Mark Booth1, Pamela Klaassen1, Claudia Cicone2, Tony Mroczkowski3, Martin A.
Cordiner4, Luca Di Mascolo5,6,7,8, Doug Johnstone9,10, Eelco van Kampen3,
Minju M. Lee11,12, Daizhong Liu13,14, John Orlowski-Scherer15, Amélie
Saintonge16,17, Matthew W. L. Smith18, Alexander Thelen19, Sven Wedemeyer2,20,
Kazunori Akiyama21,22,23, Stefano Andreon24, Doris Arzoumanian25, Tom J. L. C.
Bakx26, Caroline Bot27, Geoffrey Bower28, Roman Brajša29, Chian-Chou Chen30,
Elisabete da Cunha31, David Eden32, Stefano Ettori33,34, Brandt Gaches26,
Evanthia Hatziminaoglou3,35,36, Patricia Luppe37, Benjamin Magnelli38,
Jonathan P. Marshall30, Francisco Miguel Montenegro-Montes39, Michael
Niemack40, Conor Nixon41, Imke de Pater42, Yvette Perrott43, Sandra I.
Raimundo44,45, Elena Redaelli46, Anita Richards47, Matus Rybak48,49,50,
Nikolina Šarčević51, Dmitry Semenov52, Silvia Spezzano46, Sundar Srinivasan53,
Thomas Stanke46, Paola Andreani3, Maria T. Beltrán54, Bryan J. Butler55,
Sebastiano Cantalupo56, Miguel Chavez Dagostino57, Ana Duarte-Cabral18, Bjorn
Emonts58, Leigh Fletcher59, Dale E. Gary60, Stanislav Gunar61, Alvaro Hacar62,
Bendix Hagedorn2, Tomek Kaminski63, Fiona Kirton, Katherine de Kleer64, Eduard
Kontar65, Yi-Jehng Kuan66,30, John Lightfoot, Enrique Lopez-Rodriguez67,
Andreas Lundgren68, Stefanie N. Milam4, Atul Mohan69,70, Raphael Moreno71,
Galina G. Motorina72,73, Arielle Moullet58, Kate Pattle16, Alberto
Pellizzoni74, Nicolas Peretto18, Joanna Ramasawmy, Claudio Ricci75,76, Andrew
J. Rigby77, Álvaro Sánchez-Monge78,79, Maryam Saberi20, Masumi Shimojo80,81,
Aurora Simionescu50, Mark Thompson77, Alessio Traficante82, Cristian
Vignali83,33, Stephen M. White84
See Appendix E for the author affiliations.

AtLAST Science Overview Report

June, 2024

###### Contents

  1. Executive Summary
    1. 1 Introduction
      1. 1.1 Sub-mm astronomy in the context of past, present and future (sub-)mm facilities
      2. 1.2 The AtLAST Design Study, and purpose and scope of this document
        1. 2 Science Themes
          1. 2.1 Distant Universe and Cosmology
            1. 2.1.1 Unbiased Surveys of the Distant Universe
            2. 2.1.2 The Warm and Hot Universe via the SZ Effect
            3. 2.1.3 The Hidden Circumgalactic Medium
          2. 2.2 Nearby Galaxies
          3. 2.3 Our Galaxy
          4. 2.4 Solar System
          5. 2.5 The Sun
          6. 2.6 Transients and Variability
          7. 2.7 Black Hole Event Horizons
            1. 3 Telescope Requirements and Generalised Instrumentation
              1. 3.1 Telescope requirements
                1. 3.1.1 Single dish
                2. 3.1.2 Dish diameter
                3. 3.1.3 Field of view
                4. 3.1.4 Observing at Sub-mm wavelengths
                5. 3.1.5 Scanning speed
                6. 3.1.6 Pointing and Tracking
                7. 3.1.7 Beam accuracy
              2. 3.2 Instrumentation requirements
                1. 3.2.1 Continuum Observations
                2. 3.2.2 Spectral Line Observations
                3. 3.2.3 Polarimetry
                4. 3.2.4 Very Long Baseline Interferometry
              3. 3.3 Telescope Operations
                1. 3.3.1 Observing modes
                2. 3.3.2 Data flow and legacy value
                  1. 4 Summary
                    1. A Frequency Ranges of ‘ALMA Bands’
                    2. B Confusion Limit
                    3. C AtLAST Sensitivity Calculator and Simulator
                      1. C.1 Underlying Equations, Assumptions and Caveats
                        1. C.1.1 Assumptions
                        2. C.1.2 Caveats
                      2. C.2 Implementation
                        1. C.2.1 Source code
                        2. C.2.2 Web Client
                      3. C.3 Imaging and Spectroscopic simulator
                      4. C.4 Related simulators
                    4. D Acronyms
                    5. E Author affiliations

## Executive Summary

Submillimeter and millimeter wavelengths provide a unique view of the
Universe, from the gas and dust that fills and surrounds galaxies to the
chromosphere of our own Sun. Current single-dish facilities have presented a
tantalising view of the brightest (sub-)mm sources, and interferometers have
provided the exquisite resolution necessary to analyse the details in small
fields, but there are still many open questions that cannot be answered with
current facilities. Using sub-mm wavelengths to observe the universe allows us
to unlock the motions of gas and dust throughout its lifecycle. Many of these
populations are faint and diffuse, and so missed by current observatories. In
order to make major advances in (sub-)mm astronomy, what is needed now is a
facility capable of rapidly mapping the sky spatially, spectrally, and
temporally, which can only be done by a high throughput, single-dish
observatory. This is driving the design of the Atacama Large Aperture
Submillimeter Telescope (AtLAST).

In the distant universe, the goal is to study galaxy populations as a whole,
using a number of techniques (e.g. SZ, line intensity mapping) to enhance our
understanding of the processes that shaped these galaxies and shapes their
evolution. With AtLAST we will finally be able to probe the size and
sensitivity regime required to find ‘normal’ galaxies at high-redshift. The
circumgalactic medium, at all redshifts, is showing itself to be a faint but
key ingredient in that evolution, acting as a bridge between the interstellar
and intergalactic media enabling the flow of fresh materials between them. At
lower redshifts, in the nearby universe, we are able to start resolving the
processes of star and planet formation embedded within the galaxies that
started forming much earlier in the evolution of the Universe. With a better
understanding of the magnetic field structures, the heating, cooling,
chemistry and dust composition of nearby galaxies, including the Magellanic
Clouds and the Milky Way (the best resolved structures to study), we can
quantify the balance of forces on smaller (molecular cloud) scales that then
drive the evolution on galaxy and cluster scales. Closer in, high cadence
observations of the Solar chromopshere will further our understanding of space
weather. We will be able to resolve the planets in our solar system, study
their weather patterns and, in some cases, even resolve their moons. Studies
of the chemistry of comets will allow us to probe the origins of Earth’s
water.

A single-dish facility with a 50 m diameter provides the sensitivity to trace
the gas and dust from protostellar to circumgalactic scales (both near and
far), and the resolution to spatially resolve weather patterns on the planets
and moons of our solar system. This level of sensitivity and resolution will
enable the systematic detection of the relativistic and kinetic SZ effects
and, for the first time, the emission from ‘normal’ galaxies at high redshift,
which is hidden below the confusion limits of current and planned single-dish
observatories.

The large field of view proposed for AtLAST ($\geq$1 deg2), will enable the
large format (multiplexed) instrumentation that will be able to detect
structures at these scales. The benefits are two fold here: detecting emission
on larger scales, and increased survey efficiency. The sensitivity to large
scale structure these instruments provide comes from the requirement to take
background measurements at much larger offsets instead of within the
instantaneous field of view of the instrument. The ability to fill the large
field of view with highly multiplexed instruments increases mapping efficiency
and serendipity. There are a number of large scale survey science cases
described here (i.e. the Galactic plane, the Magellanic clouds and nearby
galaxies) that will require 1000s of times fewer pointings if the AtLAST FoV
were filled when compared to current generation facilities. With this large
FoV, efficient transient/variability detection algorithms built into the data
processing algorithms will enable a wide variety of serendipitous detections.

Similarly, high bandwidth receivers (in both the continuum and spectral
domain) will enable rapid surveys by reaching higher sensitivities / detecting
more lines in a given observation than current generation facilities. This
exploit is similar to the reasoning for the ALMA WSU. As receiver technology
advances over the next decade, polarisation preservation is becoming in-built,
and not only will these highly sensitive receiver technologies study the
spatial distribution of gas and dust, but its magnetic properties as well.

The high sensitivity provided by AtLAST will result in deep observations on
short timescales, which enables cadenced observations to be performed on
various timescales: from a few seconds, to hours, to years depending on the
science need (i.e. solar flares, cloud motions in planetary atmospheres,
protostellar flares, AGNs and GRBs).

This study has uncovered a number of key science goals that cannot be tackled
with current generation facilities, even if their instrumentation were
upgraded. To satisfy the science needs of the community an entirely new
facility is required with high throughput (large diameter and field of view)
and the high sensitivity that comes with it. The lifecycle of gas and dust in
the Universe is proving to not be a simple, steady state process, but one full
of surprises. AtLAST will be able to unlock hidden populations of objects,
study how they interact with their environments, and how they vary over time.

One of the key deliverables for the science work package of the AtLAST Horizon
2020 design study is a robust set of science-driven requirements for the
telescope design. In this regard, we have solicited input from the astronomy
community, organised into science working groups (SWGs) covering different
themes. These SWGs have produced a collection of papers describing the key
science drivers in each area and the requirements they place on the
observatory111These have all been submitted to the journal Open Research
Europe: https://open-research-europe.ec.europa.eu/collections/atlast/about..
The purpose of this document is to summarise the results of these papers and
present the combined requirements set by them.

|  RoyalBlue!50 Where are all the baryons? |  RoyalBlue!50 How do structures interact with their environments? |  RoyalBlue!50 What does the time-varying (sub-)mm sky look like?
---|---|---|---
Mahogany!75 Detailed science goal |  Measuring the total gas and dust content of the Milky Way and other galaxies, in the interstellar, circumgalactic, and intergalactic media, reaching down to the sensitivities required to probe the typical populations of sub-mm sources. |  Understanding the lifecycle of gas and dust near and far; mapping the baryon cycle on multiple-scales; observing the interplay between gravity, radiation, turbulence, magnetic fields, and chemistry and their mutual feedback. |  Identifying the mechanisms responsible for time variability across astrophysical sources: from the Solar corona and other objects in our solar system to luminosity bursts in everything from protostars to active galactic nuclei.
Mahogany!75 Detailed technical specification |  High sensitivity to the faint signals (at sub-mK levels) on large scales ($\geq$1 deg2) from even the most diffuse and cold gas through sub-mm line tracers. Wide field ($>$500 deg2) continuum surveys capturing the plane of our galaxy and resolving 80% of the cosmic infrared background, probing typical populations and looking back over 90% of the age of the Universe. |  High spectral resolution and polarisation measurements on the relevant size scales for cores (0.1 pc, our galaxy), clumps (10 pc, nearby galaxies) and cloud complexes ($\sim$ few kpc, distant universe) to quantify the chemistry, disentangle the dynamics, and measure the magnetic fields working together to shape the evolution of structures within their larger-scale environments. |  An operations model that allows for highly cadenced and rapid response observations and data reduction pipelines with in-built transient detection algorithms; high time-resolution (few seconds) observations of our Sun and other stars.

Table 1: Key Science Drivers for AtLAST

From this process we find that there are three key themes where AtLAST will
provide transformational science, which are described in Table Executive
Summary. To accomplish this, AtLAST will require high sensitivity, high
angular resolution, continuum and spectral line instrumentation which
preserves polarisation, solar observing capabilities and a highly flexible
scheduling model capable of snapshot observations, large surveys, rapid
response to transient triggers and a transient pipeline capable of sending
transient triggers.

With a telescope design and operations model capable of fulfilling these
requirements and an innovative sustainability plan, AtLAST will be capable of
serving the astronomy community for decades to come.

##  1 Introduction

###  1.1 Sub-mm astronomy in the context of past, present and future (sub-)mm
facilities

Observations in the sub-mm and mm wavelength regime (by which we are roughly
referring to wavelengths between 0.3 and 10 mm) provide the opportunity to
view a wide range of astrophysical phenomena, from cold dust and gas to high
energy sources. Dust at 10 K has a peak wavelength of around 300 $\mu$m and
carbon monoxide (one of the most common molecules in the Universe) has many
emission lines at (sub-)mm wavelengths. These are all redshifted to longer
wavelengths the more distant the galaxy. They provide ways to measure the dust
and gas mass of galaxies, their circumgalactic media (CGM) and clouds, cores
and circumstellar discs within our own galaxy. Chemical signatures allow us to
study not just the chemistry of astrophysical sources but also their dynamics.
Hot gas can also be studied in the (sub-)mm due to the scattering of CMB
photons by free electrons in the gas, which can be detected in galaxy
clusters, protoclusters and even around individual galaxies. (Sub-)mm
observations are even possible of the brightest source in the sky – our own
Sun.

Figure 1: Comparison of Planck 353 GHz (Planck Collaboration et al. 2015;
left) and Herschel 1.2 THz (250 $\mu$m; Bresnahan et al. 2018; right) mapping
of the Corona Australis molecular cloud. While the Planck data shows the
overall shape of the magnetic field in the cloud, it cannot resolve individual
sites of star formation. With AtLAST, we will map magnetic fields across the
entire Galactic plane with a resolution at 345 GHz that is four times better
than that shown in the right-hand panel, thereby resolving individual star-
forming cores within molecular clouds. (Figure reproduced from Klaassen et
al., 2024)

AtLAST builds upon 60 years of sub-mm astronomy. In the 1970s and 80s the sub-
mm sky first opened up to astronomers through dedicated facilities such as the
Kuiper Airborne Observatory (KAO) and the James Clerk Maxwell Telescope
(JCMT), based on the ability to push far-infrared technologies to longer
wavelengths while simultaneously pushing radio and mm-wave technologies to
shorter ones. Their pioneering efforts led to the discovery of sources as
close as the rings of Uranus (Elliot et al., 1977), and as far away as
entirely new populations of galaxies - first dubbed ‘sub-mm galaxies’ because
they had never been detected at other wavelengths (Blain et al., 2002).
Combining the JCMT and the Caltech Sub-mm Observatory (CSO) into a sub-mm
interferometer (Carlstrom et al., 1994) then also paved the way for sub-mm
interferometers like the SMA, NOEMA and ALMA to be seen as technically
feasible.

Since those early days, and as our understanding of the sub-mm universe has
grown, the interest in observing it has blossomed - to the point of ALMA
routinely being the most oversubscribed ground based observatory in the world.
Hand in hand with this scientific curiosity has been the technological
development required to obtain those science goals.

At sub-mm wavelengths, we are sensitive to gas and dust across cosmic time. We
often observe this through the thermal emission of cold dust and low energy
transitions of molecules. But, at these wavelengths, we can also probe atomic,
forbidden and ionised gas lines, the warm/hot ionised medium (WHIM) of
galaxies and even the magnetic fields permeating these gas and dust
populations (see Figure 1 for an example). Some of the signals we detect emit
on small scales (of order arcsec and smaller) but some emit on very large
scales (few-10s of degrees) as shown in the rest of this document.

Figure 2: Angular resolution versus observing frequency for AtLAST and current
and planned telescopes with fields of view $\geq$1 deg2.

We are now a long way from the single pixel continuum and spectroscopic
cameras that first populated those early observatories, with single-dish
facilities like APEX, ASTE, the IRAM-30m, and upgrades to the CSO (now
decomissioned) and JCMT hosting continuum cameras with 100s to 1000s of pixels
to detect dust continuum and polarisation, 10s of pixel spectroscopic
instruments able to resolve gas kinematics with precisions below 0.01 km/s (R
$>$ 108), and starting to deploy low spectral resolution integral field units
(IFUs) with broad bandwidths aimed at determining the redshifts of all objects
in their fields of view. The above are general purpose observatories with
broad suites of instruments adaptable to the changing needs of their diverse
science communities. In parallel to this, and on the horizon are a number of
6-12 m class observatories dedicated to specific science goals/surveys.

Single dish observatories have shaped our understanding of the sub-mm sky, and
have found some of the most interesting objects for interferometric followup.
As we learn more and more about the sub-mm sky through single dish
observations there are some questions that can only be answered with higher
spatial resolution observations, which requires the use of interferometers.
However, interferometers trade high angular resolution for their ability to
detect large scale structures due to the lack of short baselines. This is a
trade-off worth making when the structures being observed do not emit on large
scales, information about large scale structure is not important, or that
information can be folded in latter using dedicated total power telescopes.
This latter technique requires combining data from observatories like the JCMT
and SMA, IRAM-30m and PdBI/NOEMA, or within a single observatory as is done in
ALMA: combining the 12 m array with its compact array and total power
antennas. Single dish observatories are powerful observatories in their own
right, but can also improve the results of interferometers by detecting the
large scale emission that interferometers are blind to, yet the current
observatories are limited in this regard as they lack the sensitivity and dish
size necessary to complement ALMA, particularly for faint, extended sources
(Plunkett et al., 2023).

Looking ahead to the 2030s and beyond, the landscape of long wavelength
astronomy is set to change again. With interferometric facilities like the
ALMA Wideband Spectroscopic Upgrade (WSU; Carpenter et al., 2023) and
SKAO222https://skao.int coming online at the beginning of that timeframe and
potentially the ngVLA333https://ngvla.nrao.edu along with it in combination
with dedicated large area survey telescopes like SO, CMB-S4 and FYST, we are
entering a new golden age for unlocking the secrets of the long wavelength
sky. In addition to the main science drivers of the latter set of single dish
facilities, they will provide finder charts for followup observations with
these next generation interferometers. But these surveys will be inherently
shallow (i.e. limited to 0.1 - 1 mJy over 100 - 600 GHz) because of their
limited collecting areas and high confusion limits dictated by their 6-10 m
individual apertures (see, e.g., Figure 13 of Blain et al., 2002). A 50 m
class single-dish facility will be able to probe deeper (with confusion limits
10x lower at 100 GHz, and 1000x lower at 600 GHz), at higher spatial
resolution (3x the linear resolution of JCMT, 4x that of APEX, and 8x that of
the upcoming 6 m facilities – see Figure 2) and much more quickly (single
point sensitivities 9x, 16x and 64x those of the JCMT, APEX and upcoming 6 m
class facilities per unit integration time) than other upcoming observatories.
This combination of capabilities means with a 50 m class sub-mm telescope that
has a large field of view, we can push to the ‘normal’ population of galaxies
at high redshift (well below the confusion limits of current and upcoming
facilities), see the large scale, but diffuse and faint, interstellar and
circumgalactic media of galaxies including our own, resolve variability across
the nearby universe and weather patterns on planets and moons in our solar
system, while ensuring that the next generation interferometers will have many
more interesting objects to observe over their lifetimes (i.e. they will not
become ‘source-starved’), while putting those objects into their environmental
context.

###  1.2 The AtLAST Design Study, and purpose and scope of this document

The AtLAST design study is a 3.5 year Horizon 2020 funded project with the aim
of deriving a next generation single-dish sub-mm telescope designed with
science and sustainability at the forefront (Klaassen et al., 2020). The
project spans six work packages including (WP1) overall management, (WP2)
telescope design, (WP3) site selection, (WP4) telescope operations, (WP5)
sustainability studies and (WP6) science. Since the beginning of the project,
a detailed optical and telescope design has been developed (Gallardo et al.,
2024; Mroczkowski et al., 2024; Kiselev et al., 2024; Puddu et al., 2024;
Reichert et al., 2024), two sites on the Atacama plateau have been monitored
and studied for suitability444https://www.atlast.uio.no/design-
study/wp3-site/d3.1_site_selection_criteria.pdf, and an operations plan taking
equality, diversity and inclusion principles into account has taken shape.
Research into powering the telescope using sustainable methods has been
conducted (Viole et al., 2023, 2024a, 2024b), including consultations with the
local San Pedro de Atacama and astronomical communities to ensure their views
on the future of sustainable astronomy are heard and understood (Valenzuela-
Venegas et al., 2023). Refining the science case is the work package discussed
in this report.

The purpose of this document is to summarise the work of the science community
over the Horizon 2020 AtLAST design study to derive and refine the science
drivers and understand what we could achieve with such a disruptive and
innovative leap in technology. In Section 2 we summarise the science drivers
for the telescope. In Section 3 we describe the constraints placed on the
telescope by the many and varied science drivers and summarise this in a
matrix of requirements.

##  2 Science Themes

The science consultation for the AtLAST design study began with polling the
worldwide community to submit individual science use cases. This resulted in
the participation of over 100 scientists from 22 countries submitting 28 use
cases (see figure 3 for their geographic distribution). These use cases saw
the community requesting both spectroscopy and imaging, and everything from
limited area mapping to large scale surveys. From these submissions, four key
science categories were identified: The Sun and Solar System, the Milky Way,
Nearby Galaxies, and the Distant Universe and Cosmology. Embedded in these
themes was also the need for cadenced (and target of opportunity) observations
to measure variability in the sub-mm sky. Together these use cases showed a
great desire from the community for a new single dish facility with higher
spatial resolution, greater mapping efficiency, orders of magnitude increases
in sensitivity, and highly multiplex instruments being able to observe
continuum and line emission, often requiring polarisation preservation.

Based on these themes, working groups were created to do more in-depth studies
into the feasibility of answering new questions in these subject areas. Where
gaps were identified in the initial submissions, under-represented science
communities were targeted for more participation to ensure the derived science
cases were as well rounded as possible. Alongside this process, a sensitivity
calculator was developed to allow the working groups to quantify their science
goals with observing time estimates. This calculator is described in more
detail in Appendix C. These working groups met periodically throughout the
design study and the results of their investigations are captured in a series
of white papers (van Kampen et al., 2024; Di Mascolo et al., 2024; Lee et al.,
2024; Liu et al., 2024; Klaassen et al., 2024; Cordiner et al., 2024;
Wedemeyer et al., 2024; Orlowski-Scherer et al., 2024). Below, we summarise
the work captured in these white papers along with some highlights of the
science presented in the original use case studies Ramasawmy et al. (2022) and
related works.

Figure 3: Map showing the contributions to the initial use case submission
part of the community consultation for defining the key science drivers for
AtLAST. Figure reproduced from Ramasawmy et al. (2022). Figure 4: Exploring
the parameter space of a 1000 deg2 survey using 1000 hours of AtLAST observing
time. Top panel: SFR vs. redshift survey limit. Second panel: $dn/dz$. Third
panel: fraction of (mock) sources detected when multiple bands are used.
Bottom panel: expected AtLAST accuracy in inferring $L_{FIR}$ and $M_{dust}$
as a function of redshift, demonstrating why multi-band detections are
important. (Figure reproduced from van Kampen et al., 2024)

###  2.1 Distant Universe and Cosmology

####  2.1.1 Unbiased Surveys of the Distant Universe

See van Kampen et al. (2024) for further details.

Single dish sub-mm telescopes like the JCMT were the first to identify an
entirely new family of galaxies in the distant universe: dusty star forming
galaxies (DSFGs, also known as sub-mm galaxies), which are too deeply embedded
in their own dust and gas to be detectable at shorter wavelengths (Blain et
al., 2002; Casey et al., 2014). They are the most actively star forming
galaxies in the early universe, responsible for the bulk of star formation
over cosmic time, and seem to be more abundant than expected from models. Much
of our understanding of them as a population has come from sub-mm observations
- whether through the pioneering sub-mm single dish and interferometric
observations from the ground (e.g. JCMT, CSO, SMA, ALMA), or from space (i.e.
Herschel, Planck). These surveys are currently biased towards the most massive
or gas-rich galaxies (10-100x more massive than the Milky Way; Hodge & da
Cunha, 2020). Because of the confusion limits of current facilities, we cannot
currently probe the underlying population of ‘normal’ galaxies, i.e. the
galaxies that constitute the bulk of the sources contributing to the cosmic
infrared background (CIB) in an unbiased manner. Understanding this population
is important for understanding the relationship between star formation and
black-hole growth (Carraro et al., 2020; Mountrichas & Shankar, 2023) and the
evolution of dust over cosmic time (Sommovigo et al., 2022; Dayal et al.,
2022; Drew & Casey, 2022; Hirashita & Il’in, 2022; Di Cesare et al., 2023).

In order to create blind, unbiased spectroscopic surveys of the ‘normal’
population of dusty galaxies at high redshift requires a telescope with a high
throughput, and the sensitivity and resolution that come with a large primary
mirror. With highly multiplexed instrumentation on AtLAST, this will, for the
first time, be possible. With the types of surveys described below, we will be
able to derive number counts, clustering estimates down to small scales and
redshift distributions of the normal population of dusty galaxies down to
$L^{\star}$, the luminosity of a ‘typical’ galaxy555More precisely, it is the
characteristic luminosity at which the luminosity function exhibits a rapid
change in slope (Schechter, 1976) and it happens to be very close to the
luminosity of the Milky Way.. Such a survey, functionally similar to the SDSS
(Geach et al., 2019), would complement upcoming optical and near-IR surveys
like those from the Rubin Observatory as it would allow us to measure
redshifts as well as the gas and dust contents of hundreds of thousands of
galaxies, many of which will have optical counterparts detected by the Rubin
Observatory.

As confusion limits are proportional to the resolution of the telescope,
AtLAST’s 50 m aperture means its confusion noise limits at sub-mm wavelengths
will be much better (i.e. lower noise) than any other current or planned
observatory. For instance, the photometric confusion noise at 350 $\mu$m (see
Appendix B) will be a factor of over 10,000 lower than that of 6 m telescopes
like ACT and FYST. With a multi-chroic camera, a 1000 deg2 continuum survey
could be conducted in 1000 hours that reaches a sensitivity limit of 570
$mu$Jy at 350 $\mu$m, thereby resolving 82% of the CIB. This will give the
most comprehensive view of the infrared luminosities and dust masses of
galaxies, measure star formation rates out to a redshift of $z$=5, and give
new insights into galaxies at a wide range of redshifts, masses,
metallicities, environments and morphologies (see Figure 4). With the
appropriate cadence, this will also enable transient science (see Section
2.6).

To accurately measure the redshifts of these galaxies requires detecting
spectral lines, which requires much deeper observations. Taking a ‘wedding
cake’ approach enables smaller area, deep spectroscopic studies within regions
of the larger continuum survey described above. By layering from shallower but
larger area surveys to deeper and smaller areas we build up a deeper
understanding of an unbiased sub-sample of the normal dusty galaxy population,
pushing out to redshifts as high as 7 in these deeper spectroscopic layers. In
a 1 deg2 area, it is expected that around 100,000 galaxies would be detected
in around 3000 hours. In addition to measuring their redshifts, this would
provide information on the gas content, cooling budget, star formation rate,
dust mass and dust temperature of these galaxies.

Methods like line intensity mapping (LIM, to probe the power spectra of
galaxies), Baryon Acoustic Oscillations (BAO, for measuring fluctuations in
the density of baryonic matter in the primordial universe Cole et al., 2005)
and related redshift space distortions (RSD Peacock et al., 2001) using low
resolution spectrometers can also be derived from these surveys to address the
tension between the Hubble constant measured at different cosmic epochs (Di
Valentino et al., 2021). With a spectroscopic survey of [CII] and CO covering
1000 deg2 it will be possible to measure the Hubble constant to within 0.7%
and the growth rate to a precision of 7.3%.

Galaxy evolution in proto-clusters is currently not well understood due to the
small and heterogeneous samples from current surveys. A distinct large survey
of distant galaxy clusters with AtLAST will produce a high-redshift
counterpart to local large surveys of rich clusters like the well-studied
Abell catalogue, to understand the impact of environment on the formation and
evolution of these distant cluster galaxies. To guarantee spectroscopic
redshifts from $z=4.5-10$ and a comprehensive study of the most prominent
lines emitted from the cold ISM, an instrument with spectral coverage from 180
to 345 GHz will be necessary. If this is extended up to 700 GHz then it will
be possible to map the peak of the star formation and black hole activity of
the Universe at $z=2$ (Madau & Dickinson, 2014).

####  2.1.2 The Warm and Hot Universe via the SZ Effect

See Di Mascolo et al. (2024) for further details.

Warm and hot ($\gtrsim 10^{5}$ K) ionised gas makes up a large part of the
matter budget across a range of scales, from intergalactic filaments, to the
intracluster medium (ICM) and down to the circumgalactic medium (CGM) of
individual galaxies. This hot gas provides a powerful way of understanding the
interconnected evolution of populations of galaxies and large-scale
cosmological structures, the matter assembly of the Universe and its thermal
history. Such hot gas has typically been the realm of x-ray observations, but
the Sunyaev-Zeldovich (SZ) effect provides a way for this gas to be studied in
the (sub-)mm. This is a redshift independent spectral distortion of the cosmic
microwave background (CMB) through inverse Compton scattering by high energy
electrons (Sunyaev & Zeldovich, 1970, 1972, 1980), with the most dominant
contributions coming from its thermal and kinetic components (referred to as
tSZ and kSZ, respectively).

The tSZ effect is the distortion of the CMB spectrum due to thermalized
reservoirs of high-energy electrons – as, e.g., in galaxy clusters – resulting
in a shift in intensity at higher frequencies and is sensitive to ICM
pressure. The kSZ effect is the Doppler shift in CMB photons caused by the
proper motion of hot electron halos (or part thereof) relative to the CMB rest
frame and is sensitive to their peculiar velocity. These effects distort
observed emission over wide frequency ranges, which requires the use of highly
sensitive multi-chroic cameras that operate across the (sub-)mm wavelength
range to be detected. The targets of interest – from individual galaxies and
galaxy clusters to large-scale cosmic web – are diffuse and extend across many
degrees, which requires large fields of view and highly sensitive
observations. A telescope like AtLAST, with its large aperture, large FoV and
instrumentation across the (sub-)mm wavelength range would be ideal for using
the SZ effects to probe fundamental physics like the thermodynamical state of
the ICM, the warm/hot component of the circumgalactic medium (CGM) and
warm/hot intergalactic media, when they form, how they evolve and what impact
they have on star formation and galaxy evolution (see e.g. Mroczkowski et al.,
2019), as described below.

Observations of the morphology and thermodynamics of the hot gas in galaxies,
groups, and clusters gives a historical record of the physics that shaped it
(i.e. AGN feedback, dynamics, mergers, accretion and cooling). The effects of
these physical processes can be seen as deviations from radially symmetric
thermodynamic structure (predicted by simple models of gravitational collapse)
in the pressure of the hot gas. The tSZ gives a direct proxy for the pressure
deviations and measuring the tSZ effect gives us a calorimetric view of the
thermal properties of the hot gas hosted within haloes of individual cluster
galaxies, the extended ICM, and out to the large-scale IGM outskirts. Higher
sensitivity and greater resolution with AtLAST is required to get to the point
of being able to detect the full extent of the pressure distributions in the
ICM of galaxy clusters. With that information, we will probe how the pressure
distributions in high-z galaxy clusters deviate from universal pressure models
(which feeds into understanding multi-scale processes like mergers) and how
the pressure profiles in the outskirts of clusters vary, which feeds our
understanding of how well-virialised these clusters are, and where their
accretion shocks are. Understanding these processes is key to our
understanding of large-scale structure formation in the Universe.

On smaller scales, a calorimetric view of the gas in a galaxy tells us about
feedback processes on the size scales of galaxy haloes to full clusters. For
instance, AGN-driven outflows contribute to heating on these scales, and it is
still not clear how and through what specific mechanisms supermassive black
holes contribute to evolution of their host galaxies. Current studies of
‘integrated’ tSZ effects are starting to probe the imprints of these feedback
mechanisms (Crichton et al., 2016; Hall et al., 2019; Yang et al., 2022), but
are hampered by the necessity of stacking the data, which dilutes their
usefulness.

Including the corrections for relativistic effects on the tSZ effect, we can
directly measure the temperature of ICM electrons. Adding this temperature
measurement to the pressure and density measurements derived above, offers the
ability to measure entropy in the ICM, which is modified by the evolutionary
effects / physical processes mentioned above. This significantly improves our
physical models of galaxy clusters and their use as cosmological probes.
Measuring the relativistic thermal SZ effect for individual targets is
generally beyond the capabilities of current facilities. Higher resolution,
highly sensitive observations across the (sub-)mm band are required to reach
the sensitivities necessary to detect this effect.

Using instead the kinematic SZ (kSZ) effect gives insights into the peculiar
motions of cosmic structures like galaxy clusters. However, the shape and the
relatively weaker signal in hot clusters can make it more challenging to
detect, as does the fact that the spectral distortion characteristic of the
kSZ effect is consistent with Doppler shifting of CMB photons, which makes it
hard to distinguish from CMB anisotropies. In order to break this degeneracy
in individual observations, sensitivity to a wider range of spatial scales
(300 $\lesssim\ell\lesssim$ 20000 in terms of spherical harmonics) than
possible with current or planned facilities (including dedicated CMB
facilities) is required.666Note, however, that pairwise stacking has been used
extensively to separate the kSZ from both CMB and tSZ signals in a statistical
sense. Further, such observations benefit from a 220 GHz channel, where the
kSZ signature is dominant. Once detectable, the kSZ effect can be used to
probe the amplitude and growth rates of cosmological density perturbations
which then translates to being able to distinguish $\Lambda$CDM from
alternative cosmologies (Kosowsky & Bhattacharya, 2009; Mueller et al., 2015;
Bianchini & Silvestri, 2016). Further to this, detailed spatial mapping can be
used to characterise turbulent motions and thus the important driving
dissipation scales relevant for the feedback processes.

Being able to study the warm-hot intergalactic medium (WHIM) in detail through
both the tSZ and kSZ effects will give access to a better understanding of the
‘missing baryon’ budget as the dominant fraction at $z<3$ is expected to be
beyond the virial radii of their host galaxies, where the WHIM is (i.e. in the
Cosmic Web). To date, the most accurate way of probing the (high temperature
end of the) WHIM is through stacking of SZ images and X-ray measurements. The
latter suffers from the inherent biases described above, while the former
suffers from averaging out details in the process of stacking.

####  2.1.3 The Hidden Circumgalactic Medium

See Lee et al. (2024) for further details.

Galaxies are surrounded by a reservoir of gas and dust extending beyond the
interstellar medium (ISM) out to the virial radius, commonly referred to as
the circumgalactic medium (CGM777We do not impose a strict definition of the
CGM. In our loose definition, the CGM includes the matter at the boundary
between ISM and CGM.). Understanding the CGM allows us to evaluate the
feedback and feeding mechanisms that impact the galaxy’s ability to sustain
star formation. These feedback mechanisms leave observable signatures in the
density, temperature, metallicity, and morphology of the CGM. The total mass
of the CGM reservoir and the contribution from different gas phases
(neutral/cold versus highly ionised/hot) that exist within it are largely
unconstrained, as are the physical processes that affect and shape the CGM.
Observing the CGM is challenging because of its primarily faint, diffuse, and
extended nature (the emission can extend to many hundreds of kiloparsecs –
corresponding to an angular size of degrees for galaxies within 10 Mpc of the
Milky Way), with the ability to resolve sub-kiloparsec scale clumps and a
dynamic range capable of distinguishing the CGM from the significantly
brighter ISM.

As suggested above, there are a number of key areas where probing the CGM can
better constrain models of galaxy formation and evolution. The first is the
nature of cosmic accretion, which digs into how galaxies can sustain high
levels of star formation across cosmic time. It is still unclear how fresh
material is accreted onto a galaxy to sustain this star formation; is it
through direct accretion of cosmic streams in the intergalactic medium (IGM),
or is it reprocessed through the CGM? The second is the missing baryon
problem, which focuses on reconciling the difference between the expected
baryon mass of galaxies and their measured values. Is this missing baryonic
mass, which is comparable to or can even excess the amount in the disc itself,
hidden in the cold CGM? The third is the impact of feedback on galaxy
evolution. We know that feedback from star formation and AGN must play a role
in galaxy evolution, but what is their cumulative effect over time and size
scales? The answers to all these questions are likely hidden in the CGM.

Studies of the CGM in distant (z$\geq$1) galaxies have discovered large
reservoirs of cold gas surrounding the discs of galaxies, especially around
quasars and dusty star-forming galaxies (DSFGs, see Section 2.1.1).
Protoclusters at higher redshift (z$\sim$4) show starburst events coordinated
across galaxies separated by hundreds of kiloparsecs, hinting at the large
scale gas supply on to the protocluster region to sustain the massive
starburst configuration.

In clusters, Brightest Cluster Galaxies (BCGs) and interacting galaxies at
intermediate to high redshift seem to show these types of gas reservoirs (e.g.
McDonald et al., 2014; Dunne et al., 2021), sometimes extending to tens of
kpcs, which is at odds with the relatively gas-poor galaxies found in dense
clusters in the local universe (e.g. Haynes et al., 1984; Zabel et al., 2019).
In the local universe, galaxy interactions, tidal tails and streamers are
indicated as extended gas reservoirs and they could supply gas to feed star
formation and AGNs. Further, the feedback from star formation and AGNs through
galactic outflows and heating changes the chemical and dynamical properties of
the CGM (over time). However, observing the CGM is tricky because of its
highly extended nature and inherently low surface brightness.

Observations with current facilities, such as JCMT, APEX, ALMA and ATCA have
shown hints of extended atomic and molecular gas emission around galaxies
across cosmic time. However, their low surface brightness and their large
emitting areas (especially for nearby galaxies) make detailed observations
challenging, if not impossible with current facilities (see Figure 5). For
interferometers, the large scales involved in CGM emission prove difficult to
detect and current generation single dish facilities do not have the
sensitivity required to probe this emission. These two aspects combine to show
the need for a new facility, a 50m class single dish facility with a large
field of view like AtLAST, with the sensitivity to detect low surface
brightness emission on size scales of up to a degree. Current facilities are
able to pick up some of the brightest components of the CGM, such as the tidal
streams and fountains mentioned above, but cannot yet detect the underlying
gas mass reservoir, the extent and detailed chemical composition.

Figure 5: Mock observations of CO(3-2) emission from the ISM and CGM of a
star-forming galaxy at $z=0.02$ with an extent of $\sim$80 kpc (200”). The
input simulation is shown in panel (a). Panels (b), (c) and (d) show
observations assuming 10 hours of on-source time for a 50 m single-dish using
a single pixel ALMA detector, ACA and ALMA respectively. (Figure reproduced
from Cicone et al., 2019)

High-frequency forbidden lines (i.e. [CI], [CII] and [OIII]) are good tracers
of the CGM of distant galaxies (Schimek et al., 2024a, b) and [CI] is still
detectable in the sub-mm even for nearby galaxies, although requires a very
high sensitivity that is only possible with a large aperture. This motivates
the need for the ability to detect this line at low redshifts, which in turn
sets a requirement for the telescope to be able to observe in the highest
frequency sub-mm atmospheric windows. In addition to these spectroscopic
observations, observing the continuum emission from dust across the (sub-)mm
wavelength regime probes the underlying dust content of the CGM as well.

###  2.2 Nearby Galaxies

See Liu et al. (2024) for further details.

Studying the physics and chemistry of nearby galaxies allows us to understand
the interplay between the many processes that govern galaxy evolution across a
much wider range of physical environments and extremes than in our own Galaxy,
but at resolutions as close to Galactic as possible. Star formation theories
based on Galactic observations cannot convey the complexity or diversity of
environments in which other galaxies form and evolve (e.g. with different
metallicities, total masses, positions within and size of galaxy cluster,
etc.).

The closest companions to the Milky Way, the Magellanic Clouds (MCs), offer a
unique perspective on low-metallicity environments with near galactic scale
resolutions. From current generation facilities, we can quantify their physics
and chemistry on sub-parsec scales and understand, in general terms, the
relative contributions of the different components of their ISMs. But
observations of [CII], for instance, suggest much of the molecular gas in the
Large Magellanic Cloud (LMC) is CO-dark: we are not detecting it using
standard sub-mm methods (Chevance et al., 2020). Understanding the molecular
gas content of the LMC then leads us to understand how much star formation can
be expected, and whether that is driven by the contents of the observed giant
molecular clouds, or the ISM surrounding them. Current facilities do not have
the sensitivity, FoV and resolution to explore the CO-dark gas in the MCs.
With the ability to observe in the sub-mm, AtLAST will be capable of observing
[CI] lines (at rest-frame frequencies of 492 and 809 GHz) across the MCs that
trace this CO-dark gas.

Figure 6: Dust emission at 100 $\mu$m in the LMC from the Herschel space
observatory (Meixner et al., 2013) The left panel shows the full $\sim
7^{\circ}\times 7^{\circ}$ area of LMC. The grid boxes represent the
$1^{\circ}\times 1^{\circ}$ fields of view proposed for the AtLAST continuum
camera. The right panel shows the zoom-in $30^{\prime}$ view around the 30
Doradus Hii region. The fields of view of AtLAST, ALMA (12-m array and 7-m
array), JWST NIRCam and Roman Space Telescope WFI are shown for comparison.
Only the future-generation AtLAST and Roman are able to map the full near to
far-infrared and sub-millimeter dust emission in the LMC. (Figure reproduced
from Liu et al., 2024)

Quantifying the dust (in terms of mass, density and temperature) as a function
of the environment, energy injection and metallicity tells us about dust
properties in extreme environments. There are discrepancies between
theoretical expectations of dust properties at low metallicity and those seen
in the MCs (Bot et al., 2010; Galliano et al., 2011). Resolving this
discrepancy pushes us closer towards a complete model of cooling, heating,
star formation and dust grain growth in different metallicity environments.
Understanding low-metallicity environments in nearby galaxies allows us to
improve our understanding of high-redshift galaxies, which are typically also
lower metallicities.

Pushing out to larger distances, the physics and chemistry of the ISM in
nearby galaxies tells us about the evolution of that galaxy, and by deriving
both the star formation rate and gas surface density in nearby galaxies (the
Kennicutt-Schmidt relation; Schmidt, 1959; Kennicutt, 1989), we understand
different types of galaxies in the context of the general population of
galaxies. These types of relations break down on roughly kpc scales, and
understanding that deviation from the KS relation puts interesting constraints
on the timescales of star formation. Through observations of chemical tracers
tracing different densities and in multiple transitions, we can further
explore the parameter space in the KS relation, probing to higher and higher
gas densities. To do this well, in multiple chemical tracers, across
statistically significant samples of galaxies, requires highly multiplexed
spectroscopic instruments with high sensitivity and resolution.

Threading through these themes is the ability to measure the polarised
emission of these galaxies to quantify how magnetic fields influence the
evolution of galaxies. Magnetic fields regulate star formation, influence the
dynamics of the ISM and are important for a number of other processes, but
observing them consistently remains a significant challenge. Small samples of
polarimetric observations of dust in nearby galaxies have shown a
relationships between the fields in the molecular gas and star forming
regions, but measurements are difficult because of either the small fields of
view of sensitive interferometers like ALMA, or the relative insensitivity /
low resolution of current generation single-dish facilities where observing
polarisation fractions to uncertainties of a few percent requires a
significant investment of telescope time.

Creating a polarisation survey of 100 nearby galaxies, reaching resolutions of
20-300 pc to an rms noise level of 16 $\mu$Jy/beam will enable detections of
1% polarisation in the dust emission. This is 100 times deeper than possible
with current facilities while at the same time offering significantly higher
spatial resolution (by a factor of $\sim$3 in each direction) which means more
than 1000 independent polarisation vectors can be measured. While impossible
with current facilities, observations like this could be done with AtLAST in
18-50 hrs per galaxy.

###  2.3 Our Galaxy

See Klaassen et al. (2024) for further details.

From the vantage point of our Solar System being in the plane of the Milky
Way, we are able to probe the cycle of interstellar matter in a detailed way
that is impossible to capture for external galaxies. We can view the assembly
of clouds on the largest scales down to dense cloud core scales and
protostars, then follow the subsequent dispersal of clouds by the feedback of
the formed stars and the eventual enrichment from the death throes of stars at
the end of their lives. We can delve into the minutiae, but to put those
details into their galactic context requires large area surveys. With a large
aperture, wide-field of view (sub)millimeter single dish telescope, we can
explore and quantify the spatial distribution of gas and dust, the kinematics
of the gas, the gas chemical composition (and its evolution) and magnetic
fields from the scales of giant molecular clouds down to individual nearby,
faint debris discs, analogous to our own Kuiper belt.

Large scale surveys of the gas and dust in the Galactic Plane unlock a
multitude of science questions. How do cool Giant Molecular Clouds (GMCs) form
within the largest scale structures in the ISM of our galaxy? How do those
GMCs then continue to fragment into filaments and star forming clumps and
cores? What then triggers or inhibits star formation in those clumps and
cores? What is the interplay between magnetic fields and stellar feedback on
star-formation? From these questions, we can see that the chemistry, dynamics,
and recycling of material are all interrelated in our understanding of the
ecology of our Galaxy and tie directly to the gas, dust and magnetic field
properties of the ISM. The ISM is turbulent and magnetized and to describe the
evolution of the ISM and star and planet formation we need to probe the
density distribution of matter as well as the velocity structure (to probe the
dynamics) and the magnetic fields. AtLAST will be able to study simultaneously
the evolution of dust and gas from large scales to small (from $\sim$100 pc to
$<$0.1 pc) and from low densities to high (1 to 106 cm-3) spanning orders of
magnitude in scale and densities.

A polarised dust continuum survey of the Galactic plane will allow us to
constrain, in a statistical and holistic way, the relationship between gravity
and magnetic fields in the collapse and support of structures, from the scales
of entire GMCs down to those of star-forming cores. This will allow us to
understand the flow of material through and between states, and the nature of
star forming cores throughout the plane – as we will reach from filament to
individual core scales (0.1 pc) resolution on the far side of the Galaxy at
the highest frequencies obtainable with AtLAST.

Spectroscopic surveys of the Galactic plane add to the picture through
chemistry, turbulent velocity dispersions, temperatures and heating and
cooling rates of the gas. Wide bandwidths will be essential for observing
multiple species and transitions simultaneously. Here we can use a similar
‘wedding cake’ approach to the Distant Universe studies above. In a wedding
cake scenario, we begin with larger area surveys (i.e the whole Galactic
Plane) to sensitivities sufficient to detect common molecular and ionised gas
transitions (CO, CS, N2H+, HCO+, SiO, radio recombination lines, etc), and
then follow up with deeper / more sensitive observations in a smaller tier
towards regions deemed chemically complex (i.e. individual star forming
regions) in order to derive a census of complex organic molecules (defined as
carbon bearing species with at least six atoms). Only following up with highly
sensitive observations in regions known to be chemically complex will save
hundreds of hours of observing time.

With the resolution, sensitivity to the full range of spatial scales, and
confusion limits enabled by a 50m dish antenna, we can study the most common
mode of star formation: low mass stars in clustered environments which are
also in the process of forming high mass ones. With current technologies, we
can study the details of isolated core collapse in nearby low mass star
forming regions, but with a 50m class telescope, the same resolution achieved
in nearby (100-200 pc) regions will be possible in the richer and more
complicated high-mass star forming regions within 2 kpc - where a full Initial
Mass functions (IMF) worth of low and intermediate mass stars are also forming
but cannot currently be detected or disentangled from their brighter
companions.

Figure 7: Observations of $\epsilon$ Eridani with LMT (Chavez-Dagostino et
al., 2016) and ALMA (Booth et al., 2023) along with a predicted image using
AtLAST for observation with an on source time equivalent to the ALMA image.
For a nearby disc like this, AtLAST can reach scales of a few au much more
rapidly and without loss of large-scale flux that hinders such observations
with ALMA. (Figure reproduced from Klaassen et al., 2024)

With large format cameras, entire nearby star forming regions can captured in
a single pointing, taking a single ‘snap-shot’. This makes it possible to not
only trace the dust and gas in the star forming cores that then settles into
protoplanetary discs and as the stars emerge and finally evolve into debris
discs as the gas is dispersed from the system, but to go back to these regions
in a set cadence to monitor for transient behaviour (i.e. accretion flares).
With (sub-)mm observations of the early evolution of these planetary systems,
the masses, morphologies, grain-size distributions and dispersal mechanisms
can be tracked (Lovell et al., 2021). With repeated observations of these
star-forming regions we can also expect to find large numbers of variable
sources, informing us of the star formation process (see also Section 2.6).
The mass assembly of protostars is known to be episodic and for the youngest,
most deeply embedded sources, during the main mass assembly stage, brightness
variability in the far-IR and sub-mm provide the only direct measure of both
the timescale and amplitude of these accretion episodes - only AtLAST has the
ability to monitor 1000s of sources with monthly cadence over decade
timescales.

As the circumstellar discs evolve into debris discs, their dust masses greatly
decrease, meaning that only the brightest debris discs are detected (e.g.
Wyatt, 2008). Observing discs as faint as the Solar System’s Kuiper belt,
requires observing the nearest stars (within around 20 pc) down to
sensitivities on the order of 10s of $\mu$Jy/beam. Given their proximity, the
discs can easily extend well beyond the maximum resolvable scale of
interferometers like ALMA and so only a large single-dish telescope is capable
of detecting them (see Figure 7 for an example). Observations in multiple
bands then allow the characterisation of their dust properties (e.g. Löhne,
2020).

###  2.4 Solar System

See Cordiner et al. (2024) for further details.

Understanding our place in the Universe requires understanding the formation
and evolution of our own Solar System. (Sub-)mm spectroscopy of giant planets
teaches us about the chemistry and dynamics of their atmospheres, from which
we obtain information about their origins and temporal evolution (Encrenaz et
al., 2005). From spectroscopy of terrestrial planets, moons and comets, we can
learn about their habitability and the availability (and origins) of chemical
ingredients for life in the Solar System. Deeper insights into our own
planetary system are also essential in our quest to understand the properties
of exoplanets, and the relationship of our solar system to other planetary
systems.

Figure 8:  Angular diameters of major Solar System bodies and Saturn’s largest
moon, Titan. The vertical extents of the coloured bars for each body represent
their range of angular sizes due to differing geocentric distances throughout
the year. The second y-axis shows the angular resolution a 50 m single dish
facility would achieve for a range of representative frequencies. (Figure
reproduced from Cordiner et al., 2024)

Observations with current (sub-)mm facilities have produced ground-breaking
results, including measuring the molecular inventories and dynamical motions
of planetary and cometary atmospheres. However, current observatories suffer
from limited dynamic range, spatial and temporal coverage, and sensitivity
(Cordiner et al., 2024). Instantaneous mapping over the entire face of a body
is critical for detailed studies of the rapidly rotating and evolving
atmospheres of solar system bodies (Lellouch et al., 2019; Cordiner et al.,
2020).

Interferometric observations are not ideally suited to solar system studies
due to the spatial filtering of the large angular scales that are a common
feature of planets and comets. Furthermore, the repositioning of antennas (and
consequent variation in $uv$ coverage) hinders our ability to reliably and
consistently study and map temporally variable phenomena such as storms,
outbursts or seasonal changes.

In order to advance our knowledge of planetary atmospheres, a large aperture,
single-dish telescope is necessary. This will allow us to measure spatial and
temporal variations in temperature structure and molecular abundances
throughout the troposphere and stratosphere, which can be used to investigate
the circulation, dynamics and composition of the atmosphere. A sensitive
single-dish facility will enable short-term and seasonal variations to be
studied, and facilitate the discovery of new, trace molecular species and
their distribution. A 50 m dish is essential here so that Venus, Mars, Jupiter
and Saturn can be resolved throughout their orbits and across all wavelengths
(see Figure 8). With this, Uranus and Neptune can also be marginally resolved
at the shortest wavelengths. In order to conduct sensitive searches for new
spectral lines and isotopologues of interest, a spectroscopic dynamic range of
$\sim 10^{5}$ is needed due to the continuum brightness of the planets — to
realise this, close attention should be paid to achieving the flattest
possible spectral bandpass through the careful choice of optical design,
receiver components and optimal calibration strategies.

Moons of the giant planets also offer compelling targets in the search for
life elsewhere in the solar system. Titan has a thick, carbon and nitrogen-
rich atmosphere that can be studied in a similar way to the terrestrial
planets (Nixon, 2024). Deep spectral line observations will enable the
discovery of new organic molecules (Cordiner et al., 2015; Palmer et al.,
2017; Thelen et al., 2020), as well as isotopologues of known molecules that
improve our understanding of the long-term physico-chemical evolution of
Titan’s atmosphere (Mandt et al., 2009, 2014; Nomura et al., 2023). Many of
the icy moons of the outer planets are thought to have subsurface oceans,
which can result in cryovolcanism (Nimmo & Pappalardo, 2016). Spectral studies
of the resulting plumes provide the perfect opportunity to study the
compositions of the subsurface oceans, in particular searching for organic
molecules that can reveal insights into the habitability of these oceans. The
Cassini mission to Saturn discovered such plumes around Enceladus (Hansen et
al., 2006; Porco et al., 2006; Spencer et al., 2018). Although it had a mass
spectrometer, this lacked the resolution to differentiate between certain
molecules (Waite et al., 2006, 2009). Remote observations from ground-based
facilities present the opportunity to overcome this. However, to do so
requires a high enough angular resolution (specifically $<8^{\prime\prime}$ at
1.1 mm) to avoid interference from Saturn and its rings, and with a single-
dish to avoid any interferometric artifacts. This would enable both
confirmation of the Cassini result and continued monitoring for temporal
changes in the plume composition.

Comets are thought to be largely unchanged since they accreted very early in
the Solar System’s history (Mumma & Charnley, 2011) and thus provide unique
information on the physical and chemical conditions at the time of planet
formation (Bockelée-Morvan et al., 2004). The ices and organics contained
within comets may well have played a key role in setting the stage for life on
Earth due to their delivery via impacts. Studying comets is also useful for
understanding how our Solar System relates to other planetary systems and to
the chemistry of the ISM (see Section 2.3). A huge variety of different ice
and organic lines are detectable in the (sub-)mm, necessitating a 32-64 GHz
bandwidth to be able to observe them simultaneously and thus avoiding changing
tunings during time-critical observations. Of particular interest is the
deuterium to hydrogen ratio that varies across Solar System bodies and can
tell us about the history of water in the Solar System, particularly when
compared to observations of protoplanetary discs (Willacy & Woods, 2009;
Albertsson et al., 2014; Cleeves et al., 2014). The strongest HDO line
observable from ground-based (sub-)mm observatories is the 894 GHz line,
demonstrating the importance of reaching these high frequencies. Cometary
comae are often extended over a few arcminutes, making studies of the extended
coma impossible with interferometry. In order to resolve the complex
kinematical structure of the coma, a spectral resolution of $\sim 0.1\,\rm
km\,s^{-1}$ is necessary.

###  2.5 The Sun

See Wedemeyer et al. (2024) for further details.

The Sun is the only main sequence star we can spatially resolve, and as such,
not only does it allow us to understand the engine that drives our solar
system, but gives us insights into other stars that would be impossible to
observe in such detail. At (sub-)mm wavelengths, the emission from the Sun is
dominated by the magnetised plasma in the chromosphere, which we observe as
(polarised) thermal continuum emission. Different layers of the chromosphere
dominate at different wavelengths, hence observations across the (sub-)mm
regime probe the three-dimensional structure of the Sun’s atmosphere.
Observations of the chromosphere are important for a number of reasons,
including understanding the ‘coronal heating problem’ – understanding why the
corona is much hotter than the solar atmospheric layers below it (including
the photosphere and chromosphere), solar flares, coronal mass ejections
(CMEs), prominences and the solar wind. By disentangling the importance of
different chromopsheric processes (such as magnetic reconnection and various
forms of wave heating), we start to unpick which types of physical processes
are dominant under which conditions (i.e. active or quiet Sun) and on what
size scales.

Solar flares, prominences and the solar wind all play a vital role in space
weather. Predicting them and understanding what drives them will help shape
space weather forecasting, but doing that requires detailed long term
observations of the Sun in polarised light. Because these transient phenomena
are related to the magnetic field reconnecting and reconfiguring on sometimes
explosive timescales, simultaneous observing of the full face of the Sun is
required, with repeated observations to take in variations across the solar
rotation period (around a month) and solar cycle (11 years).

Space weather directly impacts human society due to the effect it has on power
grids, satellite infrastructure and human spaceflight. The more we understand
space weather, the better we can mitigate these effects. Understanding space
weather in our own Solar System also allows us to better understand space
weather in extrasolar planetary systems and how it impacts habitability in
those systems, particularly in cases where flares and CMEs are considerably
stronger than those released from the Sun (e.g. Maehara et al., 2012).

Current generation facilities often either do not have the spatial resolution
to resolve structures like granules (Kundu, 1959; Trottet et al., 2011;
Loukitcheva et al., 2014) or, as is the case with ALMA and other
interferometers, filter out much of the full disc scale (30′) emission. For a
dynamical structure like the Sun, temporal resolution is just as vital, with
chromospheric evolution sometimes happening on timescales shorter than a
second (Kontar et al., 2018).

###  2.6 Transients and Variability

See Orlowski-Scherer et al. (2024) for further details.

As can already be seen in some of the science cases presented above, another
dimension has emerged as important in understanding the universe as it evolves
around us: time. The study of transient and variable phenomena is still in its
early stages in the sub-mm regime, with the first papers on the subject
published less than 10 yrs ago, but already we are seeing evidence for
variability in objects from AGN and Gamma-Ray Bursts (GRBs), through to
supernovae and forming stars and the atmospheres of the Sun and planets.
Transient and variability studies can help shape theoretical predictions for a
wide range of astrophysical phenomena, but care must be taken to ensure the
proper types of observations (e.g montoring, triggered ToOs and serendipitous
discoveries), and on the right cadence, for time domain astronomy to come to
prominence.

Some transient detections have come from the sub-mm (Guns et al., 2021; Lee et
al., 2021; Li et al., 2023), however most are followups of detections at other
wavelengths. With the right facility, software and operations model, AtLAST
could instead be the source of transient detections which will require
followup and monitoring by both AtLAST and other facilities. Take, for
example, the JCMT transient study of nearby star forming regions (Herczeg et
al., 2017; Mairs et al., 2017, 2024). These studies have shown the variability
of local protostars, and the observatory is continuing to monitor the same
regions to detect longer periods of variability. However, current facilities
are hampered by small fields of view and sensitivity. The former limiting how
many targets can be monitored, the latter limiting how precise the
measurements can be as well as their overall cadence.

Variability can happen on timescales of minutes to decades. Observing
strategies need to be in place to ensure that proper monitoring cadences are
maintained. In addition to planned observations of variable sources, the high
sensitivity provided by a large aperture telescope greatly increases the
chances of serendipitous detections found in datasets intended for other
purposes. To ensure that these are picked up on short timescales to enable the
triggering of follow-up observations, real-time data reduction pipelines for
transient detection are also required. These types of commensal observations,
allow for additional science goals to be met using observations designed often
for a very different purpose. Many of the upcoming CMB experiments are
expected to have such data processing software in place, which will already
start taking stock of the transient sky. With 100x the resolution, and
significantly better sensitivity, AtLAST will be able to detect more
transients that would get lost in the large beams of the CMB surveys.

Following up on transient and variable phenomena will necessitate observing
modes like target of opportunity (ToO) overrides, coordinated observations
with other facilities for multi-wavelength observations, monitoring campaigns
and the commensal observing described above.

###  2.7 Black Hole Event Horizons

See Akiyama et al. (2023) for further details.

The global VLBI network has played a significant role in our understanding of
black hole physics, showing the emission from lensed material very close to
the black hole itself for the first time (Event Horizon Telescope
Collaboration et al., 2019). The existence, nature and shape of this emission
in M87 and SgrA∗ definitively proved the existence of black holes at the
centre of galaxies and as the power sources for relativistic jets; provided
unique probes of the astrophysics of accretion and relativistic jet formation;
and enabled the exploration of properties of general relativity.

By increasing the network, the EHT will be able to image ever fainter sources,
particularly with the addition of a telescope as sensitive as AtLAST. Adding a
sensitive anchor not only improves the baselines that AtLAST is a part of, but
also those between other stations due to improvements in calibration. As a
single dish, AtLAST has major advantages over ALMA in this regard as ALMA has
to act as a phased array, which limits the sensitivity of sources it can
detect. AtLAST will also be able to host an instrument with multi-frequency
capabilities, enabling both simultaneous multi-frequency observations (e.g.
Issaoun et al., 2023) and order of magnitude improvements in sensitivity due
to the frequency phase transfer technique (Akiyama et al., 2023).

The images of M87 and SgrA∗ are right at the resolution limits of a 230 GHz
global VLBI network, yet, within those rings are an infinitely nested set of
separate rings (Johnson et al., 2020), each of which reveals new insights into
the nature of the black holes powering them. Shifting to higher frequencies
(690 GHz) will allow future EHT observations to resolve the n=0 and n=1 rings
of these black holes (at $\sim$ 6 microarcsec) and disentangle the
astrophysical and relativistic effects in that emission (e.g. Tiede et al.,
2022).

##  3 Telescope Requirements and Generalised Instrumentation

Figure 9: AtLAST computer aided design model. (Figure reproduced from
Mroczkowski et al., 2024)

The science cases presented in the previous section place various requirements
on the telescope, instrumentation and operations that we now discuss here. A
summary of the key requirements taken from each science theme is shown in
Table 2, with additional telescope requirements listed here which are implicit
in the science cases. Many of the telescope requirements described below are
already incorporated into the first designs of the telescope, as described in
(Mroczkowski et al., 2024; Kiselev et al., 2024; Gallardo et al., 2024; Puddu
et al., 2024), and shown in Figure 9.

###  3.1 Telescope requirements

####  3.1.1 Single dish

The large scale mapping described in many of the science cases requires the
use of single dish facilities because interferometers suffer from a lack of
sensitivity to large-scale emission; the size scales on which they start to
become insensitive typically being described as the maximum recoverable scale
(MRS). For ALMA in its compact (low resolution) configurations at sub-mm
wavelengths, this scale is of order a few arcseconds, much smaller than the
largest angular scales of many of the targets considered here, which range
from arcminutes (e.g. comet halos and debris discs) to degrees (e.g. GMCs
within our own galaxy, nearby galaxies, the CGM, galaxy clusters). This
problem can even impact scales smaller than the MRS (Plunkett et al., 2023),
in both spectroscopy and imaging (Hacar et al., 2018). For extended targets, a
single-dish telescope is necessary to ensure an accurate measure of the flux
and avoid any biases introduced by interferometric artifacts.

The stability in resolution offered by a single dish telescope is ideal for
consistent observations of time critical objects and those that require
cadenced monitoring. Interferometer arrays are often built to allow changes in
configuration, which is beneficial for enabling a range of angular
resolutions. The changes in configuration are typically time consuming, and
result in the arrays staying in a single configuration for a pre-determined
amount of time (e.g. ALMA spends about a month in each configuration). This is
problematic for time critical observations where the array may not be able to
provide the appropriate resolution at the requested time of the observations.
Seasonal monitoring of planetary atmospheres and the Sun are examples of such
cases that need a consistent resolution for periodic observations throughout
the year.

####  3.1.2 Dish diameter

The dish diameter impacts both the angular resolution ($\approx 1.2\lambda/D$)
and sensitivity ($\propto 1/D^{2}$) of observations. There is a fundamental
limit to how sensitive an observation an observatory can make, and that is
related to the so-called ‘confusion noise’ limit of the observatory, which is
directly related to its angular resolution. As described more in Appendix B,
as observations become more and more sensitive, fainter or more distant
objects begin to show up within an observed field of view. There comes a point
where there are so many objects detected that they can no longer be
distinguished from each other; they become confused within a single pixel. For
higher resolution observations, the confusion limit is lower, which means we
can push further. Current and planned observatories capable of observing at
sub-mm wavelengths reach confusion limits outside the galactic plane that mean
they cannot observe the ‘normal’ population of galaxies, only resolving the
brightest ones. In order to detect galaxies of luminosity L∗ out to $z=5$ and
resolve 97% of the CIB, a 50 m dish is necessary. With a dish of only 40 m,
the confusion limit would be twice as high, 15% less of the CIB would be
resolved and our ability to reach L∗ would only reach as far as $z=4$. This
confusion limit due to high redshift galaxies is important not only for
studying the galaxies themselves, but also for faint foreground targets, such
as nearby galaxies (both in terms of their ISM and CGM) and debris discs as
faint as the Kuiper belt around the closest ($<$10 pc) stars, where
sensitivities on the order of 10s of $\mu$Jy/beam at sub-mm wavelengths are
necessary, which is well below the confusion limit of current single-dish sub-
mm facilities.

A small primary beam ($\lesssim 8^{\prime\prime}$ at 1.1 mm, corresponding to
$\gtrsim$34m dish diameter) is strictly required for observations of Saturn’s
moon Enceladus, to avoid contamination and saturation from its close proximity
($<25^{\prime\prime}$) with the bright emission sources of Saturn and its
rings. The planets themselves can all be spatially resolved with a primary
beam $\lesssim 2^{\prime\prime}$ (possible with a 50m dish at 690 GHz), which
enables detailed atmospheric studies.

To disentangle star forming regions from their environments, and study their
interactions, requires the resolution to identify star forming clumps (1-0.2
pc) and cores (0.1 pc structures) within their natal clouds across the plane
of our Galaxy. With a 50m diameter, those resolutions can be achieved at the
wavelengths best suited to observing those structures out to 10 kpc (i.e.
beyond the central molecular zone): in low frequency N2H+ and HCO+ transitions
at $<$ 100 GHz we will achieve a resolution of $\sim 15^{\prime\prime}$ (0.7
pc at 10 kpc) towards cold star forming clumps, and in higher frequency CO, CS
and other species at $>$ 600 GHz we will achieve a resolution of $\sim
2^{\prime\prime}$ (0.1 pc at 1 0kpc) to probe the hotter and more chemically
rich star forming cores. Further afield, a 50 m primary allows us to resolve
GMCs ($\sim$300 pc in scale) out to 20 Mpc at 460 GHz (where we achieve $\sim
3^{\prime\prime}$ resolution).

In addition to these science drivers is the size required for combination with
ALMA data. In order to ensure overlapping coverage in the Fourier plane, a
single-dish telescope is required to have a diameter of at least
$D_{\rm{SD}}\geq 1.18L_{\rm{min}}/\kappa$ (Frayer, 2017) where $L_{\rm{min}}$
is the minimum baseline (14.6 m for the compact ALMA main array
configurations, Cortes et al., 2023) and $\kappa$ is a coefficient between 0.4
and 1, depending on the source distribution and signal-to-noise of the data.
If we conservatively take this to be 0.4 in order to cover edge cases, then we
find that a dish size of at least 43 m is necessary to facilitate full
coverage of the short baselines when combining with ALMA’s main array.

####  3.1.3 Field of view

Many of the key science cases described in this document require surveys with
a large spatial coverage either because they target large numbers of sources
or sources with a large angular scale. Surveys of high redshift galaxies fall
into the former category, with an expectation of finding around 50,000
galaxies per $1~{}\mathrm{deg^{2}}$ field in continuum observations and a
desire to cover $\geq 1000~{}\mathrm{deg^{2}}$. Whereas studies of the plane
of our own galaxy are expected to cover up to 540 deg2 (covering the full 270∘
plane visible from the southern hemisphere, with a degree above and below the
midplane), the CGM around nearby galaxies extends to a few degrees, SZ effects
have scales on the order of a degree, and the Sun is half a degree across.
Observing these large scale structures efficiently requires a large field of
view filled with detectors. In general, a larger filled FoV reduces the time
needed for surveys of such targets. It also necessitates calibration
strategies that do not lose that emission due to atmospheric fluctuations (see
next section); in the SZ case an instantaneous field of view covering
$>1~{}\mathrm{deg^{2}}$ is explicitly needed as a smaller FoV will not have
sensitivity to the degree scale structure essential for the project. Transient
and variable targets can appear anywhere on the sky and so a larger field of
view increases the chance of detecting such phenomena, especially in cases
where they are poorly localised triggered events, or are commensal (i.e.
transients to be detected by automated routines in observations derived from
other science goals).

####  3.1.4 Observing at Sub-mm wavelengths

Cold dust in the Universe has a peak rest-frame wavelength around 100 $\mu$m,
which is shifted to longer wavelengths at higher redshifts. The sub-mm
atmospheric windows therefore provide the only way to probe this emission at
or close to the peak of the distribution at high redshift from a ground based
observatory. Figure 10 also shows key atomic and molecular line frequencies as
they shift into the sub-mm bands as a function of redshift. To observe the CGM
of galaxies at z$<1$, emission too large scale and too diffuse to detect with
current generation facilities, requires the highest sub-mm wavelengths.

Figure 10: Evolution of the observed frequency of the brightest sub-mm and
far-infrared emission lines from CGM (orange/brown lines) as they get
redshifted into the sub-mm bands at different redshift values. The blue
spectrum in the background shows the top quartile of the atmospheric
transmission at the Chajnantor Plateau (about 5100 metres above sea level,
derived using the atmospheric am code, Paine 2019), where the corresponding
transmission is reported on the y-axis on the right. The coloured vertical
bands correspond to the ALMA bands (see Appendix A). The horizontal dashed
lines indicate the redshifts of a few sources discussed in the paper, and of
two additional high-z galaxies from the literature (D’Eugenio et al., 2023;
Hashimoto et al., 2018). This plot demonstrates the importance of a broad
wavelength coverage for the CGM science case, and it highlights how crucial
the high-frequency bands are to observe the brightest tracers of the cold CGM
in galaxies at $z>1$. (Figure reproduced from Lee et al., 2024)

The sub-mm atmospheric windows open up the possibility of observations of
redshifted spectral lines. For instance, being able to detect multiple tracers
of the CGM in the same source enables the study of the physical properties of
the gas, its density, temperature, metal content, mass, and so its origin.
Given the wide range of redshifts (from $z=0$ to $z=10$) that are of interest
for the CGM science case and studies of distant galaxies in general, a
telescope with sub-mm (e.g. $\lambda<600~{}\mu$m) capabilities would have
access to the brightest gas tracers of material in a larger volume of the
Universe (see Figure 10). For observations within our own galaxy, many
rotational transitions of important molecules emit at sub-mm wavelengths; some
tracing the thermal balance of the gas, others tracing its ionisation state or
the dynamics of the warmer gas traced by the higher energy rotational
transitions of molecules. For comets in the Solar System, access to the HDO
$1_{1,1}-0_{0,0}$ line at 894 GHz is useful to the study of D/H ratios that
can inform us of the origins of Earth’s water.

For these reasons, AtLAST needs to be located in a high dry site and needs to
have a high surface accuracy to access the sub-mm window.

####  3.1.5 Scanning speed

Follow-up of some transient sources that decay rapidly such as GRBs or
gravitational wave events requires rapid slewing to quickly get on target and
begin observing. A fast scanning speed is also necessary for calibrating out
atmospheric fluctuations and mapping large fields. These needs push for a
scanning speed of a few degrees per second, with one of the most stringent
requirements set by the need to detect faint emission on large scales that
would be drowned out by short term atmospheric fluctuations if not accounted
for properly; the ability to modulate the astrophysical signal faster than the
signal from Earth’s atmosphere (Morris et al., 2022; van Marrewijk et al.,
2024). For AtLAST, an average scan speed of 1.5 deg/s is needed to keep the
atmospheric signal stable in the time domain, with the ability to reach 3
deg/s needed for faster than average wind conditions (Mroczkowski et al.,
2024).

####  3.1.6 Pointing and Tracking

With many of the science cases described above, large scale mapping is key. To
ensure these maps are of high quality requires high relative and absolute
pointing calibrations (greater than the resolution of the telescope) and
tracking ability.

Ensuring proper pointing is also important for instruments with greater than
Nyquist spacings between their individual pixels which means that small maps
(i.e. jiggle, dither, etc) are required to fully sample the sky.

####  3.1.7 Beam accuracy

Accurate calibration is needed to properly account for the potential
systematics in the small-amplitude fluctuations of the SZ signal that are
associated with local pressure and velocity perturbations, or to relativistic
distortions. This results in a requirement of a sub-percent level control of
the beam stability. This is also important for cases that need a high dynamic
range such as observing moons near bright planets, planetary atmosphere
observations that need to distinguish the weak atmospheric lines from the
bright continuum or observations of the CGM, which is much fainter than the
ISM.

###  3.2 Instrumentation requirements

The science cases in this report require a mix of continuum and spectral line
observations including polarisation, solar observing and VLBI capabilities.

####  3.2.1 Continuum Observations

Cold dust throughout the Universe (from galaxies to circumstellar discs to
asteroids within our own Solar System), photons scattered by hot electrons due
to the SZ effect and emission from the atmosphere of our own Sun can all be
observed in their broadband continuum emission or as deviations from that
expected continuum. A wide wavelength coverage from 350 $\mu$m to 3 mm is
desirable to get a broad sampling of the spectral energy distribution (SED) of
the targets. Sub-mm observations, in particular, are crucial as these
wavelengths are closest to the peaks in the SEDs for many targets and due to
the improvements in resolution that they provide (see Section 3.1.4 for more
details). These continuum emission measurements often need to be combined with
observations at other wavelengths for the greatest scientific benefits. For
instance, when considering thermal emission from dust, multi-wavelength
observations enable information on the dust composition and size distribution
to be derived. At around 1 mm and longer, other contributions to the SED, such
as free-free emission, can dominate and so multi-wavelength observations allow
these distinct contributors to the overall continuum emission to be
disentangled. Multi-chroic observations are also important for distinguishing
between the thermal and kinetic SZ effects.

A large number of the science cases described above would benefit from
simultaneous multi-chroic (or multi-band) observations (see, for instance
figure 4). For transient and variable sources, simultaneity reduces the need
to distinguish between calibration and temporal effects, and in the case of
short timescale phenomena such as solar flares, it is necessary for multiple
wavelengths to be observed simultaneously to capture the evolution of the
flare in real-time. However, for deep observations (such as high redshift
galaxies or extrasolar Kuiper belts), multi-chroic observations may be less
efficient as the longer wavelength detectors will become confusion limited
much quicker than the shorter wavelengths (see Appendix B).

The bandwidth of the continuum camera is a key parameter for determining the
sensitivity of an observation, with wider bandwidths allowing for greater
sensitivity in a given integration time. To maximise observing efficiency,
getting bandwidths as close to the (sub-)mm atmospheric transmission windows
as possible would minimise the on source integration times (see Appendix A for
the bounds of those atmospheric windows).

As described in Section 3.1.3, a FoV of $>1~{}\mathrm{deg^{2}}$ is required
for many of the science cases due to either the large scale emission from the
targets (e.g. CGM and the Galactic plane) or the large number of targets (e.g.
surveys of distant galaxies) available to be observed. To properly take
advantage of that large field of view requires highly multiplexed
instrumentation. This opens up significant possibilities for synergy between
various science cases. For example, large area (1000+ deg2) continuum surveys
can be used to study distant galaxies and the clusters and protocluseters they
inhabit, whilst also picking up any nearby circumstellar discs within the
field as well as serendipitously detecting any transients or variables. A 300
000 pixel camera would be capable of a 2.2 degree FoV at 3 mm or provide a
0.26 degree FoV at 350 $\mu$m (if the pixels are well spaced). However, this
assumes one camera per wavelength. As some of the cases push for a multi-
chroic camera, this could result in trade-offs between FoV and number of
wavebands if the total number of pixels that can be placed in an instrument is
fixed (see, for instance, Figure 11).

Figure 11: Number of pixels in existing and expected continuum cameras on
current and planned facilities as a function of time. Extrapolating from the
trend seen in this plot, is is expected that a first light continuum camera
for AtLAST could have upwards of 106 pixels (as shown with red star in the top
right corner of the plot).

##### Solar Observing

Observations of the Sun also place requirements on the detectors as it is the
brightest mm source in the sky and can vary in brightness significantly, in
particular during flares. A solar filter with well-defined characteristics and
large dynamic range for the detector are needed and may be best served by a
specialised solar observing instrument. While integration times are very short
for a bright object like the Sun, spatial (or spectral) scanning in order to
cover an adequate spatial (or spectral) region is constrained by the short
dynamic timescales. The temporal cadence should be at most 1 min and ideally
less than 1 s as flares develop quickly. High detector readout rates are
required, especially for fast scanning. Given the particular requirements of
solar observing, a bespoke instrument may be better suited to fulfilling this
requirement.

####  3.2.2 Spectral Line Observations

Atomic, ionised and molecular gas fills the ISM of galaxies, extends beyond
them as part of their CGM and even populates the ICM of clusters. Studying the
spectral signatures of this gas tells us about the chemistry, temperature and
ionisation state within the gas and its dynamics both relative to other parts
of the galaxy and relative to us – allowing us to determine redshifts. Many
Solar System bodies are surrounded by atmospheres of gas, which is rich in
information about their origins and ongoing chemistry.

Large format spectrometers are necessary to observe the gas populations
described above and quantify ionisation states, dynamics, chemistry, evolution
and redshift (see, for instance figure 12). Highly multiplexed receivers are
key to any spectral science cases that aim to study wide areas, but different
types of spectroscopic instruments are better suited to the broad range of
these science cases. For many Galactic and nearby galaxy science cases, high
spectral resolution (of order $<0.1$ km s-1, or R $>10^{8}$) is required to
disentangle kinematics and hyperfine transitions of molecules that have
linewidths below 1 km s-1. Instrument design, processing power and data
transfer rates then limit the bandwidth that can be simultaneously observed
with such an instrument. For nearby galaxies, with much broader spectral lines
(of order a few 100 km s-1), the same instrument could be used, likely in a
lower spectral resolution mode, but with a greater need for stable
spectroscopic baselines to ensure detection of faint, non-Gaussian, emission
in the line wings. To ensure these baselines would require position switched
or other atmospheric corrective observing techniques. This is contrasted
against science cases dedicated to finding the redshifts of galaxies within
the FoV of the instrument where much lower spectroscopic resolution is
required (R$\sim$ few hundred). Because these science cases can get by with
lower spectral resolution integral field units (IFUs) they can simultaneously
observe much larger spectral bandwidths which enables covering larger redshift
ranges, which increases the efficiency of such ‘redshift machines’.

Based on current technologies, and extrapolating to future spectroscopic
instruments (see, for instance Groppi et al., 2019), 1000 pixel spectroscopic
instruments will become feasible by the time AtLAST has first light.

Figure 12: Footprints of various assumed instruments (from left to right): a
919-beam (18-layer) heterodyne receiver, a 217-beam (9-layer) heterodyne
receiver, an IFU with $50\times 50$ pixels, and an IFU with $900\times 900$
pixels, on top of the nearby galaxy NGC 3627’s CO(2–1) emission (Leroy et al.,
2021). The heterodyne receiver’s hexagon units and the IFU pixels are assumed
to have a spacing/size of $4^{\prime\prime}$, which is about the AtLAST 50-m
angular resolution at 400 GHz. A large-format IFU as shown in the last panel
can most efficiently observe multiple galaxies, with each galaxy fully sampled
at resolutions as shown in the first three panels. (Figure reproduced from Liu
et al., 2024)

####  3.2.3 Polarimetry

Polarimetric observations are important for understanding the interplay
between the various forces (i.e. gravity, turbulence, magnetic fields) at work
in the clouds and galaxies and clusters under study. Such capabilities for the
continuum and spectral cameras are a key requirement unlocking a hidden
dimension of understanding, particularly for studies of magnetic fields in
galaxies (including the Milky Way) and our Sun. Continuum polarimetry (as
shown in the left panel of Figure 1) gives an understanding of the
directionality of the dust grains in the object(s) being observed.
Spectropolarimetry preserves the polarisation information present in molecular
lines, and thus can differ from that in the dust continuum.

To achieve polarimetric observations not only requires polarisation preserving
instruments, but optical paths which do not confuse the polarisation signals.
The former can often be integrated into MKID and other detector-on-chip type
instrument technologies, while the latter needs to be incorporated into the
observatory at conceptual design time (see, for instance, Puddu et al., 2024).

####  3.2.4 Very Long Baseline Interferometry

The incorporation of a large, single-dish telescope capable of sub-mm
observations into EHT will significantly increase the number of viable targets
and the science return of the Event Horizon Telescope. For AtLAST to become a
VLBI station, a specialised heterodyne receiver is necessary that is capable
of high sensitivity observations at multiple frequencies (at least 90, 230,
345 and 690 GHz) with a wide bandwidth.

###  3.3 Telescope Operations

To address the ambitious and diverse scientific programme outlined in this
report, and to accommodate future user-driven research initiatives, a
versatile and sustainable operational scheme must be developed. This range of
programmes requires a flexible scheduling system with suitable observing
modes. Additionally, the powerful capabilities of AtLAST, such as its large
field of view, high sensitivity and multi-beam and multi-pixel
instrumentation, necessitate efficient methods for processing, storing, and
providing the community with high-quality data for optimal exploitation.

####  3.3.1 Observing modes

The variety of science cases require different types of observing strategies.
Telescope scheduling must accommodate requirements from large, automated
surveys of the sky to dedicated principal investigator (PI) projects requiring
only a few hours on sky, and everything in between. This diversity in
observing modes calls for fully dynamic scheduling to quickly adapt to the
weather conditions and to make the best use of the available time, allowing
for flexibility and on-the-fly changes of the schedule.

PI-driven programmes can range from a few to a few hundreds of hours (in the
form of Large Programmes), depending on the nature of the sources to be
observed and scientific goals being pursued. Depending on the case, granting
the PI the possibility to interact with the observatory staff, to check the
quality of the data while it is being taken and to refine the observing
strategy, when necessary, will help maximise the scientific output of the
facility.

To better meet the requirements for Solar studies, a combination of long-term
synoptic observations with a very low daily load on observing time with
campaign-based observations is needed. The observations of active regions
should be schedulable with just a few days warning, as the coordination with
other space-borne and ground-based observatories will optimise the scientific
result.

Transient and variable phenomena from the Sun and other Solar System objects
to GRBs and AGN are time critical and require special observing modes to be
considered. Target-of-opportunity observations, coordination with other
telescopes/space missions, and accommodation for regular monitoring of science
targets with a predefined cadence (over a timescale of hours, days, weeks,
months or years), will all be crucial for maximum science return on time-
variable sources. Scheduling the most critical observations to within an
accuracy of a few minutes will therefore be desirable. Responding to
unexpected transient phenomena such as cometary outbursts or gravitational
wave events will necessitate flexible telescope scheduling, for example, with
turnaround times from a few minutes to a few days.

If AtLAST is to not only follow-up transient triggers from other telescopes
but also to send its own triggers when transients are serendipitously
detected, then a transient search pipeline with an automated alerting system
should be set up. This can be implemented with an “opt-out” option for PIs in
cases where their science goals include searches for transients.

Coordination with other telescopes will also be necessary for VLBI
observations. This places extra constraints on the operations with specialised
equipment in addition to the heterodyne instrumentation including a highly
phase-stable frequency standard, backend digital electronics for recording
data at rates up to 256 Gbps and specialised control software.

####  3.3.2 Data flow and legacy value

AtLAST is expected to have a tremendous legacy value given the variety of
different imaging and spectroscopic surveys the facility will address. The
commitment is to enhance this by providing users with fast, reliable and
useful access to AtLAST data, as well as maintaining high data quality
standards. This requires supplying an adequate infrastructure to store the
data and give easy access to the community. The final goal is to make data
accessible, properly described and compatible with other analysis tools and
services that facilitate comprehensive research across wavelengths, in line
with the FAIR principles (Wilkinson et al., 2016).

The large format continuum and spectral line cameras onboard AtLAST will
generate large data volumes, especially if the spectral line cameras can
achieve a 0.01 km/s spectral resolution across a 32 GHz bandwidth. This will
require re-thinking calibration and data reduction techniques and implementing
more efficient and sustainable data processing and storage schemes.

AtLAST will not be alone with this problem. With ALMA’s Wideband Sensitivity
Upgrade (Carpenter et al., 2023), and the even larger fields of view possible
with the SKAO, data storage, reduction and analysis will become significantly
more challenging. This will open new opportunities, on the one hand to explore
and incorporate new technological developments, like novel distributed
computing systems or federated cloud services, and on the other hand
establishing synergies among different research infrastructures facing these
common challenges.

##  4 Summary

In this document, we have summarised the key science cases developed for
AtLAST by the community through the consultation process held during the
AtLAST design study. The community consolidated around 28 use case submissions
in the first part of the consultation (as summarised in Ramasawmy et al.,
2022) and further delved into the details of these science cases in five
working groups which produced the 7 science white papers summarised in Section
2. The technical requirements these science cases place on the observatory
were then discussed in Section 3, with the key science drivers for AtLAST,
derived from these consultations presented in Table Executive Summary. An
overview of the technical constraints from the overarching science themes are
consolidated into a matrix of requirements in Table 2.

The strong support of the community has been vital to this work, with experts
in various (sub-)mm domains volunteering their effort towards understanding
the technical needs of future facilities in order to advance our understanding
of the Universe around us. Similarly, the thoughts and views and priorities of
this community will evolve along with our understanding. As such, this
document should be viewed as a snapshot of the anticipated science return that
an observatory like AtLAST could provide. As new and better observatories and
instruments come online, models improve and science evolves, so to will the
telescope requirements for enabling great science. The future evolution of the
telescope design must take this into account and continue refining its science
goals in the light of new discoveries.

|  |  The Sun |  Solar System |  Milky Way |  Nearby Galaxies |  CGM |  Distant Universe |  The SZ effect |  VLBI |  Transients
---|---|---|---|---|---|---|---|---|---|---
Telescope |  |  |  |  |  |  |  |  |
| Diameter (m) |  $>$50 |  $>$34 |  $>$50 |  $>$53 |  $>$50 |  $>$50 |  $>$43 |  $>$50 |  $>$50
| Field of View (deg) | 0.5 | 0.2 |  $>$1 | 1 |  $>$1 |  $>$1 |  $>$1 | - |  $>$1
|  LASa (deg) | 0.5 | 0.2 | 10 | 70 | 1 | - | 1 | - | -
|  Angular resolutionb (′′) |  $<$5 |  $<$7 |  $<$5 | $<4.5$ |  $<$5 |  $<$5 | $<6$ |  (6e-6)d | -
|  Spatial Coverage (deg2) | 0.25 | 70.5 | 540 | 1000+ | 1000+ | 1000+ | 1000+ | - | 540
Wavelengths |  |  |  |  |  |  |  |  |
|  Atmospheric Windowsc | 1-10 | 3-10 | 3-10 | 3-10 | 1-10 | 1-10 | 1-10 | 3-9 | 1-10
Spectral Requirements |  |  |  |  |  |  |  |  |
| Spectral observations |  | ✓ | ✓ | ✓ | ✓ | ✓ |  | ✓ | ✓
| Polarisation |  |  | ✓ |  |  | ✓ |  | ✓ | ✓
| Spec. Bandwidth (GHz) | - | 64 | 16-32 | 32 | 4-8 | 400 | - |  $>$8 |  -e
| Spec. Resolution (km/s) | - | 0.1 | 0.01 | 0.5 | 5 | 400 | - | - |  -e
| Spatial Sampling | Full | Full | Full/MOS | Full/MOS | Full | Full/MOS | - | - |  -e
Continuum Requirements |  |  |  |  |  |  |  |  |
| Bandwidth | Full atmospheric window
| Resolution | 1+ channel per band
| Spatial Sampling | Full | Full | Full | Full | Full | Full | Full | - | Full
| Polarisation | ✓ |  | ✓ | ✓ |  |  |  |  | ✓
Timing |  |  |  |  |  |  |  |  |
| Repeats | ✓ | ✓ | ✓ | ✓ |  |  |  |  | ✓
| Time Variability | ✓ | ✓ | ✓ | ✓ |  |  |  |  | ✓
| Targets of Opportunity | ✓ | ✓ | ✓ | ✓ |  |  |  |  | ✓
| Cadenced observations | ✓ | ✓ | ✓ | ✓ |  |  |  |  | ✓
| Time Critical | ✓ | ✓ | ✓ | ✓ |  |  |  |  | ✓
Table 2: Generalised telescope and instrument requirements defined by the
science working groups. aLargest angular scale. bAngular resolution at 1 mm.
cThe Atmospheric windows are listed with respect to the ALMA observing bands,
for which a translation to wavelength/frequency is found in Appendix A. d
Resolution of VLBI observations not set by AtLAST alone. e Spectral and
spatial sampling requirements vary according to the type of variable target
being observed within the other science categories presented here.

##  Appendix A Frequency Ranges of ‘ALMA Bands’

The current and anticipated receivers on the ALMA antennas are well aligned to
the (sub-)mm atmospheric windows. The values in Table 3 describe the rough
edges of the frequency bands defined within the ALMA community, and correspond
to the vertical bands presented in Figure 10.

ALMA Band | Wavelength | Frequency | ALMA Band | Wavelength | Frequency
---|---|---|---|---|---
| (mm) | (GHz) |  | (mm) | (GHz)
1 | 6-8.5 | 35-50 | 6 | 1.1-1.4 | 211-275
2 | 3.3-4.5 | 65-90 | 7 | 0.8-1.1 | 275-373
3 | 2.6-3.6 | 84-116 | 8 | 0.6-0.8 | 385-500
4 | 1.8-2.4 | 125-163 | 9 | 0.4-0.5 | 602-720
5 | 1.4-1.8 | 163-211 | 10 | 0.3-0.4 | 787-950
Table 3: Mapping of ALMA Observing bands to their Atmospheric windows

##  Appendix B Confusion Limit

‘Confusion’ refers to the issue of lacking the angular resolution necessary to
spatially distinguish between sources on the sky that are physically separate.
The deeper an observation, the more sources are detected and the worse this
problem becomes leading to a limit on the sensitivity which is referred to as
the confusion limit. Confusion can come from both Galactic and extragalactic
sources but in the (sub-)mm, the main source of confusion is distant galaxies
(e.g. Blain et al., 2002). There are then multiple criteria that can be used
to define the confusion limit (e.g. Dole et al., 2003; Ermash et al., 2020).
Here we consider both the photometric criterion, which is most relevant when
considering the confusion as a source of background noise, and the source
density criterion, which is most relevant when studying the galaxies
themselves.

The confusion defined by the so-called photometric criterion corresponds to
signal fluctuations in the beam of the observation due to extragalactic
sources below the detection threshold $S_{\rm lim}$. Following Dole et al.
(2003), it is given by

$\sigma_{\rm c}=\int_{\Omega}f^{2}(\theta,\phi)\,d\theta d\phi\int_{0}^{S_{\rm
lim}^{\rm c}}S^{2}\frac{dN}{dS}dS,$ (1)

where $f(\theta,\phi)$ is the two-dimensional beam profile, $S$ is the source
flux density, $dN/dS$ is the differential number counts, $\sigma_{\rm c}$ is
the photometric confusion noise and $S_{\rm lim}^{\rm c}$ the photometric
confusion limit. These last two variables are then implicitly linked by

$S_{\rm lim}^{\rm c}=q_{\rm phot}\times\sigma_{\rm c},$ (2)

where $q_{\rm phot}$ measures the desired photometric accuracy and is usually
taken equal to 5. As the calculation of this photometric confusion limit
relies on the knowledge of differential number counts down to infinitely low
flux densities, it requires the use of a galaxy population model. Here we used
the empirically-motivated galaxy population model of Béthermin et al. (2017),
which synthesizes state-of-the-art constraints on the extragalactic cosmic
infrared background (CIB), reproducing a wide variety of number count
measurements in the mid-infrared (e.g., Spitzer), far-infrared (e.g.,
Herschel), and (sub-)mm (e.g., SCUBA2). The extragalactic photometric
confusion limits measured for a 50 m telescope assuming $q_{\rm phot}=5$ are
summarized in Table 4. The very large diameter of AtLAST gives very low
photometric confusion limits, of the order of a few tens of micro-Jansky, or
even lower than the flux density range of our model in the two AtLAST bands
with the highest angular resolution (i.e., 350 and 450 $\mu$m). These are
about 100 times lower limits than the photometric confusion limits reached by
current 10-m class single-dish (sub-)mm telescopes or even $>10\,000$ times
lower than those of the Herschel space telescope. Only the LMT can match these
low photometric confusion limits, but it is limited to a very small field of
view ($\sim 8^{\prime}$) and lacks the essential sub-mm coverage.

The confusion defined by the so-called source density criterion (SDC; e.g.
Dole et al., 2003) is driven by the fraction of sources lost in the detection
process because the nearest neighbor whose flux is greater than $S_{\rm
lim}^{\rm SDC}$ is too close to be separated. Unlike the photometric
confusion, the SDC confusion does not act as a background noise against which
sources must be detected, but it affects the completeness of the catalog of
sources of interest. It is therefore more relevant when studying the sources
themselves and, for example, the SDC confusion limit set by distant galaxies
does not directly affect galactic science (although galactic source catalogs
are also affected by the same phenomena). Following Dole et al. (2003), the
SDC confusion is defined using the probability $P$ given a source density
$N_{\rm SDC}$ (Poisson distribution) of having the nearest source of flux
$S_{\rm lim}^{\rm SDC}$ located at a distance less than $\theta_{\rm min}$,

$P(<\theta_{\rm min})=1-e^{-\pi N_{\rm SDC}\theta_{\rm min}^{2}},$ (3)

where $\theta_{\rm min}$ is usually defined in term of the FWHM of the beam as

$\theta_{\rm min}=k\,\theta_{\rm FWHM}.$ (4)

It follows that $N_{\rm SDC}$ is given by

$N_{\rm SDC}=-\frac{{\rm ln}\,(\,1-P(<\theta_{\rm min})\,)}{\pi
k^{2}\theta_{\rm FWHM}^{2}}.$ (5)

By choosing values for $P(<\theta_{\rm min})$ and $k$, $N_{\rm SDC}$ can be
calculated and then translated into $S_{\rm lim}^{\rm SDC}$ using a source
number count model. Although $P(<\theta_{\rm min})$ depends on the desired
completeness and $k$ on the exact source extraction method used, it is
generally set to 10% and 0.8, respectively (Dole et al., 2003). This gives a
$N_{\rm SDC}$ of 1 source per 16.7 beam area. Using the galaxy population
model of Béthermin et al. (2017), we translated this source density into the
extragalactic SDC confusion limits for a 50 m telescope (see Table 4). In most
AtLAST bands, the extragalactic SDC confusion limits are greater than those
deduced from the photometric criterion, with the transition between the
photometric and source density criteria occurring only at 3 mm. At these
confusion limits, a very large fraction of the extragalactic CIB is
nevertheless resolved into individual galaxies (see the third column of Table
4), enabling unprecedented advances in our understanding of galaxy evolution
(see, e.g., Section 2.1.1). We also note that these are rather conservative
SDC confusion limits, as source extraction methods using prior information on
source position (from, for example, complementary optical/near-infrared
surveys) can separate sources down to $k=0.5$ (e.g., Magnelli et al., 2009,
i.e., a $N_{\rm SDC}$ of 1 source per 6.5 beam area). This lowers the
extragalactic SDC confusion limits by a factor $\sim 3$ and leads to a
transition between the photometric and source density criteria at 1.3 mm. In
this more optimistic case, more than $80\%$ of the extragalactic CIB is
resolved in the AtLAST bands shortwards of 1.3 mm.

Our calculations unambiguously demonstrate the ability of a 50 m telescope to
achieve unrivalled low confusion limits, enabling deep galactic observations
and resolving into individual galaxies most of the extragalactic CIB in the
0.35 to 1.3 mm range. As shown in Section 2.1.1, such deep observations over a
wide area of the sky ($>1\,000\,$deg2) will in particular enable us to study
the far-infrared/(sub-)mm emission of a large and comprehensive sample of
typical star-forming galaxies (i.e. at $L^{\star}$) from $z=0$ to $z\sim 5$.
By comparison, a 40 m telescope would have $\sim 1.8$ $\times$ higher
confusion limits, resolving in each band about 15% less of the extragalactic
CIB into individual galaxies, and allowing the study of typical star-forming
galaxies only up to $z\sim 4$.

Wavelength | Photometric Confusion Limit | SDC Confusion Limit | CIB Resolved
---|---|---|---
$\mu$m |  $\mu$Jy |  $\mu$Jy | %
350 |  $<$1 | 56 | 97
450 |  $<$1 | 112 | 92
750 | 30 | 195 | 76
850 | 34 | 166 | 70
1100 | 44 | 141 | 61
1300 | 65 | 138 | 50
2000 | 57 | 68 | 36
3000 | 39 | 31 | 16
Table 4: Extragalactic confusion limits based on the photometric (assuming
$q_{\rm phot}=5$) and source density (SDC; assuming $P(\theta_{\rm min})=0.1$
and $k=0.8$) criteria. The third column gives the fraction of the
extragalactic CIB resolved into individual galaxies at the higher of these two
limits (highlighted in bold). These calculations are based on the galaxy
population model of Béthermin et al. (2017).

##  Appendix C AtLAST Sensitivity Calculator and Simulator

To allow astronomers contributing to the science working groups the ability to
understand whether their planned observations were realistic, and to ensure
consistent calculations were being done across working groups, a sensitivity
calculator and telescope simulator was developed.

Having a single source of truth, from which astronomers can derive integration
times required to reach their desired sensitivity (or conversely, achieved
sensitivity in a given on-source integration time), is vital for ensuring that
astronomers are comparing like-for-like when deriving observing cases.

At such an early stage in the development of the telescope, when so many
details (large and small) are still under constant review, the goal of the
sensitivity calculator was to be as flexible as possible, even including the
ability to modify the diameter of the telescope when running the calculator
from the source code.

The sensitivity calculator was designed with simplicity and flexibility of use
in mind. As such, users can interact with it either through a simplified web
interface, with most telescope parameters fixed, or at the command line (or
within a python instance) and configure any and all of the input parameters to
the sensitivity calculation to their specification. The former lowers the
barriers to participation by removing the complexity in setting up the
calculation, the latter allows expert users to explore the parameter space
available to them through the calculation to better understand how varying a
telescope property (like dish diameter) will affect their results.

It is through the command line (or python) interface that users can create
simulated observations using an input FITS dataset (2D or 3D), with examples
of how to do that, in the form of Jupyter Notebooks, provided within the
sensitivity calculator source code. With these tools, astronomers can simulate
observing with an idealised telescope under various weather conditions with
the caveats and assumptions provided below.

###  C.1 Underlying Equations, Assumptions and Caveats

As explained in the sensitivity calculator documentation, we used the system
equivalent flux density (SEFD) of a telescope to derive the sensitivity
achievable in a given integration time or vice versa:

$\Delta S=\frac{SEFD}{\eta_{s}\sqrt{2\Delta\nu t}}\hskip
20.0pt\textrm{or}\hskip 20.0ptt=\left(\frac{SEFD}{\Delta
S\eta_{s}}\right)^{2}\frac{1}{n_{pol}\Delta\nu}$ (6)

where $\eta_{s}$ is the system efficiency, $n_{pol}$ is the number of
polarisations, $\Delta\nu$ is the bandwidth under consideration (of a single
channel for spectral line observations, or the full bandwidth for continuum)
and the SEFD is defined as:

$SEFD=\frac{2kT_{sys}}{\eta_{A}A_{g}}$ (7)

where $k$ is the Boltzmann constant, $T_{sys}$ is the system temperature,
$\eta_{A}$ is the aperture efficiency, and $A_{g}$ is the geometric dish area.
Further breakdowns of the equations used in the sensitivity calculator can be
found in the online documentation that accompanies the code and web client.

####  C.1.1 Assumptions

With the telescope design still actively under construction, and the final
efficiencies and transmission coefficients unknown, a number of assumptions
about these parameters, as well as those relating to as yet to be defined
instrumentation suites.

  * •

Telescope efficiencies such as system, dish, forward, spillover, illumination,
polarisation and blocking were assumed based on parameters from other similar
facilities

  * •

Atmospheric and sky transmissions and temperatures were based on measurements
of conditions on the Chajnantor Plateau

  * •

No specific instrument efficiencies were used because they depend on evolving
underlying technology

  * •

Receiver temperatures were assumed to be near their fundamental quantum limit
at $T_{rx}=5h\nu/k$

####  C.1.2 Caveats

Because there are a number of unknowns with respect to the final operational
model, instrument suite, and calibration strategies of the observatory, the
following caveats apply to the sensitivity calculations:

  * •

Only on-source integration times are provided. This is because the operational
model, calibration strategies and mapping patterns/speeds are not yet known
for the observatory.

  * •

No pixel spacing information is inherent in the calculations because that is
highly dependent on how the future instrumentation will be implemented. For
the simulations, a pixel size consistent with diffraction limited seeing is
applied to the input image or spectrum.

###  C.2 Implementation

The code was implemented with a python backend. For the web client, a
javascript front end has been implemented. The source code is available
through github with the web front end can be accessed through the AtLAST
website. In both instances, best software engineering practices were employed,
with an emphasis on modularity, testing and documentation to allow for future
development of the calculator as the project evolves.

####  C.2.1 Source code

Python was chosen for backend development because of its prevalence within the
astronomical community: enabling interested astronomers to modify parameters
of the simulation with relative ease. The backend code is highly modular to
allow for as much code reuse as possible and to simplify updates. A
description of the public and REST APIs as well as UML class diagrams are
available as part of the documentation.

All parameters relevant to the sensitivity calculation have been abstracted so
that they can be modified by the user using dedicated input files, or
specified at runtime at the command line, or within a python instance. This
latter mode allows astronomers to provide the calculators with a set of
telescope properties which vary (i.e. telescope diameter ranging from 40 to 60
m) in a loop to generate tables of output sensitivities or time estimates.
This configurability is useful as the telescope properties (and eventually,
the instrument properties) are refined. Semantic validation is performed using

All of the code is subject to unit and functional testing to ensure that as
code is updated and modified, the underlying functionality changes in expected
ways (i.e. if increasing calculation efficiency, the output results do not
change). For maintainability, the code is linted and adheres to PEP8 and Black
formatting principles.

Extensive readthedocs documentation was created for 1) understanding the
individual code components 2) for installing the code in its own conda, venv
or poetry environment and 3) user guides for using the calculator and web
client.

####  C.2.2 Web Client

Javascript was chosen for the web client because of its ease of use for simple
projects and is consistent with other calculators being developed by other
facilities. As with the backend, best software engineering principles were
used, and a minimal set of parameters are passed between the front and back
ends of the project to minimise traffic and lag time. Client side validation
is performed before any properties are passed to the backend for processing.

###  C.3 Imaging and Spectroscopic simulator

Included in the source code repository are two Jupyter Notebooks which show
how the sensitivity calculator can be used to derive telescope simulations.
These notebooks take as input the set of observing properties required for the
calculation (the same ones as can be found in the web version of the
calculator) and an input FITS file in either two or three dimensions,
depending on the notebook being used. The notebooks then use the sensitivity
calculator to derive the noise levels to apply to the input image based on the
telescope parameters and on-source integration times specified. These noise
levels (as Gaussian noise) are then applied to a new version of the input
image which has been convolved with a Gaussian beam corresponding to the
diffraction limit of a 50m telescope at the input wavelength.

###  C.4 Related simulators

The Maria simulator (van Marrewijk et al., 2024) was created alongside the
sensitivity calculator described above to simulate the time variable
atmosphere from various sites around the world used for (sub-)mm observations.
It can be used to optimise observing strategies. Its performance has been
tested against the actual performance of the Mustang-2 instrument on the Green
Bank Telescope, and predictions have been made for what can be expected for
large scale source recovery with AtLAST. Maria creates timeseries data which
can then be transformed into sky maps using data reduction software. The
timestream data can be exported into FITS or HDF5 file formats to allow for
ingest into the relevant packages for the telescope being simulated.

##  Appendix D Acronyms

Acronym | Meaning
---|---
ACA | Atacama Compact Array (Morita Array, 7-metre component of ALMA)
ACT | Atacama Cosmology Telescope
AGN | Active Galactic Nucleus
ALMA | Atacama Large Millimeter/submillimeter Array
APEX | Atacama Pathfinder Experiment (APEX)
AtLAST | Atacama Large Aperture Submillimetre Telescope
ASTE | Atacama Submillimetre Telescope Experiment
BICEP | Background Imaging of Cosmic Extragalactic Polarization
CGM | Circumgalactic medium
CIB | Cosmic infrared background
CMB | Cosmic microwave background
CMB-S4 | The US Dept of Energy backed “Stage IV” proposed CMB experiment
CSO | Caltech Submm Observatory
EHT | Event Horizon Telescope
FoV | Field of View
FYST | Fred Young Submm Telescope (formerly CCAT)
GMC | Giant Molecular Cloud
GRB | Gamma Ray Burst
ICM | Intracluster Medium
IFU | Integral Field Unit
IRAM | Institut de Radioastronomie Millimétrique
ISM | Interstellar Medium
JCMT | James Clerk Maxwell Telescope
JWST | James Webb Space Telescope
KAO | Kuiper Airborne Observatory
KID | Kinetic Inductance Detector
LAS | Largest Angular Scale
LMT | Large Millimeter Telescope
LSST | The Vera C. Rubin Observatory, formerly “Large Synoptic Survey Telescope”
MOS | Multi-Object Spectragraph
MRS | Maximum Resolvable Scale
ngVLA | The next-generation Very Large Array
NOEMA | NOrthern Extended Millimeter Array
SKAO | Square Kilometer Array Observatory
SMA | Submillimeter Array
SO | Simons Observatory
TAO | (University of) Tokyo Atacama Observatory
WHIM | Warm/Hot Ionised Medium

##  Appendix E Author affiliations

1UK Astronomy Technology Centre, Royal Observatory Edinburgh, Blackford Hill,
Edinburgh EH9 3HJ, UK 2Institute of Theoretical Astrophysics, University of
Oslo, PO Box 1029, Blindern 0315, Oslo, Norway 3European Southern Observatory
(ESO), Karl-Schwarzschild-Strasse 2, Garching 85748, Germany 4Astrochemistry
Laboratory, Code 691, NASA Goddard Space Flight Center, Greenbelt, MD 20771,
USA 5Laboratoire Lagrange, Université Côte d’Azur, Observatoire de la Côte
d’Azur, CNRS, Blvd de l’Observatoire, CS 34229, 06304 Nice cedex 4, France
6Astronomy Unit, Department of Physics, University of Trieste, via Tiepolo 11,
Trieste 34131, Italy 7INAF – Osservatorio Astronomico di Trieste, via Tiepolo
11, Trieste 34131, Italy 8IFPU – Institute for Fundamental Physics of the
Universe, Via Beirut 2, 34014 Trieste, Italy 9NRC Herzberg Astronomy and
Astrophysics, 5071 West Saanich Rd, Victoria, BC, V9E 2E7, Canada 10Department
of Physics and Astronomy, University of Victoria, Victoria, BC, V8P 5C2,
Canada 11Cosmic Dawn Center (DAWN), Denmark 12DTU-Space, Technical University
of Denmark, Elektrovej 327, DK2800 Kgs. Lyngby, Denmark 13Max-Planck-Institut
für extraterrestrische Physik, Giessenbachstrasse 1 Garching, Bayern, D-85748,
Germany 14Purple Mountain Observatory, Chinese Academy of Sciences, 10 Yuanhua
Road, Nanjing 210008, China 15Department of Physics and Astronomy, University
of Pennsylvania, 209 South 33rd Street, Philadelphia, PA, 19104, USA
16Department of Physics and Astronomy, University College London, Gower
Street, London WC1E 6BT, UK 17Max-Planck-Institut für Radioastronomie (MPIfR),
Auf dem Hügel 69, D-53121 Bonn, Germany 18School of Physics & Astronomy,
Cardiff University, The Parade, Cardiff CF24 3AA, UK 19Division of Geological
and Planetary Sciences, California Institute of Technology, Pasadena, CA
91125, USA. 20Rosseland Centre for Solar Physics, Institute of Theoretical
Astrophysics, University of Oslo, Postboks 1029 Blindern, N-0315 Oslo, Norway
21Haystack Observatory, Massachusetts Institute of Technology, 99 Millstone
Road, Westford, MA 01886, USA 22National Astronomical Observatory of Japan,
2-21-1 Osawa, Mitaka, Tokyo 181-8588, Japan 23Black Hole Initiative, Center
for Astrophysics $\mid$ Harvard & Smithsonian, 20 Garden Street, Cambridge, MA
02138, USA 24INAF Osservatorio Astronomico di Brera, via Brera 28, 20121
Milano, Italy 25National Astronomical Observatory of Japan, Osawa 2-21-1,
Mitaka, Tokyo 181-8588, Japan 26Department of Space, Earth, & Environment,
Chalmers University of Technology, Chalmersplatsen 4 412 96 Gothenburg, Sweden
27Université de Strasbourg, CNRS, Observatoire astronomique de Strasbourg, UMR
7550, F-67000 Strasbourg, France 28Academia Sinica Institute of Astronomy and
Astrophysics, 645 N. A’ohoku Place, Hilo, HI 96720, USA 29Hvar Observatory,
Faculty of Geodesy, University of Zagreb, Zagreb, Croatia 30Academia Sinica
Institute of Astronomy and Astrophysics (ASIAA), No. 1, Sec. 4, Roosevelt
Road, Taipei 106216, Taiwan 31International Centre for Radio Astronomy
Research (ICRAR), The University of Western Australia, M468, 35 Stirling
Highway, Crawley, WA 6009, Australia 32Armagh Observatory and Planetarium,
College Hill, Armagh, BT61 9DB, UK 33INAF, Osservatorio di Astrofisica e
Scienza dello Spazio, via Piero Gobetti 93/3, 40129 Bologna, Italy 34INFN,
Sezione di Bologna, viale Berti Pichat 6/2, 40127 Bologna, Italy 35Instituto
de Astrofisica de Canarias (IAC), E-38205 La Laguna, Tenerife, Spain
36Universidad de La Laguna, Dpto. Astrofısica, E-38206 La Laguna, Tenerife,
Spain 37School of Physics, Trinity College Dublin, College Green, Dublin 2,
Ireland 38Université Paris-Saclay, Université Paris Cité, CEA, CNRS, AIM,
91191, Gif-sur-Yvette, France 39Departamento de Física de la Tierra y
Astrofísica e Instituto de Física de Partículas y del Cosmos (IPARCOS).
Universidad Complutense de Madrid, Av. Complutense, s/n, 28040 Madrid, Spain
40Department of Astronomy, Cornell University, Ithaca, NY 14853, USA;
Department of Physics, Cornell University, Ithaca, NY 14853, USA 41Planetary
Systems Laboratory, NASA Goddard Space Flight Center, Greenbelt, MD, 20771,
USA 42Departments of Astronomy and of Earth and Planetary Science, University
of California Berkeley, Berkeley, California, 94720, USA 43Victoria University
of Wellington, Wellington, New Zealand 44Physics and Astronomy, University of
Southampton, Highfield, Southampton, SO17 1BJ, UK 45DARK, Niels Bohr
Institute, University of Copenhagen, Jagtvej 155, Copenhagen N, 2200, Denmark
46Max Planck Institute for Extraterrestrial Physics, Garching bei München,
85748, Germany 47Jodrell Bank Centre for Astrophysics, University of
Manchester, Manchester, UK 48Leiden Observatory, Leiden University, P.O. Box
9513, 2300 RA Leiden, the Netherlands 49Faculty of Electrical Engineering,
Mathematics and Computer Science, Delft University of Technology, Mekelweg 4,
2628 CD Delft, The Netherlands 50SRON – Netherlands Institute for Space
Research, Niels Bohrweg 4, 2333 CA Leiden, The Netherlands 51School of
Mathematics, Statistics and Physics, Newcastle University, Newcastle upon
Tyne, NE1 7RU, UK 52Max Planck Institute for Astronomy, Koenigstuhl 17, 69117