers-guideTechnical Handbook

http://almascience.org/documents-and-tools/latest/alma-technical-handbook

Cycle 6 Science & Software Tools

http://almascience.org/tools/

Proposing Guidance quick links

http://almascience.org/proposing/learn-more/

Current ALMA Status

http://almascience.org/observing/alma-status-pag

ALMA Archive and Data

http://almascience.org/alma-data/

JAO Public Web

http://almaobservatory.org

ALMA Users’ Policies

http://www.almascience.org/documents-and-tools/latest/alma-user-policies

Common ALMA acronyms

http://www.almaobservatory.org/en/about-alma/acronyms

North American ARC

http://science.nrao.edu/facilities/alma/

European ARC

http://www.eso.org/sci/facilities/alma/arc.html

East Asian ARC

http://alma.mtk.nao.ac.jp/e/ea-arc/

Intro to Interferometry

https://science.nrao.edu/science/meetings/2014/14th-synthesis-imaging-workshop

Reducing Data with CASA

http://casa.nrao.edu/

Simulating ALMA Data

http://casaguides.nrao.edu/index.php?title=Guide_To_Simulating_ALMA_Data

Observation Support Tool

http://almaost.jb.man.ac.uk/

Splatalogue

http://www.splatalogue.net/

ALMA Publications

http://almaobservatory.org/en/science-at-jao/articles-a-publications

Press Releases

http://www.almaobservatory.org/en/press-room

ALMA Instructional Videos

https://science.nrao.edu/opportunities/courses/alma-instructional-videos

Cycle 6 Quick Facts

http://science.nrao.edu/facilities/alma/didyouknow

ALMA Web Site for Children

http://kids.alma.cl/?lang=en!37

!

The Atacama Large Millimeter/submillimeter Array (ALMA), aninternational astronomy facility, is a partnership of the European
Organization for Astronomical Research in the Southern Hemisphere (ESO), the U.S. National Science Foundation (NSF) and the
National Institutes of Natural Sciences (NINS) of Japan in cooperation with the Republic of Chile. ALMA is funded by ESO on behalfof its Member States, by NSF in cooperation with the National Research Council of Canada (NRC) and the National Science Councilof Taiwan (NSC) and by NINS in cooperation with the Academia
Sinica (AS) in Taiwan and the Korea Astronomy and Space Science
Institute (KASI).
ALMA construction and operations are led by ESO on behalf of its
Member States; by the National Radio Astronomy Observatory(NRAO), managed by Associated Universities, Inc. (AUI), on behalfof North America; and by the National Astronomical Observatoryof Japan (NAOJ) on behalf of East Asia. The Joint ALMA Observatory (JAO) provides the unified leadership and management of theconstruction, commissioning and operation of ALMA.

!38