Doc 8.1, ver. 2 | 8 Mar, 2021

Observing with ALMA – A Primer(Cycle 8 2021)

ALMA, an international astronomy facility, is a partnership of ESO (representing its member states), NSF (USA) and NINS (Japan),
together with NRC (Canada), MOST 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
2

5 March 2020
8 March 2021

G. Schieven
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., 2021, Observing with ALMA – A Primer, ALMA Doc. 8.1, ver. 22

Table of Contents
Purpose of this Document
Some Acronyms Used in this Document
What is ALMA?
ALMA Regional Centers (ARCs)
What is Interferometry?
Science with ALMA

2
2
3
4
5
6

ALMA Full Operations Specifications

6

ALMA Cycle 8 2021 Capabilities

7

Receivers and Sensitivities This Cycle
ALMA Correlators
Did You Know? In Cycle 8 2021 ALMA can…

8
8
9

Science During Cycle 8 2021

10

Before You Propose

10

Examples of Cycle 8 2021 Observing With ALMA
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
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

Interferometry Concepts for ALMA: A Glossary of Terms
A Few Useful Equations
A Summary of “Quick Links”

1

22
23
27
28
29

31
36
37

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

EHT

Event Horizon Telescope

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

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. Decades-long monitoring studies of the sky above
this site 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 ~55th and 25th
percentile conditions respectively averaged over the year. Thismeans that ~55% 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 1/4 (25%) 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 this cycle.

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 European ARC (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, 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 https://researchers.alma-telescope.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 isequivalent to measuring the Fouriertransform of the sky. In a single integration (typically a few seconds orless), each pair of antennas, called abaseline, samples a single point inthis power spectrum, at a position in
Fourier space related to the distance
between the pair of antennas andthe position angle of the baseline
vector. Antennas which are close
together (short baselines) sample
large-scale angular structure, while
long baselines sample very smallscale angular structure. By combining these data, called visibilities, overFigure 6: An areal view of ALMA in a relatively compact configuration. Thea large number of baselines (the uvmaximum distance between the 12-m Array antennas in the most compact concoverage), the Fourier plane is samfiguration is about 150 m, and the minimum distance a mere 15 m. Indicatedpled, which can then be inverted(using a Fourier transform) to realso are the AOS Technical Building, and the ACA 7-m Array. (Credit: C. Padilconstruct an image.
The reconla/ALMA (ESO/NAOJ/NRAO))
structed image quality is very sensitive to the uv-coverage — how completely 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 this cycle provides good coverage. During longer observations, more of the uv-plane isfilled in by the rotation and foreQuick Links
shortening of baselines as the Earthrotates on its axis. Antennas are
For a delightful 13 minute introduction to radio interferometry,
periodically moved (reconfigured)
see Dale E. Gary’s 1991 video on Radio Synthesis Imaging at:
to provide a wide range of array
https://vimeo.com/425868315
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 is revolutionizing 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 - 16.2 km~0.2” × (300/ν GHz) × ( 1 km / maximum baseline )~19.4” × (300/ν GHz)
~33.3” × (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 35 GHz when Bands 1 and 2 are 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,
and galactic structure. The resolution of ALMA reveals the kinematics of obscured active galactic nuclei and6

Figure 8: ALMA image of the planet-forming disk around the young star RU Lup.
The inset image (lower left, red disk) shows a previous (DSHARP) observation of thedust disk with rings and gaps that hint at the presence of forming planets. The newobservation shows a large spiral structure (in blue), made out of gas, that spans far
beyond the compact dust disk. Credit: ALMA (ESO/NAOJ/NRAO), J. Huang and S.
Andrews; NRAO/AUI/NSF, S. Dagnello; 2020, ApJ, 869, L42)

quasars 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 can enable the accretion of cloud material onto an accretion disk to be imaged and can trace the formation andevolution of disks and jets in young protostellar systems. For older protostars as well pre-main sequence stars, ALMA
can 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 has the large total bandwidth, high spectral
resolution, and sensitivity needed to detect the myriad lines associated with the heavy, pre-biotic molecules that
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 can 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. Many recent exciting results are shown in figures throughout this document.

ALMA Cycle 8 2021 Capabilities
Capabilities offered for this cycle 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 (3.6mm to 300 µm), using receiver Bands 3 through 10 (see Figure 4)• Baselines up to 8.5 km for Bands 3 - 7 (configurations C-9 and C-10 will not be offered in Cycle 8 2021); andup to 3.6 km for Bands 8, 9 & 10
• Stand-alone ACA observations may be requested in all receiver bands (new!)• Single field interferometry at all receiver bands, and mosaics at Bands 3 through 9• Full polarization capability, including circular polarization, at Bands 3 through 7 in both continuum andspectral line mode using the 12-m Array, plus single-field full polarization with the 7-m Array (new!)• Mosaicking of continuum linear polarization observations at Bands 3 through 7 using the 12-m Array (new!)• Spectral scans using the 12-m Array or the 7-m Array (new!)• Solar observations at Bands 3, 5 (new!), 6 and 7• Millimeter-wavelength VLBI of faint science targets (new!) at Bands 3 and 6• ALMA-only phased array mode at Band 3 useful primarily for pulsar observations (new!)• See the Proposer’s Guide on the Science Portal for a full list of capabilities this cycle

Quick Links
Starting in Cycle 8 2021, ALMA proposals will need to be written using a dual anonymous procedure, andmost proposals will be refereed using a distributed peer review process. To find out how this affects you,see:
https://almascience.org/proposing/alma-proposal-review
7Receivers and Sensitivities This Cycle
Table 1 shows the receiver bands which are available for Cycle 8 2021, and various properties. The sensitivities assume 43 antennas, 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 8 2021 Receiver Bands

Most Compact

Most Extended‡Band

Frequency(GHz)

Wavelength(mm)

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

AngularResolution(“)

Approx.
Maximum
Scale (“)
(see P.24)

Spectral
Sens.
∆Tline(K)

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

Spectral
Sens.
∆Tline(K)

384-116

3.6-2.6

69-50

0.08

4.0-2.9

34-25

0.08

114-83

1.69-1.23

96

4

125-163

2.4-1.8

46-36

0.09

2.7-2.1

23-18

0.08

77-59

1.14-0.87

1045158-211

1.9-1.4

37-28

0.12

2.1-1.6

18-13.5

0.11

61-46

0.90-0.67

137

6

211-275

1.4-1.1

28-21

0.12

1.6-1.2

14-10

0.10

45-35

0.67-0.52137

7275-373

1.1-0.8

21-16

0.22

1.23-0.91

10.4-7.6

0.2

35-26

0.52-0.38254

8385-500

0.78-0.6

15-12

0.40

0.88-0.68

7.4-5.7

0.37

55-42

0.67-0.52969602-720

0.5-0.42

9.7-8.1

1.4

0.56-0.47

4.7-4.0

1.8

35-29

0.43-0.36475

10787-950

0.38-0.32

7.4-6.1

3.2

0.43-0.36

3.6-3.0

4.3

27-22

0.33-0.271150‡This cycle, the maximum baselines extend only to 8.5 km.

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
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 this cycle.
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 8 2021 ALMA can…
…resolve molecular structures in M83:
clouds of excited CO(J=3-2) gas across the central 400 pc of M83 at 6pc resolution

in 1.7 hours…detect the ISM in high-redshift galaxies:
dust emission in a “normal" 1011 L⊙ galaxy between z=1 and z=6

in 4.3 hours
HC…reveal the behavior of solar system objects:
obtain wind patterns in the atmosphere of Mars with 300 km resolution

in 26 minutes

measure Kuiper Belt Object sizes from their thermal emissionin 48 minutes…survey Galactic Clouds and star forming regions:
measure the polarization of dust in 30 protostars in a single star-forming regionin 3 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 hours…measure stellar activity from low-mass to high-mass stars:
investigate heating mechanisms of red giant stars

in 5 minutes…study black holes and their environments, near and far:measure black hole mass of NGC 4526 from molecular gas kinematics

in 42 minutes…trace the formation of galaxy clusters, cosmic structure:
survey clustering in a sample of 23 Lyman-α Blobs (LABs) at z=3.1

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 hour

Science During Cycle 8 2021
From the start of science observing in late-September 2011 (Cycle 0),
Quick Links
ALMA was already a powerful millimeter/submillimeter interferometer, and is now essentially in full operation mode. By early 2020, overFind a list of published ALMA papers at
1700 papers featuring ALMA data had already appeared in refereed
http://telbib.eso.org
journals. Some exciting recent results are shown in figures throughout
this document. A few example projects suitable for this cycle havebeen compiled by the world-wide ALMA science staff and can be found on the previous page, or worked out inmore 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 to create, requiring careful positioning of the spectral windows in more than one
Science Goal, which strongly affects the time required.
• 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 this cycle 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 primary beam
decreases with distance from the field center. It isstrongly recommended that if the map area of interest (i.e. the LAS) is larger than about 1/3 theprimary beam, that you consider making a small(3 point) mosaic instead of a single pointing.
• Required sensitivity. Care must be taken when
Figure 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: This image (left) of the event horizon from the supermassive black hole at the center of the M87 galaxy was created by a consortium of observatories across the globe, and won the Event Horizon Telescope (EHT) team the $3M 2020 Breakthrough Prize in Fundamental Physics.
ALMA played a critical role in the EHT, providing both greatly enhanced sensitivity and critical baselines; theright panel shows how the image would have appeared without the Chilean observatories (ALMA + APEX). Astronomers can apply to
use the EHT for 1 mm Very Long Baseline Interferometry using the OT. (Credits: EHT Collaboration)

estimating the sensitivity needed per synthesized beam, particularly if estimating the source brightness fromsingle-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 to estimate required ALMA sensitivities.
• Dynamic range needed. The dynamic range (ratio of brightest to the faintest emission) in an image (or ratio ofstrongest to weakest detectable features in a spectral line image(spectral dynamic range)) improves with the length of the observation and the number of baselines (which goes nearly with the
Quick Links
square of the number of antennas). Observations of fields where
Resources for proposers, including the Call forthere is a bright nearby source but where the faint emission is of
Proposals, the OT, Technical Handbook, andinterest (e.g. faint sources near OMC1) will require great care,other links can be found on the Science Portal atand may not achieve the theoretical noise level. (See p. 26.)
http://almascience.org/proposing/With at least 43 antennas, ALMA will have 903 or more baselines. Thus the array should be able to image most “simple”
fields with good fidelity. Imaging dynamic ranges are typically
expected to reach up to around 50-100 depending on band and configuration (extended configurations or
Bands 8, 9 and 10 are likely to be closer 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 acomplex field or one with a nearby bright source, then the uv-coverage can be as important to consider as sensitivity. If you have a reasonable model for your source 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 oftime, so that surveys of objects scattered across the sky are much less efficient than surveys of objects nearby inthe 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 8 2021 Observing With ALMA
In the following sections we provide a few examples of observations that could be done with ALMA this semester.
Note that these examples use the sensitivities and capabilities of ALMA as they were known at the time of the Cycle 8 2021 Pre-Announcement. Anyone 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 experimentbecause 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”, corresponding to array configurations C-4/5/6.
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 is very simple. In the OT in the Spectral Setup frame, we request a start and end
Figure 11: Examples of molecusky frequency of 86.0 and 115 GHz respectively, which is nearly the entire frequency
lar absorption lines detected torange of the Band 3 receiver (84-116 GHz). The OT automatically sets up a set of fiveward PKS1830-211. These obtunings to cover the requested interval. To reach our target line sensitivity, a total
servations were made by ALMAin late 2010 when there were
integration time of 15 minutes (plus overheads of approximately 48 minutes) will be
only a few antennas.
required to cover 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 gravitationally
the Plateau de Bure interferometer. Both figures are from Swinbank etlensed into two “images” (Figure 12) that have a
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: A resolution of 0.1" (spectral line [CO(9-8)] and continuum, Band 7) is sufficientto spatially resolve the components revealed by the Harvard-Smithsonian Submillimeter Array (SMA), andperhaps resolve individual lenses. Though the total extent of emission is ~5”, only structures of order 1” areof 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 42 minutes of integration to reach 1 σ = 0.25 mJy/beam plus
another 36 minutes for calibration/overheads. This amount of time will also produce a sensitivity of 34 µJy/
beam in a map of the remaining 5.625 GHz of line-free continuum providing a S/N ~29 for the continuum signal (~1 mJy).
Figure 13: (left) ALMA Band 6 continuum
image of the Cosmic Eyelash 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 theexistence 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 optical ornear-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 andmore compact 12-m Array configurations are not needed forthese observations since the sources are small. An LAS of 1” isselected.
Spectral Resolution: These are purely continuum detections,
so only the TDM mode (15.6 MHz channels) is required.
Continuum (Band 7) Sensitivity: In This cycle, up to 150
sources (see Proposers Guide for restrictions) may be observed in each science goal. As an example, there are ~150
sources in the COSMOS-AzTEC catalogue (Aretxaga et al. 2011)
Figure 14: The inset of this HST image galaxy
with (de-boosted) flux densities > 3.3 mJy. Furthermore thecluster Abell 2744 shows an ALMA image of thesources are all within a few degrees of each other, making thez=8.38 gravitationally-lensed galaxy A2744_YD4.
observations highly efficient since the calibration can be shared
At this distance, the universe was a mere 600 milamong many sources. Given the exceptional atmospheric conlion years old, and yet the ALMA observations
ditions at ALMA, we choose to pinpoint these sources at ashow that the galaxy is rich in dust. (Credit:
higher frequency (344 GHz or 0.8 mm) where they are signifiALMA (ESO/NAOJ/NRAO); NASA; ESA; ESO;
cantly brighter (S α νβ, typically β~2 at z~2). Therefore, thesesources should have S0.8mm > 5 mJy. It is possible that these
D. Coe (STScI)/J. Merten (Heidelberg/Bologna);sources may be extended or may resolve into more than one
Laporte et al., 2017, ApJ, 837, L21)
source, so we aim to get a S/N ~ 40, adequate to identify thec o u n t e rparts and obtain excellent relative astrometric accuracy, which
is usually estimated 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 onsource integration, for a total of 12.7 hours including overheads
and calibration.
Figure 15: ALMA observed the earliest known example of a merging(non gravitationally lensed) galaxy. This image of B14-65666, a Lyman break galaxy at z=7.15, shows the distribution of dust (in red),oxygen (in green), and carbon (in blue), observed by ALMA, plus stars(in white) from the HST. (Credits: ALMA (ESO/NAOJ/NRAO),
NASA/ESA Hubble Space Telescope, Hashimoto et al., 2019, PASJ, 71,
71)
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 8 2021, there is no limit on the size of the time window,
but whether 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: GRB161219B was observed by ALMA at Band 3 over five epochs ranging from 1.3 days to 78 days after the burst.
On the left above, the Band 3 flux density is plotted as a function of time. The rebound, or reverse shock, triggered by the GRB’s
powerful jets slamming into surrounding debris, lasted thousands of times longer than expected. On the right is an artist’s impression of the “reverse shock” echoing back through the jets of the GRB. (Credits: Laskar et al., 2018, ApJ, 862, 94L; NRAO/
AUI/NSF, S. Dagnello)
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 starformation 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 molecular emission 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 to large angular scale (via the ACA)
make a wide-field, high-resolution image of a galaxy like M100 a viable project. Such data would probe the dynamical effects of spiral arms and the nuclear bar on molecular gas, map the density distribution of the molecular
ISM, and allow comparisons of molecular material to the distributions of dust, recent star formation, stellar populations, and other phases of the ISM.
Receiver(s): Band 3 (115 GHz): We consider a mosaic to observe CO (1-0) emission (~115 GHz, Band 3) acrossmost of M100.
Spectral Sensitivity: Large molecular clouds like the Orion complex have masses ~5 × 105 M⨀ or more. We design our survey of M100 to detect such clouds. For a CO-to-H2 conversion factor of ~2 × 1020 cm-2 (K km/s)-1 andassuming 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 10 km/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.
Largest Angular Scale: 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 8 2021 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 ~2.1 minutes per mosaic position. With 137 fields we expect that the program willneed about 4.8 hours on source, or justover 7.6 hours with overheads. The
OT estimates that 52 overlapping
pointings will be needed for the 7-m
Array requiring about 40 hours total,
and about 78 hours for the total power observations.
Figure 17: Six of the galaxies imaged usingmosaics of 12-m, 7-m, and TP Arrays in the2017 Large project to image 100,000 giant
molecular clouds in 74 nearby galaxies.
These results have been presented in a seriesof papers in the Astrophysical Journal andAstrophysicalJournalLetters.(Credit:
ALMA (ESO/NAOJ/NRAO); NRAO/AUI/
NSF, B. Saxton; J. Sun et al., 2018, ApJ,
860, 172)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, andeventually, perhaps, form planets within the disk. The evolutionary state of the disk
may be traced by determining the continuum SED of a coherent sample of protostellar
disks. Wide frequency coverage is necessary to measure the SED of the disk. The SED
is determined by a combination of the dust temperature and density distribution, andthe optical properties of the dust grains. Although this cycle ALMA will be able to resolve nearby protostellar disks in exquisite detail, for this survey we will only measureglobal properties but at great sensitivity in relatively short integrations. Here we discuss a potential project to observe a set of six Class II protostellar disks in Ophiuchus,selected from the catalogue presented by Evans et al. (2009, ApJS, 181, 321).
Receiver(s): Bands 3, 5, 7, & 9: (98, 203, 344, & 679 GHz respectively)
Figure 18: ALMA Band 7
Angular Resolution & LAS: The typical size of a protostellar disk is ~ 100 AU. At theimage of the dust in PDS 70,
distance to the nearby Ophiuchus molecular cloud (125 pc), 100 AU subtends 0.8”. Ina star system located approxthis experiment we aim to measure global properties of the disks, including the SED. We
imately 370 light-years fromtherefore seek observations at each frequency at 0.4”. Using the OT, we simply request a
Earth. Two faint smudges in
the gap region of this disk areresolution of 0.4” for each frequency band. The observations would then be executed inassociated with newly
more compact configurations for the higher frequencies, while Band 3 data would be obformed planets. One suchserved in a more extended configuration. If necessary, we can fine-tune the angular resoconcentration of dust is alution achieved in each band by applying a uv taper during the data analysis. It should
circumplanetary disk, thebe noted however that such a taper suppresses the contribution from the longest basefirst such feature ever detectlines and changes the sensitivity. For this experiment, a configuration which gives an aned around a distant star.
gular resolution of 0.4” (see Technical Handbook Chapter 7) will have recoverable angu(Credits: ALMA (ESO/
lar scales of about 3”, so most likely all emission will be recovered.
NAOJ/NRAO); A. Isella etSpectral Resolution: These are continuum observations, so this is not relevant.
al., 2019, ApJ, 879, L25)
Sensitivity and Observing Time: We calculate the expected brightness of a typicaldisk 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 consideration thatthe 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 are ~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 the seven bands respectively. Using the ALMA
sensitivity calculator with 43 main array antennas (12 m diameter)and 7.5 GHz of bandwidth, we find that these sensitivities can beachieved in 21 minutes per disk (Band 3), 2.4 minutes per disk
(Band 5), 40 seconds per disk (Band 7), and 2.6 minutes (Band 9)
for each science target, or 3.4h/30m/21m/39m respectively including calibrations and overheads.
Figure 19: The VLA/ALMA Nascent Disk and Multiplicity (VANDAM)
survey imaged more than 300 protoplanetary disks in the Orion molecular
cloud (OMC). This image shows the OMC in blue (from Herschel), the observed protostars in the survey (yellow dots), and inset are a few of the disks
imaged by ALMA (blue) and the Jansky VLA (orange). (Credits: ALMA
(ESO/NAOJ/NRAO), J. Tobin; NRAO/AUI/NSF, S. Dagnello; Herschel/
ESA; Tobin et al., 2020, ApJ, 890, 130)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 20: 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 will be ~6
mJy/beam. We request a sensitivity of 100 µJy/beam toachieve a 60σ detection of the polarization over the entire
core.
Observing Time: It will take 5.0 minutes on-source with a
7.5 GHz bandwidth to achieve the sensitivity. However, theOT imposes a minimum time of three hours for full-polarization observations in order to get sufficient parallactic angle coverage for calibration.

Figure 21: ALMA was used to detect the magnetic field in the jetfrom a protostar, using polarized SiO line emission. The top image(right) shows an optical H2 image of the HH 211 jet (from Hirano etal. 2006). The blue and red features show the blue- and red-shifted
SiO J=8-7 emission. The inset image shows the inner jet, with therecently detected disk (Lee et al. 2018), the solar system scale for comparison, and yellow bars showing the SiO line polarization vectors.
The vector orientations are consistent with a model of a helical magnetic field threading the jet. (Credit: ALMA (ESO/NAOJ/NRAO);
Lee et al., 2018, NatCo, 9.4636)
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 observedin 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 been found to be made of thousands of small (Diameter < 1"), dense (n ~ 105 cm-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
Figure 22:
The asymptotic giant branch star
ALMA began science observing the highest angular resoluW43A was observed by ALMA at Band 6, to undertion millimeter molecular line observations had beam sizes
stand how bipolar porto-planetary nebulae are gengreater than 3" (see Huggins et al. 2002, ApJ, 573, L55).
erated. Here CO(J=2-1) emission clearly shows a
Receiver(s): Band 6 (230.538 GHz) For this project wehigh velocity jet (blue) ejected from the aged central
choose to observe the CO (2-1) line.
star, within a lower velocity (green), with the dusty
Angular Resolution & LAS: 0.3" Taking the Helix Nebula
envelope of the star entrained by the jets shown inresults as a starting point, the angular resolution desiredorange. (Credits: ALMA (ESO/NAOJ/NRAO);
should be below the fragmentation scale of ~1" (which we set
Tafoya et al., 2020, ApJ, 890, L14)
as the LAS). We attempt to gain a factor of one hundred inangular resolution over previous observations.
Mosaic Required: The Helix is quite large (diameter ~ 25’)and highly fragmented. However, the diameter of the primary beam at 1.3 mm is only about 25". It would takean enormous mosaic of pointings to map the entire Helix and thus in this proposal one pointing each toward theSE and NW portion of the nebula are chosen.
Spectral Resolution: The spectral resolution is chosen to match the expected line profiles (~0.8 km/s FWHM)
within the Helix Nebula. We choose the 117 MHz bandwidth spectral mode and spectral averaging factor 4 togive a resolution of 0.158 km/s.
Channel (Line) Sensitivity: 0.5K. The fragments observed in the Helix Nebula are quite faint, with peak < 5 K.
In this scenario, a moderate 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 3.6hours to reach 0.5 K (or 6.6h including calibration and overheads). The two separate pointings will require about 13 hours
of ALMA observing time.
Figure 23: (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 spiralshell pattern (indicated by a white box) results from a highly ellipticalbinary 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 asprobes of the thermal response of asteroid surfaces. ALMA can resolve thermal emission froman asteroid's surface to measure the temperature
distribution across the surface and estimate theregolith thickness and composition. Tracking
the asteroid over time will reveal its rotational
period and 3D shape. (Note this project is not feasible in Cycle 8 2021 because the most extended configurations will not be offered.)
Receiver(s): Bands 6: (233 GHz or 1.3 mm)
Angular Resolution & LAS:
Near-infrared
adaptive optics imaging combined with modelling puts the size of Juno at ~240 km. To obtain at least eight resolution elements across theasteroid, a resolution of ~29 km is required
which translates to 20 mas at the distance toJuno 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 MaxFigure 24: Asteroid 3 Juno was observed in October 2014 durimum Recoverable Scale of 0.21” with the mosting the commissioning period of the ALMA long baseline camextended configuration (not available this cycle),paign. Over the course of the observations, the achieved resoluand so all relevant size-scales will be recovered.
tion varied from 32 x 24 mas to 42 x 36 mas, differing mostly
Continuum Sensitivity: Juno has a single-dishdue to changes in phase stability as the observations spanned
flux density of 240 mJy at 250 GHz. At a resoluthe transition from night through dawn and into daytime. Thetion of 20 mas, there are ~64 synthesized ALMA
baselines used to achieve this resolution ranged from 26 m tobeams across the 240 km (i.e. 168 mas) wide as13 km. The observations spanned 10 epochs, covering 60% of
teroid. Thus the flux density per beam will be
the 7.2h rotation period. (Credit: ALMA Partnership et al.,
240 mJy / 64 beams = 3.8 mJy/beam. The sensi2015, ApJ, 808, L2)
tivity required to reach a signal-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 can be achieved in 11 minutes of on-source integration time (but with approximately 37 minutes of 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 not necessarily need to be observed contiguously, and indeed contiguous observations longer than 2-3 hours cannot be guaranteed this cycle (see the Proposer’s Guide Appendix A.10), and can be scheduled over several
days using the Time Constraint feature in the OT.
Figure 25: Spherical image of Jupiter showing the distribution of ammonia gas below Jupiter’s cloud deck. (Credit:
ALMA (ESO/NAOJ/NRAO), I. de Pater et al.; NRAO/AUI
NSF, S. Dagnello)
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.
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 8 2021 offers observations of the Sunthat use a hybrid array comprised of both the fixed-station ACA 7 m antennas and the 12-m Array in configurations corresponding to angular resolutions of approximately 3.4”, 1.8”, 1.2” and 0.92” in Band 3(corresponding to configurations C-1 to C-4), approximately 1.7”, 0.91” and 0.61” in Band 5 (configurations C-1 to C-3), approximately 1.5”, 0.8”, and 0.5” in Band 6 (configurations C-1 to C-3), and approximately 1” & 0.7” at Band 7 (configurations C-1 & C-2). 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 43 12-m and 10 7m antennas available in a hybrid array, the sensitivity is significantly better.

Figure 26: The left panel shows a
230 GHz (band 6) image of the Sun
on 7 December 2016. In order toemphasize structure on the disk, the230 GHz image color display
ranges from 5300 to 7400 K. Lowlevel contours are plotted at 300,
600, 1200, and 2400 K in order toshow features above the limb. The
right panel shows disk profiles
through the Poles and on a diameter through the active region in thesouthwest quadrant, but with theblue curve offset by 800 K in orderto show structure in both. (Credit:
ALMA (ESO/NAOJ/NRAO; S.
White et al., 2017, SoPh, 292, 88)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, or register as an ALMA user on the Science Portal.
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 to which users
will be directed for data retrieval and helpdesk support. Users from Chile or from non-ALMA member regions may selectany of 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. (Starting in Cycle 8 2021, proposals will behttps://almascience.org/proposing/observing-tool
reviewed using a dual anonymous procedure, which
Instructional videos on using the OT can be found at
requires that PIs write their proposals in a way that
http://almascience.org/proposing/observing-tool/video-tutorials
preserves anonymity. For more information, see
https://almascience.org/proposing/alma-proposalMore info on the distributed peer review process and thereview) Phase 1 consists of a detailed observing produal anonymous procedure can be found at
posal with a scientific and technical justification subhttps://almascience.org/proposing/alma-proposal-review
mitted to the Observatory through the OT. The OT
includes a sensitivity calculator and viewers for assisting with correlator setups and mapping parameters while preparing Science Goals (SGs). Users can use the ALMA
Helpdesk, available from the ALMA Science Portal, to get assistance from ARC staff at any stage of the preparation andsubmission 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 thesetools.)
Proposal Review Process: Starting in Cycle 8
2021, ALMA has adopted a distributed peer review
process for most proposals. In this system, for eachsubmitted proposal the PI (or one of the delegated coIs) will be responsible for reviewing up to 10 othersubmitted proposals. Proposals requesting between
25 to 50 hours on the 12-m Array and all Large Pro-

Quick Links

The spectral line catalogue tool Splatalogue is available athttps://splatalogue.online
Information on observing with ALMA, including the Project
Tracker (SnooPI) and allocated projects, can be found at:
http://almascience.org/observing/*For the non-expert, a list of Interferometry Concepts is included, starting on page 31.*22

Figure 27: Composite image of Uranus’s atmosphere and rings obtained using ALMA in December
2017. The image shows thermal emission from the rings of Uranus for the first time, enabling scientists
to determine their temperature is a frigid 77 K. Dark bands in Uranus’s atmosphere at these wavelengths are due to different optical depths of the bands, with the brightest regions (e.g. north polar region) indicating the greatest depth and highest temperature. (Credits: ALMA (ESO/NAOJ/NRAO);
E. Molter and I. de Pater; Molter et al., 2019, AJ, 158. 47)grams will still be reviewed by a single international committee that is divided into a
number of science-themed review panels (Scientific Review). Priority status isawarded based on the proposal’s scientific ranking and available time. The time
available for projects will depend on the PI’s institutional affiliation, with 33.75%,
33.75%, 22.5%, and 10% made available to projects associated with the North American, Europe, East Asian, and Chilean partners, respectively. (Large programs, i.e. thoserequesting more than 50 hours of 12-m Array time or 150 hours of ACA time (see Proposer’s Guide section 3.3) will havetheir time split among the partner regions of the co-PIs.) Time will also be available to projects from non-ALMA regions
through the Open Skies policy. Users should consult 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 distributed peer review process
and on the dual anonymous procedure at: https://almascience.org/proposing/alma-proposal-review)
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 2. During Phase 2, successful proposers will havethe opportunity to review their SGs and make minor changes 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 tothe array 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 throughthe 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 theSBs are permitted during Phase 2. Any significant change requires a major change request through the Helpdesk, andmay result in delays to implementing the SBs and/or may not be granted. See the Proposer’s Guide section 6.2 for details.
Archive & Data Delivery:
After each SB has been sucQuick Links
cessfully observed, the data
A fuller description of ALMA observing considerations can be found inwill be calibrated, reduced
and quality assessed by
the Technical Handbook
ALMA/ARC staff and dehttp://almascience.org/documents-and-tools/latest/alma-technical-handbook
posited into the ALMA Science Archive, available fromthe Science Portal, where they may be retrieved by observers. ALMA data have a one-year proprietary period from thedate when they are placed in the ALMA Science Archive and made available to the PI. Archived data products include theraw 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, field-of-view, spectral resolution,and sensitivity. These quantities must be specifically defined and justified for a given project in a proposal, and carefulchoices are required to ensure that the project's scientific aims can be met. These quantities are also used during Phase 2, toguide in planning the 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 of23

finer 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 “uvcoverage”, 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 meters and ν is the observing frequency in GHz.
MRS is a guideline for the largest angular structure on which some of the flux of a smooth structure can be reasonably recovered by the interferometer. This
rule-of-thumb applies to the size
Quick Links
scale of smoothly varying structures in both dimensions. Smooth
A video explaining Maximum Recoverable Scales can be viewed atstructures larger than the MRS will
https://science.nrao.edu/opportunities/courses/alma-instructional-videos
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 emission that has been "resolved out," additional observations are needed, including observations with morecompact 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 the use 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 a remarkable 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 cubeswith up to 8192 independent spectral channels. The width of these channels can be defined from 3.8 kHz to 31.25 MHz. (Inthis cycle, 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 gal24

Figure 28: This figure links Desired Angular Resolution to the Maximum Recoverable Scale (MRS) in Cycle 8 2021. 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 oversize-scales from ~1” to ~11” (cyan region). Structures larger than the “Maximum Recoverable Scale” (MRS; see p. 24) are “resolved
out” by the interferometer. Structure larger than the MRS may be recovered either by adding observations using a smaller 12-m Arrayconfiguration, 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 first need 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 beneeded to recover all angular scales. (The OT will automatically determine the required combination of arrays/configurations to achieve
the resolution/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. Note that angularresolutions finer than (0.21” x 100 GHz/ν) are not permitted for Bands 8-10 (Limit Bands 8-10 on figure above). The two largest configurations (C-9, C-10) are not offered in Cycle 8 2021.
25

axies), wide bandwidths and low spectral resolution channels 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 theexpense of the spectral resolution. Channel averaging can be set up during Phase 1 in the OT, which has the added benefit ofreducing 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. The OT will select a spectral 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 todetermine 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 (∆K) is proportional to the sensitivity in Jy (∆S) divided by the angular size ofthe beam when the source is resolved. For a given ∆S, the corresponding ∆K increases with decreasing beam size; it isharder to detect extended line emission at high angular resolution. The quantity ∆S itself varies as the inverse-square rootof 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 andhow 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. Imaging dynamic ranges are typically expected to reach up to around 50-100 depending on band and configuration (extended configurations, or Bands 8, 9 and 10 are likely to be closer to 50). In somecases 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 that26

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’sbeam. To estimate the flux density within ALMA’s usually much smaller beam, you’ll need an estimate of the angular size of thesource. If the source is small compared to the ALMA beam, for example, then all of the flux will fit into the beam. 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 nearby galaxy, and detected a 200 km/s wide CO J=2-1
line with an integrated intensity of 22 K km/s. IRAM has an antenna diameter of 30m and an aperture efficiency ηA~ 0.45
(from the IRAM website) at 230 GHz. Plugging these values into the equation above, we find the measured flux density is
191 Jy km/s per beam. Dividing by the line width, the average intensity of the line is about 0.9 Jy/beam. We want to observe
it with ALMA at an angular resolution of 1”. What flux do we expect in a 1” beam? As before, that depends on how large thesource is. In the worst case, the emission is spread evenly over the 11” IRAM beam, or over 121 ALMA beams. In this casethe expected flux density would be about 7 mJy/beam. A 10-sigma detection per spectral channel would require a sensitivityof 0.7 mJy/beam. Conversely, if the source were only 1” or smaller in angular size, the expected flux density would be the full
0.9 Jy/beam. A 10-sigma detection of the line emission per channel would require a sensitivity of 90 mJy/beam.
As you can see, some a priori knowledge of source size can be quite important when estimating the required sensitivity.
In some cases you simply will not know this before getting ALMA data. In absence of other data or a solid physical argumentthat leads you to expect a particularsize, the most conservative estimate,that the source fills the single-dishQuick Links
beam, may be the best. In any
A video about estimating ALMA sensitivities can be viewed atcase, be sure to explain in yourtechnical justification how you
https://science.nrao.edu/opportunities/courses/alma-instructional-videos
estimated the required sensitivity.
28

Creating Images From Your Data
Once the data are taken, ALMA data will be reduced by staff from the JAO and ARCs using the ALMA Data Reduction 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 ALMA Pipeline-reduced calibrated data cube), and the scripts used by the ALMA Pipeline.
Once the ALMA 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 conceptsof reducing interferometer data.
This process can be distilled
Quick Links
down to two stages, calibrationand imaging, and we discuss
The CASA homepage
these below in turn. Casaguides
with links to documentation, downloads and tools is available at
http://casa.nrao.edu/
are available which will guide
you step-by-step in reducing real
Casaguides
ALMA Science 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) before, during and after the target source
observations are made. The calibrator data are 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 sources of known flux density and angular extent. The brightness of 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 fromthese 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. Typically flux calibrators are also used as bandpass calibrators.
Phase (or gain) calibration corrects for phase errors due to differential atmospheric changes above each antenna ontimescales of seconds tominutes, and for ampliFigure 30: The upper blue portion of this graph shows the spectral lines ALMA detected in thetude fluctuations as well.
NGC6334I star-forming regionThe lower black portion shows the lines detected by the European
Atmospheric water vaSpace Agency’s Herschel Space Observatory. The ALMA observations detected more than ten times
por can cause very rapid
as many spectral lines. Note that the Herschel data have been inverted for comparison. Two molecularfluctuations
overlines are labeled for reference. (Credit: NRAO/AUI/NSF, B. McGuire et al., 2018, ApJ, 863, L35)
timescales of seconds.
Water vapor radiometers(WVRs) are installed on
all antennas on the 12-m
Array to monitor thesevariations, which areroutinely applied to thedata during ALMA Pipeline processing. Longer
timescale phase drifts aremonitored using periodic observations of a moderately bright, very compact source at relatively
29

small angular distance
Quick Links
from the target. The
Learn the principles of interferometry.
best sources for phasecalibration are unre- Lectures and tutorials from the 2014 Synthesis Imaging Workshop (NRAO) are found at:solved at the angular
https://science.nrao.edu/science/meetings/2014/14th-synthesis-imaging-workshop/
scales probed by thearray; since such objects are point sources, their data have intrinsically zero phase (no emission at any angular offsets), and any phasechanges recorded in the data are due only to changes in the system and/or atmosphere. The cadence at which thephase 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 at higher frequencies
the variations have larger magnitude. The ALMA antennas "fast-switch" between the target and gain calibrators fromevery 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 the ALMA Pipeline so that project teams will be able to see
preliminary results quickly. More complex projects may have their data calibrated by members of the JAO and the
ARCs using a combination of the ALMA Pipeline and by hand using CASA. There are many approaches to reducing
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.

Quick Links
Angular Resolution (or Synthesized Beam)
The effective angular resolving power (equivaFor a brief introduction to radio interferometry,lent to a point spread function) provided by
see Dale E. Gary’s 1991 video on Radio Synthesis Imaging at:
the ensemble of transformed visibilities given
https://vimeo.com/425868315
its range of spatial frequency coverage. Essentially this is roughly proportional to λ/Lmax,
where λ is the observed wavelength and Lmax is thesize 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 canalso be adjusted through the choice of an imaging weighting function (see Imaging page 29). A typical observation
will result in a synthesized beam with a primary feature that can be approximated by a Gaussian whose FWHM istypically 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 this cycle.
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 receiverand correlator) which is digitized at the antenna. Up to four 2 GHz wide basebands
are delivered to the ALMA correlators (seeFigure 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 it is possible to separate out the sig- Figure 31: A graphical view of basebands and sidebands. Basebands may be
nal from the unwanted image sideband. tuned to overlap if the user wishes, or may be located so as to maximize the(See Sideband)) The user-selected correlator total bandwidth (as shown). Each baseband may be further subdivided into
configuration determines how many base- as many as 8 spectral windows. Up to four spectral windows per baseband
bands are ultimately used, where they are are available.
placed in the available IF range, and whichcorrelation products are produced (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
31

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

Array will have 1225 baselines. In Cycle 8 2021, atleast 43 antennas provide 903 or more baselines in
the 12-m Array.
Clean Image or Cube A deconvolved image (orcube of images), where the emission in each has
been modelled in some manner so that distortions
induced by secondary features to the synthesized
beam are minimized. The optimal method of deconvolution depends on the science goals of theobservation.
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 for a given angular resolution. An extended configuration provides a very small synthesized beam though at the expense of larger scale structure, and vice versa (see Figure 28).
See the Proposer’s Guide for the configuration schedule for this cycle.
Correlator A powerful computer which cross-correlates the amplified, down-converted signals from each antenna pair toproduce 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 this cycle) 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 3-7, ≳50 for Bands 8, 9 &10 can be difficult to achieve because of 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 interFigure 33: An open cryostat fitted with receivers covering Bands 3 through 10.
The two “empty” slots will eventually be occupied by Bands 1 and 2 receivers,
giving ALMA complete spectral coverage from ~30 GHz to nearly 1 THz.
(Credit: ALMA (ESO/NAOJ/NRAO)
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)

est 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.
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 8 2021; see the Proposer’s
Guide in 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 thesource 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 a Gaussian of FHWM equal to ~1.13(λ/D), where λ isthe 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 bycareful design and construction. The primary beam sets the field ofview (FOV) for an observation with the array, unless a larger mosaic
Figure 35: ALMA observations of of the protoplanetary disk around HD169142 whichshows an outer region composed of thin rings and a double gap. The image also reveals abright inner ring with signs of being subject to dynamical perturbations, possibly by a“mini-Neptune”-size planet. Solar system size approximated by Pluto’s orbit is shownfor comparison. (Credits: N. Lira – ALMA (ESO/NAOJ/NRAO); S. Pérez – USACH/
UChile.)33

is made.
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 (e.g. a bright compact source).
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 fromboth sidebands simultaneously. By default for Bands 9 and 10, 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 this cycle. (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 and affecting the uv-coverage.
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 of channels together. Note that, because the correlator applies a Hanning filter tothe 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 during calibration and imaging. The OT default is to use an
averaging factor of 2.
Spectral Window A spectral window is a frequency subrange of a baseband. Each baseband
may be divided into one or more spectral windows by allocating a fraction of the correlator
resources to each window (see Table 2 page8). The properties of the spectral window depend on the fraction of the correlator resources
allocated to it. A single window will have the

Figure 36: Four spectral windows were set up to observe these fourlines simultaneously in one baseband using the OT. The windows areeach 234 MHz wide (with 0.43 km/s channels), and are centered at345.80, 345.45, 346.11 and 345.09 GHz. The spectral windows must
fit within the 2 GHz wide baseband.
34

Figure 37: New analysis of ALMA observations reveal that Proxima Centauri emitted a
powerful flare that would have created inhospitable conditions for planets in that system. Here is shown the brightness of Proxima Centauri as observed by ALMA over thetwo minutes of the event on March 24, 2017. The massive stellar flare is shown in red,with the smaller earlier flare in orange, and the enhanced emission surrounding the flarethat could mimic a disk in blue. At its peak, the flare increased Proxima Centauri’s
brightness by 1,000 times. (Credit: M. MacGregor et al., 2018, ApJL, 855, 2)

maximum bandwidth and minimum channel spacing, while multiple spectral windows provide reduced bandwidth per spectral window and/or larger channel spacings. Spectral windows can be placed anywhere within the~2 GHz wide baseband. Each baseband may contain up 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 by a 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 atlarge scales (low spatial frequencies) can be restored to images by combining array data with those from singledish telescopes and/or a more compact configuration of antennas. For example, data from 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(N-1), where N is the number of antennas.
Visibility An interferometric observation of a source made at a specific spatial 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 the brightness and position of the emission relative to theposition where the antennas are pointed. These amplitudes and phases need to be calibrated during observations byobserving 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: Chilean fox relaxing near the array. ALMA assumes responsibility for the habitat located around the observatory facilities, under archaeological advisement with San Pedrode Atacama and Chilean authorities. (Credit: Jaime Guardia /ALMA (ESO/NAOJ/NRAO))
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 ofFor 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

Science & Software Tools

http://almascience.org/tools/

Proposing Guidance quick links

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

FAQ on Distributed Peer Review

https://almascience.org/proposing/alma-proposal-review

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

Common ALMA acronyms

http://www.almaobservatory.org/en/siglas

North American ARC

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

European ARC

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

East Asian ARC

https://researchers.alma-telescope.jp/e/ea-arc/

Intro to Interferometry

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

Intro to Radio Interferometry Video

https://vimeo.com/425868315

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

https://splatalogue.online

Did You Know…?

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

ALMA Publications

http://telbib.eso.org

Press Releases

http://www.almaobservatory.org/en/catagory/press-releases

ALMA Instructional Videos

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