∎

11institutetext: C. Evans22institutetext: European Space Agency (ESA), ESA
Office, Space Telescope Science Institute, 3700 San Martin Drive, Baltimore,
MD 21218, USA
UK Astronomy Technology Centre, Royal Observatory, Blackford Hill, Edinburgh,
EH9 3HJ, UK
22email<EMAIL_ADDRESS>
S. Cristiani, V. D’Odorico, P. Di Marcantonio, M. Franchini and A. Trost
33institutetext: INAF - Osservatorio Astronomico di Trieste, Via G. B. Tiepolo
11, 34143, Trieste, Italy
C. Opitom and C. Snodgrass44institutetext: Institute for Astronomy, University
of Edinburgh, Royal Observatory, Blackford Hill, Edinburgh, EH9 3HJ, UK
G. Cescutti55institutetext: Dipartimento di Fisica, Sezione di Astronomia,
Università di Trieste, Via G. B. Tiepolo 11, 34143 Trieste, Italy
J. M. Alcalá and M. Marconi 66institutetext: INAF-Osservatorio Astronomico di
Capodimonte, via Moiariello 16, 80131 Napoli, Italy
S. H. P. Alencar77institutetext: Departamento de Fisica - ICEx - UFMG, Av.
Antônio Carlos, 6627, 30270-901, Belo Horizonte, MG, Brazil
S. A. Balashev88institutetext: Ioffe Institute, Polyteknicheskaya 26, 194021
Saint-Petersburg, Russia
HSE University, Saint-Petersburg, Russia
B. Barbuy, M. Diaz, H. Ernandes and S. Rossi 99institutetext: Universidade de
São Paulo, IAG, Rua do Matão 1226, Cidade Universitária, São Paulo 05508-900,
Brazil
N. Bastian 1010institutetext: Donostia International Physics Center (DIPC),
Paseo Manuel de Lardizabal, 4, 20018, Donostia-San Sebastián, Guipuzkoa, Spain
IKERBASQUE, Basque Foundation for Science, E-48013 Bilbao, Spain
U. Battino 1111institutetext: E. A. Milne Centre for Astrophysics, Department
of Physics and Mathematics, University of Hull, Hull, HU6 7RX, UK
P. Cambianica, G. Cremonese and A. Grazian 1212institutetext: INAF-
Osservatorio Astronomico di Padova, Vicolo dell’Osservatorio 5, 35122, Padova,
Italy
R. Carini and P. Ventura 1313institutetext: INAF, Osservatorio Astronomico di
Roma, Via Frascati 33, 00078, Monte Porzio Catone (RM), Italy
B. Carter 1414institutetext: Centre for Astrophysics, University of Southern
Queensland, University Drive, Springfield Central, QLD 4300, Australia
S. Cassisi 1515institutetext: INAF, Osservatorio Astronomico di Abruzzo, Via
M. Maggini, I-64100 Teramo, Italy
INFN - Sezione di Pisa, Largo Pontecorvo 3, I-56127 Pisa, Italy
B. V. Castilho 1616institutetext: Laboratório Nacional de Astrofísica MCTI,
Rua Estados Unidos 154, Itajubá, MG 37504364, Brazil
N. Christlieb, A. Quirrenbach and J. Stürmer 1717institutetext:
Landessternwarte, Zentrum für Astronomie der Universität Heidelberg,
Königstuhl 12, 69117 Heidelberg, Germany
R. Cooke 1818institutetext: Centre for Extragalactic Astronomy, Durham
University, Durham DH1 3LE, UK
S. Covino and M. Genoni 1919institutetext: INAF - Osservatorio Astronomico di
Brera, via Bianchi 36, 23807, Merate (LC), Italy
K. Cunha 2020institutetext: Observatório Nacional, Rua Gen. José Cristino 77,
São Cristóvão, 20921-400, Rio de Janeiro, Brazil
Steward Observatory, University of Arizona, 950 N. Cherry Ave., Tucson, AZ,
85719
A. R. da Silva, R. E. Giribaldi, B. Pilecki and R. Smiljanic
2121institutetext: Nicolaus Copernicus Astronomical Center, Polish Academy of
Sciences, ul. Bartycka 18, 00-716, Warsaw, Poland
V. D’Elia 2222institutetext: Italian Space Agency - Space Science Data Centre,
via del Politecnico snc 00133 Rome, Italy
INAF - Osservatorio Astronomico di Roma, via di Frascati 33, 00040,
Monteporzio Catone, Rome, Italy
A. De Cia 2323institutetext: Department of Astronomy, University of Geneva,
Chemin Pegasi 51, 1290 Versoix, Switzerland
G. De Silva and T. Zafar 2424institutetext: Australian Astronomical Optics,
Macquarie University, 105 Delhi Road, North Ryde, NSW 2113, Australia
A. Fitzsimmons 2525institutetext: Astrophysics Research Centre, School of
Physics and Astronomy, Queen’s University Belfast, Belfast BT7 1NN, UK
B. Gänsicke 2626institutetext: University of Warwick, Department of Physics,
Gibbet Hill Road, Coventry, CV7 4AL, UK
C. J. Hansen 2727institutetext: Goethe University Frankfurt, Institute for
Applied Physics, Max-von-Laue-Str. 12, 60438 Frankfurt am Main, Germany
F. La Forgia and M. Lazzarin2828institutetext: Dipartimento di Fisica e
Astronomia-Padova University, Vicolo dell’Osservatorio 3, 35122 Padova, Italy
W. Marcolino 2929institutetext: Observatório do Valongo, Universidade Federal
do Rio de Janeiro, Ladeira Pedro Antônio, 43, Rio de Janeiro, Brazil
A. Migliorini 3030institutetext: INAF - Institute of Space Astrosphysics and
Planetology, via Fosso del Cavaliere 100, 00133, Roma, Italy
P. Noterdaeme 3131institutetext: Franco-Chilean Laboratory for Astronomy,
Camino El Observatorio 1515, Las Condes, Santiago, Chile
3 Institut d’Astrophysique de Paris, CNRS-SU, UMR 7095, 98bis bd Arago, 75014
Paris, France
C. B. Pereira 3232institutetext: Observatório Nacional, Rua Gen. José Cristino
77, São Cristóvão, 20921-400, Rio de Janeiro, Brazil
S. Randich3333institutetext: INAF - Osservatorio Astrofisico di Arcetri, Largo
E. Fermi 5, 50125, Firenze, Italy
E. Vanzella3434institutetext: INAF - OAS, Osservatorio di Astrofisica e
Scienza dello Spazio di Bologna, Via Gobetti 93/3, 40129, Bologna, Italy
D. Wright 3535institutetext: Centre for Astrophysics, University of Southern
Queensland, West Street, Toowoomba, QLD 4350, Australia

# The CUBES Science Case

Chris Evans  Stefano Cristiani  Cyrielle Opitom  Gabriele Cescutti  Valentina
D’Odorico  Juan Manuel Alcalá  Silvia H. P. Alencar  Sergei Balashev  Beatriz
Barbuy  Nate Bastian  Umberto Battino  Pamela Cambianica  Roberta Carini  Brad
Carter  Santi Cassisi  Bruno Vaz Castilho  Norbert Christlieb  Ryan Cooke
Stefano Covino  Gabriele Cremonese  Katia Cunha  André R. da Silva  Valerio
D’Elia  Annalisa De Cia  Gayandhi De Silva  Marcos Diaz  Paolo Di Marcantonio
Heitor Ernandes  Alan Fitzsimmons  Mariagrazia Franchini  Boris T. Gänsicke
Matteo Genoni  Riano E. Giribaldi  Andrea Grazian  Camilla Juul Hansen
Fiorangela La Forgia  Monica Lazzarin  Wagner Marcolino  Marcella Marconi
Alessandra Migliorini  Pasquier Noterdaeme  Claudio Pereira  Bogumil Pilecki
Andreas Quirrenbach  Sofia Randich  Silvia Rossi  Rodolfo Smiljanic  Colin
Snodgrass  Julian Stürmer  Andrea Trost  Eros Vanzella  Paolo Ventura  Duncan
Wright  Tayyaba Zafar

(Received: 20 April 2022 / Accepted: 09 July 2022)

###### Abstract

We introduce the scientific motivations for the development of the Cassegrain
U-Band Efficient Spectrograph (CUBES) that is now in construction for the Very
Large Telescope. The assembled cases span a broad range of contemporary topics
across Solar System, Galactic and extragalactic astronomy, where observations
are limited by the performance of current ground-based spectrographs
shortwards of 400 nm. A brief background to each case is presented and
specific technical requirements on the instrument design that flow-down from
each case are identified. These were used as inputs to the CUBES design, that
will provide a factor of ten gain in efficiency for astronomical spectroscopy
over 300-405 nm, at resolving powers of $R$ $\sim$ 24,000 and $\sim$7,000. We
include performance estimates that demonstrate the ability of CUBES to observe
sources that are up to three magnitudes fainter than currently possible at
ground-ultraviolet wavelengths, and we place its predicted performance in the
context of existing facillities.

###### Keywords:

instrumentation: spectrographs – comets – stars – galaxies – transients

##  1 Introduction

The four 8.2m telescopes of the Very Large Telescope (VLT) at the European
Southern Observatory (ESO) are the world’s most scientifically productive
ground-based observatory in the visible and infrared. Looking to the future of
the VLT there is a long-standing aspiration for an optimised ultraviolet (UV)
spectrograph (e.g. bar14 ). Here we introduce the scientific motivations for
the Cassegrain U-Band Efficient Spectrograph (CUBES), that has been designed
to provide high-efficiency observations in the near UV (300-405 nm) at
spectral resolving powers of $R$ $\geq$ 20,000 and $R$ $\sim$ 7,000.

CUBES will provide the VLT with a unique capability and will open-up new
discovery space for a broad range of astrophysics. The near UV contains a
tremendous diversity of iron-peak and heavy elements in stellar spectra, as
well as some lighter elements (notably beryllium) and light-element molecules
(CO, CN, OH). The near-UV is also important for extragalactic science, such as
studies of the interstellar and circumgalactic medium (CGM) of distant
galaxies, in measuring the contribution of different types of galaxies to the
cosmic UV background, and follow-up of explosive transients. Closer to home,
significantly improved sensitivity at the shortest (ground-based) wavelengths
will also provide exciting new opportunities in solar system science.

The CUBES project completed a Phase A conceptual design in June 2021 and has
now entered the detailed design and construction phase; first science
operations are planned for 2028. Alongside an overview of the CUBES design in
this Special Issue ExA_Zanutta , this article introduces the science cases
developed during Phase A that motivated the instrument design and informed its
top-level requirements (TLRs). The cases assembled by the Science Team were
grouped into four fields: solar system, Galactic, extragalactic and
transients. Table 1 summarises the topics within each field, and highlights
those that are presented in more depth elsewhere in this Special Issue. While
the design was informed by the cases outlined here, we envisage CUBES more
generally as providing the VLT with an exciting new capability that will
enable an even broader range of applications and serve a diverse user
community once in operation.

In Sections 2, 3, 4 and 5 we briefly introduce the cases considered during
Phase A and some of their key instrument requirements. In Section 6 we
identify the science cases used to specify the TLRs, and in Section 7 we
compare the predicted performance of CUBES with other facilities worldwide.

Table 1: Summary of high-level science cases developed during the CUBES Phase A study. Entries in the third column refer to the expanded articles presented elsewhere in this Special Issue (SI). Field | Science Case | SI Contrib.
---|---|---
Solar System | S1: Cometary Science | ExA_Opitom
| S2: Icy Satellites | …
Galactic | G1: Accretion, winds & outflows in YSOs | ExA_Alcala
| G2: Exo-planet composition | …
| G3: Stellar astrophysics & exoplanets | …
| G4: Beryllium in metal-poor stars and stellar clusters |  ExA_Smiljanic , ExA_Giribaldi
| G5: Lithium production in novae | …
| G6: Metal-poor stars & light elements |  ExA_Bonifacio , ExA_Hansen
| G7: Neutron-capture elements | ExA_Ernandes
| G8: Precise metallicities of metal-poor pulsators | …
| G9: Horizontal branch stars in Galactic GCs | …
| G10: Early-type companions in binary Cepheids | …
| G11: Extragalactic massive stars | ExA_Evans
Extragalactic | E1: Primordial deuterium abundance | …
| E2: Missing baryonic mass in the high-$z$ CGM | ExA_DOdorico
| E3: Cold gas at high redshift | ExA_Balashev
| E4: Reionisation | …
Transients | T1: GRBs | …
| T2: Kilonovae | …
| T3: Superluminous supernovae | …

##  2 Solar System Cases

###  2.1 S1: Cometary science

The search for water in our solar system is far from complete and is
tremendously difficult with ground-based facilities given the water content of
Earth’s atmosphere. Infrared spectroscopy can detect ice on the surface of
distant bodies, but only for the largest/brightest objects. For smaller bodies
the ice is sub-surface and we must look for outgassing water escaping into
space, e.g. cometary comae and tails. The most powerful probe of this
outgassing water from the ground is OH emission at 308 nm (Fig. 2). While
observations in the far-UV or IR from space could be more sensitive, they
would also be substantially more expensive, and ground-based OH observations
are the most compelling next step (see snodgrass17a ); CUBES will provide this
capability.

Observation of the OH line is only possible with current facilities for a few
active comets while they are near the Sun and Earth. This severely limits
studies of water production in comets around their orbits and we miss seasonal
effects that the Rosetta mission revealed to be important (e.g. kramer17 ). It
also prevents study of the large majority of comets which are simply too
faint. Even more tantalising are studies of main-belt comets – bodies in
asteroidal orbits that undergo activity (usually detected by a dust tail or
trail) that is thought to arise from sublimation (e.g. recurrent activity near
perihelion). As shown by the X-Shooter spectrum in Fig. 2, constraining the OH
emission of such objects is well beyond our current capabilities. Main-belt
comets have typical sizes that are very common in the asteroid belt, so
detection of outgassing water would point to a potentially large population of
icy bodies, hence a large reservoir of water, of considerable interest in the
context of models of the formation and evolution of the inner solar system
(e.g. obrien18 ).

Enhanced sensitivity in the near UV will also allow us to measure the
deuterium to hydrogen ratio, using the same OH band, in a much wider range of
comets than currently possible, providing invaluable constraints to
investigate the fate of water in the solar system and the origin of
terrestrial water. Finally, CUBES will enable us to detect and measure the
abundance of critical species in the comae of comets, and constrain their
composition; with its greater sensitivity than current instrumentation it will
also give us access to more distant comets. This capability will be
particularly timely in the late 2020s and early 2030s, as it will enable
characterisation of potential targets for the ESA Comet Interceptor mission,
which is due to launch in 2029 snodgrass19 .

Figure 1: OH emission at 308 nm is a potentially powerful probe of water
outgassing from main-belt comets. (Image credit: ESO.)

Figure 2: VLT/X-Shooter spectrum of main-belt comet P/2012 T1 (black) compared
with an active Jupiter family comet (9P/Tempel 1, in blue). Significantly
better sensitivity is required for the OH line for such objects snodgrass17b .

Quantitative simulations from the end-to-end simulator have demonstrated that
CUBES observations will be able to detect outgassing water down to a level ten
times lower than the current upper limit possible with X-Shooter. For details
on these calculations and studies of the chemistry of cometary comae (D/H
ratio, N2/CO ratio, and detection of metals), we refer the reader to
ExA_Opitom  in this Special Issue.

####  2.1.1 S1: Key requirements

  * •

Spectral resolution: Essential: $R$ $>$ 20,000 (with a goal of $R$ $\sim$
40,000). For OH detections in the faintest targets, $R$ $>$ 5,000 will suffice
(and provide even greater sensitivity), see ExA_Opitom .

  * •

Spectral coverage: Essential: 305-400 nm; Goal: 300-420 nm. Ensures well
sampled continuum on both sides of the OH(0,0) and CO+(3,0) bands.

  * •

Simultaneous coverage at longer $\lambda$: Desirable with UVES to place
constraints on the general composition of the coma. Note that observations
close in time (rather than simultaneous) would be almost as valuable.

  * •

Flux calibration: Required to $\pm$10% accuracy (goal of $\pm$5%).

  * •

Exposure times: Long exposures (of up to 2 hr) will be needed to detect faint
signatures.

  * •

Position angle: It is important to be able to configure the slit position
angle to sample different parts of the coma, e.g. jets or tails for active
comets.

###  2.2 S2: Icy satellites

####  2.2.1 Surface composition

The near-UV offers a unique opportunity to investigate the surface properties
of the satellites of our solar system, including the Galilean and Saturnian
satellites, Triton and several Kuiper Belt Objects (KBOs).

Sulphur-bearing materials are quite common on the surface of Europa, the third
Galilean satellite. Europa is one of the targets of the ESA JUICE mission
(together with Ganymede and Callisto), and the target of the NASA Europa
Clipper mission. Sulphur is expected to be found on Europa’s surface from Io’s
plasma torus lane81  and its presence was inferred in the past from ground-
based observations in the near-UV and visible jm70 ; wamsteker72 ; mcfadden80
; nelson87 ; noll95 ; spencer95 . A spatially-varying surface distribution of
SO2 on Europa was also reported, with a clear absorption band at 280 nm on the
trailing face, and less intense (or maybe absent) absorption on the leading
face lane81 . This absorption band has not been formally identified as
condensed SO2 but the detection of polymeric sulphur around 360 nm with CUBES,
due to its improved sensitivity in the near-UV, would provide confirmation.

Other peculiar features that have been identified in the UV spectral range
could be ascribed to organic compounds, as reported for the South Pole of
Callisto with the Galileo/UVS spectrograph hendrix08 . The red slope observed
in the UV-visible in the spectra of the icy satellites of Saturn is not
typical of a water-dominated surface, and hence it is an open question which
species or process is responsible. If subject to radiolytic/photolytic or
thermal processing, hydrocarbons in the outer Solar System develop a very
distinctive absorption feature at about 220 nm, as a consequence of H loss
hendrix16 . They also present a distinctive spectral slope longwards of 300 nm
that can be used with CUBES to better constrain the surface composition of
these objects.

The Galilean satellites have exospheres mcgrath04 , or thin atmospheres, which
are thought to be the result of ion-induced sputtering and sublimation of the
surface materials. Any further information we can derive from observations of
these exospheres will provide insights into the processes acting on the
surface, as well as important constraints on the surface compositions. The
search for exospheric compounds is a key scientific goal in the near future
for satellites of the outer planets, with particular attention on the Galilean
moons. Greater sensitivity to features from Al (396.1 nm) and Fe (371.9, 374.5
and 395.9 nm), similar to those identified in Mercury’s exosphere bk17 , will
provide the opportunity to detect these elements for the first time around a
Galilean moon.

UV observations of the Galilean and Saturnian satellites have revealed a lot
about their surface compositions and exospheric properties. Unfortunately, the
JUICE mission will not cover the 300-370 nm range with its selected payload. A
similar UV spectrometer to that on JUICE will be part of the Europa Clipper
mission, with the aim of investigating the atmospheric and surface properties
of Europa. Future observations at 300-400 nm with a facility such as CUBES
will therefore be essential to complement measurements by missions to the
Jovian system.

####  2.2.2 S2: Key requirements

  * •

Spectral resolution: Essential: $R$ $>$ 10,000; Goal: $R$ $>$ 20,000.

  * •

Spectral coverage: Essential: 300-400 nm; Goal: 300-420 nm.

  * •

Simultaneous coverage at longer $\lambda$: Desirable with UVES to enable
identification of other species.

##  3 Galactic Cases

###  3.1 G1: Accretion, winds & outflows in YSOs

Studies of protoplanetary disks play a pivotal role in determining the initial
conditions for planet formation, and many of their characteristics are only
now being unveiled by the new observational facilities at high angular and
spectral resolution, over a wide range of wavelengths, from X-ray to radio.
The way in which circumstellar disks evolve and form protoplanets is strongly
influenced by the processes of mass accretion onto the star, ejection of
outflows and photo-evaporation in winds of the disk material. A proper
understanding of the impact of these phenomena requires comprehensive, multi-
wavelength studies of different physical processes throughout the first 10 Myr
of the star–disk evolution. Classical T Tauri stars (CTTS) are key objects to
investigate these processes observationally. They are young (a few Myr), low-
to solar-mass stars that are actively accreting mass from planet-forming
disks.

Spectroscopic surveys of CTTS in nearby (d $<$ 500 pc) star-forming regions,
have been used to study the mutual relationships between accretion, jets and
disk structure (e.g. giannini15 ; alcala17 ; manara17 ; nisini18 ) but further
aspects remained uncertain, mainly because of the low sensitivity and spectral
resolution so far available in the UV. A more expansive discussion of how
CUBES will enable important contributions to the investigation of
accretion/winds-outflows in young stellar objects (YSOs) is given by
ExA_Alcala  in this Special Issue. Here we briefly highlight some of the key
points:

  * •

High throughput observations at R $\sim$ 20,000 will enable more detailed
study of the accretion process than currently possible, via precise modelling
of the Balmer jump, and by investigating the high-order Balmer lines (e.g.
Fig. 3). The latter and Ca II H&K lines provide diagnostics of the accretion
funnel flows and heated chromosphere in the post-accretion shock region (e.g.
alencar12 ).

  * •

Studies of YSOs with low accretion rates will help to disentangle the
contribution from chromospheric emission, which is difficult to remove solely
using e.g. X-Shooter spectra, particularly for YSOs with spectral types
earlier than K3. Accretion has been confirmed in a few such objects using HST
data (e.g. alcala19 ), and we anticipate strong synergies with the HST ULLYSES
Director’s Discretionary Program, that has observed $\sim$60 CTTS (see
ExA_Alcala  for details).

  * •

Studies of CTTS in star-forming regions at larger distances and with sub-solar
metallicity and/or in different environments to current Galactic targets will
only be possible with much greater sensitivity in the near UV. This will
enable studies of accretion at low metallicities, as well as the effects of
local UV fields on accretion and winds – weak fields should not influence the
disks, whereas strong fields may modify the ionisation rate in the disk with
an important impact on mass-loss rates.

  * •

Emission lines from [O II] 372 nm and [S II] 407 nm are useful diagnostics of
the winds of protoplanetary disks. Also, by combining observations of [O II]
372 nm with [O I] 630 nm (from e.g. UVES), we will be able to measure the
degree of ionisation of the wind. The [O I] 630 nm line has provided the
historical foundation for identifying and measuring jets and disk winds, but
the accuracy of mass-loss rates improves significantly when paired with high-
resolution spectra of the [S II] 406.8 nm line fang18 . Much of the remaining
uncertainty relates to how jets and winds are excited, and the little-studied
[Ne III] 386.9 nm line has the potential to discriminate between shocks and
X-ray heating liu14 .

  * •

There are bright [Fe II] and [Fe III] lines in the near UV with higher
Einstein coefficients than the [Fe II] 1.25 and 1.64 $\mu$m lines that are
commonly observed in YSO jets. The [Fe III] lines are particularly important
as they are produced from high energy levels so are good tracers of the
temperature of the ionised gas. Observations of [Fe III] combined with [Fe II]
allow simultaneous determination of temperature, density and degree of
ionisation, as well as a precise estimate of the Fe abundance in the gas.

  * •

Beryllium abundances for older stars will help to constrain evolutionary
models of pre-main sequence stars, complementing the Li (670.7 nm) line.

####  3.1.1 G1: Key requirements

  * •

Spectral resolution: $R$ $>$ 20,000 for detailed quantitative analysis; $R$
$\gtrsim$ 5,000 for observations of the faintest targets (see ExA_Alcala  for
further discussion).

  * •

Spectral coverage: Essential: 310-400 nm; Goal: 310-430 nm (includes [S II]
407 nm and Ca II 422 nm).

  * •

Position angle: Non-parallactic observations are essential so the slit can be
aligned along the outflow/jet direction and/or perpendicular to it.

  * •

Simultaneous observations at longer $\lambda$: Given the variable nature of
YSOs, simultaneous observations with UVES are essential to correctly
characterise the CUBES observations.

  * •

Flux calibration: Required to better than $\pm$10% (goal), $\pm$15%
(essential).

Figure 3: Example X-Shooter spectrum (red) of the classical T Tauri star GQ
Lup in the region of the Balmer jump. Note the blending of the high Balmer
series due to the spectral resolving power of only $R$ $\sim$ 5000, which
shifts the Balmer limit from 364.6 nm to an apparent jump at $\sim$370 nm. The
template spectrum of a non-accreting YSO of the same spectral type as GQ Lup
is overplotted (green). The continuum emission from a hydrogen slab is shown
by the black continuous line. The best fit to the data with the emission
predicted from the slab model is given by the blue line. (Adapted from
alcala17 .)

###  3.2 G2: Composition of exo-planets

In the past 25 years we have gone from not knowing if the Solar System is
unique to understanding that it is totally normal for stars to have planets.
Planetary systems have been identified in increasingly complex architectures,
including rich multi-planet systems around stars both more and less massive
than the Sun, circumbinary planets, and wide binaries with planets orbiting
one or both stellar components. Characterising exo-planets and understanding
their formation and evolution are now major research areas, but the
fundamental question of what these other worlds are made of cannot be answered
from studies of planets orbiting main-sequence hosts. We can only estimate the
bulk densities of transiting planets (with significant uncertainties), and
considerations of internal structure and composition are model dependent (e.g.
rs10 ).

The key to answering the above question draws on a method well-established in
planetary science: measuring the abundances of the building blocks of planets,
and their fragments. We know more about the composition of the terrestrial
planets in the Solar System from abundance studies of meteorites than from
surface samples collected by astronauts and robots. Similarly, it has been
demonstrated (z07 ) that we can accurately measure the bulk composition of
exo-planetary systems (analogous to studies of Solar System meteorites) from
spectroscopic analysis of white dwarfs which are accreting debris from
tidally-disrupted planetesimals jura03 .

White dwarfs are moderately hot stars so most of the atomic transitions that
can be used for abundance studies are located in the ground-based UV, in
particular for refractory and transitional elements (Sc, Ti, V, Cr, Mn, Fe,
Ni). White dwarfs, being Earth-sized, are intrinsically faint, so high-
throughput high-resolution spectroscopy ($R\geq 20,000$) is essential. So far
only about twenty systems have precise abundances for more than five chemical
elements, the bulk of which were observed with Keck-HIRES (e.g. klein10 ).

The current workhorses for studies of the bulk compositions of exo-planets are
VLT-UVES (south) and Keck-HIRES (north), but they are limited to targets with
$V$ $<$ 16 to 16.5 mag for such observations. High S/N data are possible with
X-Shooter (see Fig. 4) but we then lack the spectral resolution required to
study the detailed chemistry. CUBES will dramatically advance this new area of
exo-planet research, pushing $\simeq$3 mag. fainter than Keck-HIRES or VLT-
UVES and opening up a vastly larger sample of targets. For instance, $\simeq
260,000$ white dwarfs were identified from Gaia DR2, almost half of which were
new discoveries in the southern hemisphere gf19 . Low-resolution spectroscopic
confirmation will be done with SDSS-V and later with 4MOST. Statistically,
$\simeq$1% of white dwarfs are suited to detailed abundance studies, so CUBES
will be able to increase the sample of known exo-planetesimal compositions to
a level comparable to the number of meteorite samples with known abundances,
thus providing exquisitely detailed abundance information for the next
generation of planet formation models.

####  3.2.1 G2: Key requirements

  * •

Spectral resolution: $R$ $>$ 20,000.

  * •

Spectral coverage: Essential: 305-400 nm; Goal: 305-420 nm.

  * •

Position angle: Parallactic observations are sufficient (unless avoiding a
nearby bright star).

  * •

Simultaneous observations at longer $\lambda$: Desirable.

Figure 4: Example 1 hr X-Shooter spectrum for a white dwarf with $V$ $=$ 16.5
mag.

###  3.3 G3: Stellar astrophysics & exoplanets

Stellar astrophysics faces significant questions as to how stellar magnetic
fields are generated by an internal dynamo, emerge as flux tubes, and energise
the outer atmospheres and activity, including impulsive phenomena such as
flares tsd15 . Even for the Sun, our theory of atmospheric energy transport
and release is incomplete, yet Solar physics necessarily remains the key to
our explanations of the physics of stellar atmospheres and activity.
Conversely, stars provide proxies for studying otherwise unobservable Solar
evolution, and testing Solar models under more extreme conditions.

Two questions arise: 1) Can stellar activity be better understood by new
connections with Solar physics? 2) How can this new understanding of stellar
activity advance exoplanet studies?

Exoplanet surveys of young or otherwise active stars are important to
investigate the shared evolution of host stars and their planetary systems.
However, the interpretation of high-resolution visible spectroscopy (with e.g.
ESPRESSO) can be limited by our understanding of stellar activity. To address
this we need a better understanding of stellar chromospheres: the region where
stellar activity is prominent due to the dominance of non-radiatively heated
plasma, with variability from flares and other transient magnetic activity. To
characterise chromospheric activity we require near-UV observations in support
of e.g. ESPRESSO programmes, and with better sensitivity than is available
from current instruments (particularly given the relatively rapid timescales
of variability).

The ideal programme would use contemporaneous CUBES and ESPRESSO observations
to:

  * •

Monitor the Ca II H&K lines that are diagnostics of quiescent chromospheres
and transient flares klocova17 .

  * •

Include the Balmer jump as an added diagnostic of stellar chromosphere
physics, for both the quiescent chromosphere and any major flares that are
expected to change the UV continuum kowalski15 .

  * •

Accumulate chromospheric spectra to expand our knowledge of activity trends
for Solar-type stars. This would aid forecasting of future Solar activity
level extremes and their impact on space weather in the Solar System, while
also giving insights into space weather in exoplanetary systems to investigate
its impact on habitable zones.

  * •

Monitor radial velocities (ESPRESSO) and chromospheric activity (CUBES)
simultaneously to help separate genuine stellar reflex motion due to
exoplanets from jitter arising from stellar activity.

In short, CUBES and ESPRESSO together would provide a powerful combination for
stellar and exoplanet science, including an enhanced search for planetary
systems around the youngest, most active stars. Future observations would
focus on young, active and rapidly-rotating stars, spanning cool G to M dwarfs
over a wide range of $V$ $=$ 8-16 mag. Alongside radial-velocity studies with
ESPRESSO, CUBES would be used to monitor the Ca II H&K lines (and other near-
UV chromospheric variability), as well as monitoring flares to constrain
models of flare physics (e.g. k13 ).

Such a programme would be similar to ESPRESSO observations of inactive stars
studied in current exoplanet surveys, but would need to include enough
exposures to provide phase coverage of the stellar rotation rate. As a guide,
a minimum set of ESPRESSO data would comprise two to four exposures per night
of a given target star (reaching S/N $=$ 10 per pixel), for two to four near-
contiguous nights. This would enable tomographic active-region modelling to be
applied, with CUBES observations obtained roughly contemporaneously to
disentangle activity effects. If flare monitoring was also included, the CUBES
exposure times are best kept as short as possible (1-10 min). For example, a
single young active high-priority target could be continuously monitored with
ESPRESSO across a rotation period of several nights to improve the detection
of planets around the youngest, most active stars (e.g. yu17 ). Adding
contemporaneous CUBES monitoring of chromospheric variability would provide
firmer constraints on stellar activity, hence a clearer understanding of the
radial-velocity changes due to a planet.

Tomographic mapping of chromospheres requires observations at $R$ $\geq$
20,000 of the Ca II H&K and Balmer lines (e.g. donati97 ). Both the CUBES and
ESPRESSO data could also be used for tomographic imaging of different layers
of the active regions of the host star, which further helps to separate
stellar activity from planetary signals in the spectra (as used already to
e.g. detect hot Jupiters around stars as young as T-Tauri stars donati16 ),
and is of intrinsic scientific interest to those studying stellar magnetic
activity and dynamos.

####  3.3.1 G3: Key requirements

  * •

Spectral resolution: Essential: $R$ $>$ 20,000; Desirable: $R$ of up to
100,000.

  * •

Spectral coverage: Essential: 305-400 nm; Goal: 300-420 nm.

  * •

Simultaneous coverage at longer $\lambda$: Desirable to have contemporaneous
observations with ESPRESSO (using another UT).

###  3.4 G4: Beryllium in metal-poor stars and stellar clusters

Although one of the lightest, simplest elements, there remain profound
questions regarding the production of Be in the early Universe. For instance,
the upper limit on its abundance of log(Be/H) $<$ $-$14.0 in an extremely
metal-poor star with [Fe/H] $=$ $–$3.84 is consistent with no primordial
production spite19 , but larger samples of very metal-poor stars are required
to constrain its formation channels. Beryllium is also a powerful tracer of
early nucleosynthesis in a range of different contexts, including stellar
evolution, the formation of globular clusters, and the star-formation rate and
chemical evolution of the Galaxy smiljanic14 .

Furthermore, 9Be is thought to be mostly produced by spallation of Galactic
cosmic rays reeves70 . During the early stages ($<$0.5-1 Gyr) of Galactic
evolution, cosmic rays were generated and transported on a Galactic scale,
diffusing over the whole Galaxy. This implies that the production site was
widespread and, to a first approximation, the Be abundance should have one
value at a given time over the whole Galaxy beers00 ; sy01 . Beryllium could
then be used as an ideal cosmic clock pasquini05 ; pasquini07 ; smiljanic09
and is therefore a key element to trace star-formation histories. In
particular, its abundance can potentially be used to identify accretion of
external systems into the early halo, complementing astrometry from Gaia
molaro20 ; smiljanic21 .

Beryllium is destroyed in stars at the relatively low temperature of 3.5
$\times$ 106 K (compared to 2.5 $\times$ 106 K for lithium). Measurements of
both Li and Be in the same star therefore provide valuable probes of stellar
interiors and, in particular, of the non-standard mechanisms at work. Lithium
estimates are now available from recent large spectroscopic surveys (e.g. in
several open clusters) but Be has only been studied for the brighter stars in
a few clusters (down to V $\sim$ 14) smiljanic10 ; smiljanic11 ; boesgaard20 ;
boesgaard22 . CUBES will provide Be abundances for many more stars, in a range
of clusters, complementing the Li measurements, and enabling study of internal
mixing as a function of stellar age, mass, and metallicity.

The only Be lines available from the ground for abundance estimates are the Be
II 313.04 and 313.10 nm resonance lines, which require good S/N ($>$50) and
sufficient resolution ($R$ $\gtrsim$ 20,000) to clearly discern them from
nearby, relatively strong V (313.03 nm) and Ti II (313.08 nm) absorption.

There are only $\sim$200 stars with estimated Be abundances (from Keck-HIRES
and VLT-UVES), which span near-solar metallicities down to [Fe/H] $<$ $–$3
smiljanic09 ; boesgaard99 ; primas00a ; primas00b ; boesgaard11 . The limiting
magnitude of current observations is V $\sim$ 12 mag., for observations of a
few hours. An ambitious UVES programme with a total integration of 20 hrs has
observed a metal-poor star ([Fe/H] $=$ $-$3.0) with $V$ $=$ 13.2 mag.,
resulting in a S/N around the Be doublet of $\sim$200 smiljanic21 , but such
long exposures are impractical to assemble large samples of objects.

In general, CUBES will be able to provide estimates of Be in stars belonging
to different populations for the first time, complementing results for Li,
metallicities, and $\alpha$-elements from surveys such as Gaia-ESO and 4MOST,
and Gaia astrometry. Some example target magnitudes in this context include:

  * •

Essential: $V$ $=$ 16 mag.: Turn-off and main-sequence stars (FGK-type) in
open clusters where Li abundances are available from other surveys.

  * •

Essential: $G$ $=$ 16 mag.: Turn-off stars in the Gaia-Enceladus population
and/or other kinematic structures discovered by Gaia up to 2.5 kpc from the
Sun.

  * •

Goal: $V$ $=$ 16.7 mag.: Solar-type stars in one of the oldest known globular
clusters.

  * •

Goal: $G$ $=$ 17.5 mag.: Turn-off stars up to 5 kpc from the Sun.

We now briefly expand on two key areas for studies of Be abundances.

####  3.4.1 Beryllium abundances in extremely metal-poor stars

Abundances of beryllium in stars with [Fe/H] $<$ $–$3.0 seem to indicate an
increased scatter or even a flattening of the relation between [Fe/H] and
log(Be/H) (e.g. primas00b ). Possible explanations include: Be production by
inhomogeneous primordial nucleosynthesis, significant production of Be by pre-
Galactic cosmic-rays in the intergalactic medium, or the inhomogeneous nature
of the early halo having formed from mergers of stellar systems with different
[Fe/H] vs. log(Be/H) relations smiljanic21 .

Ongoing multi-band imaging surveys are finding excellent candidates for very
metal-poor stars (e.g. Pristine pristine ; SkyMapper skymapper ), as well as
wide-field spectroscopic surveys (e.g. LAMOST and the upcoming WEAVE and 4MOST
surveys). However, stars with such low metallicity are faint and hard to
observe in the near-UV with current spectrographs (see ExA_Bonifacio  and
ExA_Hansen  in this Special Issue). Only 10-15 stars have been investigated in
this region and with only upper limits to the Be abundances in several cases.
Improved near-UV sensitivity is required to access a larger sample over the
metallicity range $-$4.0 $<$ [Fe/H] $<$ $-$3.0 to investigate if there is a
plateau or a larger scatter in the [Fe/H] vs. log(Be/H) relation.

Even with enhanced near-UV sensitivity, it will be particularly challenging to
detect the Be lines in such extremely metal-poor stars, as the lines become so
weak. Detailed simulations, presented elsewhere in this Special Issue
ExA_Smiljanic , show that this will be possible with CUBES for stars that have
$\log$(Be/H) $\gtrsim$ $-$13.6, as long as a S/N of order 400 can be achieved.
Given the anticipated high efficiency of CUBES, stars down to V $\sim$ 14 mag.
will be able to be observed within a couple of nights of observations
ExA_Smiljanic , and a dedicated Large Programme could more than double the
precious sample of extremely metal-poor unevolved stars to be investigated.

####  3.4.2 Multiple populations in globular clusters

Globular clusters (GCs) are not the simple stellar populations that we once
thought there were. They exhibit star-to-star variations in specific elements
(e.g., He, C, N, O, Na, Mg, Al) that bear the hallmark of high-temperature
H-burning gratton12 ; variations are also seen in fluorine abundances (e.g.
ymc08 ; gcs19 ). These abundance variations can be observed spectroscopically
and also photometrically, with the appropriate choice of filters, due to the
changing of spectral features within the band pass. This phenomenon is
observed in nearly all of the ancient GCs and has recently been found in many
younger clusters as well.

Many scenarios have been suggested to explain this phenomenon, with most
invoking multiple epochs of star formation within the cluster, in which a
second generation of stars are formed from material contaminated by proton-
capture nucleosynthesis products. Proton-capture reactions can decrease the
abundances of Li, C, O and Mg, and enhance the abundances of N, Na and Al. The
observed correlations and anti-correlations of these light-element abundances
are thought to arise from mixing pristine and processed material, with
potential polluters thought to be first-generation massive asymptotic-giant-
branch stars (AGBs) or first-generation rapidly-rotating massive stars.
However, the proposed scenarios all fail to reproduce key observations.

An important constraint on the problem has come from lithium, that is easily
destroyed at high temperatures (like those required to burn H through the CNO
cycle). However, Li is a complicated element with potential nucleosynthetic
production happening in evolved stars through the Cameron-Fowler effect.
Beryllium, on the other hand, could provide new insights as it is not produced
in significant quantities through nuclear burning and only destroyed inside
stars. Thus, the polluting material from the first generation of evolved
stars, regardless of the exact nature of the polluting star, should be
completely devoid of Be. A star formed with only pristine material should
therefore have the original Be abundance of that material. Stars with
different amounts of polluted material should have diluted the surface Be
abundance to different levels. Stars with bigger fractions of polluted
material should have stronger correlations and anti-correlations between the
light elements and should be strongly depleted in Be.

The objective is to look for differences in the Be abundances between turn-off
stars of the same cluster. Different stars should have been polluted to
different levels, so if different Be abundances are detected we can use them
to quantify the fraction of polluting material. An expanded study of this case
is presented elsewhere in this Special Issue ExA_Giribaldi , where simulated
observations of 4 hr per star suggest that depletions of Be of $\sim$0.6 dex
will be possible in nearby GCs with turn-off magnitudes as faint as $V$ $=$ 18
mag.

Furthermore, it is not clear whether the multiple populations phenomenon is
limited to light-elements, as some studies have suggested that K, as well as
neutron-capture elements (e.g. Eu, La) may also vary. The diverse set of
elements made available through CUBES will offer the exciting potential of
discovering new variations and correlations that may finally allow us to solve
this long-standing puzzle.

####  3.4.3 Key requirements

  * •

Spectral resolution: $\geq$ 20,000 to sufficiently resolve the Be II doublet
(see e.g. smiljanic14 ).

  * •

Position angle: Observations at non-parallactic angles are desirable.

###  3.5 G5: Lithium production in novae

The detection of the beryllium UV doublet in nova outburst ejecta tajitsu15
confirmed the long-standing theoretical prediction of substantial Li
production in nova thermonuclear runaways in accreting white dwarfs. Such
runaways soon develop density instabilities and convection plumes, allowing
the decay of processed 7Be. Any lithium brought by accretion from the
secondary star sinks to deep layers, being destroyed at the onset of the
process. Therefore, Li enhancement recently observed in novae is associated
with production during the thermonuclear runaway itself. Once the 7Be line-
formation conditions and radiative transfer are known, measuring the UV 7Be
lines during outburst provides a gauge of the Li enhancement in the ejecta
(see, e.g. ExA_Izzo  in this Special Issue).

While the Li nucleosynthesis process was successfully verified, two
fundamental questions remain unanswered. The first concerns the determination
of the total ejected mass and the corresponding Li mass contribution to the
ISM from a single nova outburst. Equally important, the long-term contribution
from the diverse nova population (classical, recurrent and symbiotic) remains
to be effectively constrained by observations.

The inhomogeneous nature of novae ejecta implies geometry-dependent line
optical depths, calling for synoptic and/or large sample observations. A
statistically robust sample of Li abundances in recent novae as well as a
precise modelling of the atmosphere and ejecta are needed to evaluate the
current and past contribution of novae to the Galactic Li abundance.

####  3.5.1 G5: Key requirements

  * •

Spectral resolution: Essential: $R$ $>$ 20,000 (to study the Be doublet).

  * •

Spectral coverage: Essential: 300-400 nm; Goal: 300-420 nm.

  * •

Simultaneous coverage at longer $\lambda$: Useful but not essential.

  * •

Position angle: A configurable slit direction is essential to study resolved
shells and to enable differential techniques.

###  3.6 G6: Metal-poor stars & light elements

Characterising the first stars that formed from primordial material produced
by Big Bang Nucleosynthesis is a vibrant field. The chemical imprints of the
first stars are detectable in the second generation of stars, including long
lived low-mass stars. While these are exceptionally rare, the tell-tale sign
is the extremely low metal content. Only a handful of these so-called ultra-
metal-poor stars are known; these are the most chemically ancient objects
known. They exhibit an enormous range in their relative chemical abundance
ratios, which demand a wide variety in the properties of the first supernovae,
e.g., mass, explosion energy, rotation, mass cut, explosion mechanism. More
data at the lowest metallicities are needed to better constrain the diversity
of properties, and relative frequency, amongst the first supernovae and
characterise the properties of the parent first stars.

There are many surveys planned for the future that promise large number
statistics on ultra-metal-poor stars. These are identified using photometric
and lower resolution surveys, and are typically selected using broad
indications of metallicity rather than precise measurements. Accurate and
detailed chemical abundance patterns of such candidates are essential to gain
the full insight into the range of nucleosynthesis processes that were
available in the early Universe, as well as to probe the physical properties
and mass distribution of the very first stars.

Many potential targets in this context will be drawn from the so-called
Carbon-enhanced metal-poor (CEMP) stars, with [C/Fe] $>$ $+$1.0 (see bc05 ).
Although rare, they have a diverse range of abundances of neutron-capture
elements, commonly grouped as: ‘CEMP-no’ (no over-abundance of r-process
elements), ‘CEMP-r’ and ‘CEMP-s’ (stars with over-abundances of r- and
s-processed elements, respectively) and ‘CEMP-r/s’ (with contributions from
both processes enriching their photospheres).

Abundances of CNO bring a wealth of information on stellar evolution and the
chemical evolution of the Galaxy. In contrast to the atomic transitions of the
elements discussed in the other Galactic cases, CNO features in the near UV
are dominated by a series of molecular bands (see Fig. 5) which include the
A-X OH transitions at the shortest wavelengths that can be used to estimate
oxygen abundances.

A range of scenarios have been explored to investigate these patterns,
including rotational mixing in rapidly rotating, low-metallicity stars (e.g.
chiappini13 ; choplin16 ) and supernova models which include both mixing and
fallback of material to yield the observed abundance ratios (e.g. un02 ; un05
; t14 ). In short, CEMP stars are perfect probes to investigate
nucleosynthesis from the first stars (including production of neutron-capture
elements) as well as mass transfer in binary systems (e.g. abate15 ). However,
comprehensive near-UV spectroscopy of CEMP stars from the ground has been
limited to date by the sensitivity of current facilities to just a few targets
(e.g. placco15 ; hansen15 ; hansen19 ).

A more expansive discussion of these objects in the context of CUBES is given
elsewhere in this Special Issue (ExA_Bonifacio ; ExA_Hansen ).

####  3.6.1 G6: Key requirements

  * •

Spectral resolution: $R$ $\geq$ 20,000; Goal: $R$ $\geq$ 30,000.

  * •

Spectral coverage: Essential: 305-400 nm; Goal: 300-400 nm.

  * •

Simultaneous observations at longer $\lambda$: Desirable in terms of
observational efficiency, but not essential.

Figure 5: Red: Synthetic spectrum of a metal-poor star generated using the
turbospec radiative transfer code plez12 , adopting physical parameters as for
CS 31082–001 cayrel01 ; hill02  including [Fe/H] $=$ $–$2.9. Grey: Synthetic
spectrum of the same star, but now N-rich ($\Delta$[N/H] $=$ $+$ 2.0 dex).

###  3.7 G7: Neutron-capture elements

Abundances of the heavy elements are often described in terms of three mass
peaks (see Fig. 6), within which the respective r- and s-process elements are
thought to have common formation channels. These peaks are connected to the
nuclear magic numbers 50, 82 and 126. Nuclei with these ‘closed’ nuclear
shells are especially stable and are produced in greater abundance during
n-capture reactions. The s-process path closely follows the valley of
$\beta$-stability, whereas the r-process encounters the closed shells in
nuclei of lower proton number, Z, than the s-process. The abundance peaks of
the r-process are therefore shifted to lower-mass atoms relative to the
s-process. The s-process elements are highlighted in blue in Fig. 6. The main
s-process, responsible for Zr to Pb, is associated with stars on the
Asymptotic Giant Branch (AGB, e.g. busso99 ), while the production of lighter
s-process elements (Fe to Zr) is thought to be dominated by the weak s-process
component during core He-burning and C-shell burning phases of massive stars
($M$ $>$ 10 $M_{\odot}$) (e.g. p10 ).

The formation channels for the r-process are particularly topical following
the detection of the GW170817 kilonova from a binary neutron-star merger
pian17 ; smartt17 ; watson19 . The r-process is thought to occur both during
the merging and in the milliseconds afterwards (e.g. bovard17 ) and is thought
to play a major role in the chemical evolution of the Galaxy matteucci14 ;
cescutti15 . Other predicted sites of r-process nucleosynthesis included low-
mass (8-10 M⊙) core-collapse supernovae mc90 ; iw99 . However, according to
nucleosynthesis calculations, normal core-collapse SNe do not appear to host
the proper conditions for the production of r-process arcones07 . In contrast,
the r-process is thought to take place in a special family of core-collapse
SNe which have magneto-hydrodynamically-driven jets arising from rapidly-
rotating massive stars with a strong magnetic field winteler12 ; ntt15 .

Figure 6: Abundances of isotopes in the Solar System vs. mass number (A, where
A = Z + N, in which N is the neutron number) normalised by the abundance of
28Si to 106 (from lodders03 ). The iron peak and the regions of the three
peaks of r-process and s-process are indicated.

Determining the abundances of neutron-capture elements in metal-poor stars is
fundamental to understand the physics of the different processes outlined
above (slow, rapid and possibly intermediate), as well as the chemical
evolution of the Galaxy. Neutron-capture elements have been measured in over
1000 stars in the Galactic halo (with at least one neutron-capture element
measured, typically Ba). Looking ahead, 4MOST and WEAVE will provide Ba
estimates (at least) for several 105 stars. Most of these will be in the
Galactic disk, but their samples will also include halo stars. However, a
clear understanding of the different processes can only be reached with
detailed measurements of all the key elements, many of which are in the UV
domain (e.g. Hf, Sn, Ag, Bi) and have only been measured for a tiny sample of
stars. CUBES will be the instrument to fill this gap.

The full pattern of neutron-capture elements (coupled with stellar kinematic
data from Gaia) could also be fundamental to understand the origin of the
Galactic halo (e.g. r18 ). In the wider context of Galactic Archaeology and
chemical evolution of the Milky Way, it is important to have a sample that
evenly covers the range of metallicities of halo giants (not only the most
metal poor) with as much information on the neutron-capture elements as
possible. Gaia and 4MOST will help with target selection, and CUBES will
provide the UV spectra to estimate the abundances.

To illustrate the expected performance of CUBES in the determination of heavy
elements, detailed simulations of the Ge I 303.9 nm and Hf II 340.0 and 371.9
nm lines are presented elsewhere in this Special Issue ExA_Ernandes . These
demonstrate the gain of up to three magnitudes that CUBES will provide
compared to current observations.

Building on analysis undertaken before the Phase A study ernandes20 , we also
explored the precision limits for stellar parameters and abundances at $R$ $=$
20,000 more quantitatively using the Chem-I-Calc package swt20 . This
estimates the precision using a Fisher information matrix and the Cramér-Rao
inequality to find the Cramér-Rao Lower Bound, which is a method to find the
lower bound to the root-mean-square of an unknown, non-biased variable. We
used the full-spectrum fitting capability of Chem-I-Calc, adopting the 1D LTE
models provided by the package. Simulations for observations with S/N $=$ 100
for two template models for RGB stars with [Fe/H] $=$ $-$1.5 and $-$2.5 are
shown in the upper and lower panels of Fig. 7, respectively.

Fig. 7 shows that for the stellar parameters (Teff, log($g$), $v_{\rm micro}$)
and the majority of the elements we can expect to recover a precision to
better than 0.1 dex at $R$ $=$ 20,000 for observations with S/N $=$ 100\. For
many of the heavier elements a precision of $\sim$0.2 dex is predicted. We
note that for elements with relatively few and weak lines, it is preferable to
use fits to single lines rather than the full spectral fitting.

Figure 7: Simulations using the Chem-I-Calc package swt20  to investigate the
precision to which stellar effective temperature (Teff), gravity (log($g$)),
microturbulence ($v_{\rm micro}$) and chemical abundances ($\sigma$[X/H]) can
be recovered for observations of stars on the red giant branch (RGB) with
CUBES (at $R$ $=$ 20,000, S/N $=$ 100), for models with [Fe/H] $=$ $-$1.5 and
$-$2.5 in the upper and lower panels, respectively.

####  3.7.1 G7: Key requirements

  * •

Spectral resolution: $R$ $\geq$ 20,000.

  * •

Spectral coverage: Essential: 300-400 nm; Goal: 300-420 nm.

  * •

Simultaneous coverage at longer $\lambda$: Desirable for observational
efficiency, but not essential.

###  3.8 G8: Precise metallicities of metal-poor pulsators

Pulsating stars such as Classical Cepheids and RR Lyrae stars are
traditionally used to derive distances within the Local Group and to calibrate
the extragalactic distance scale, and secondary distance indicators, at
cosmologically interesting distances. Interest in their properties has been
renewed by the revolution in astrometry from Gaia and recent tensions in the
Hubble constant between local estimates based on distance indicators (e.g.
riess19 ) and early Universe measurements and predictions (see verde19 ).

Classical Cepheids have Period-Luminosity (PL) and Period-Luminosity-Colour
(PLC) relations. Due to the finite width of the instability strip, PL
relations are better defined towards near-IR filters, whereas using optical
filters entails a colour term via PLC relations or (reddening-free) Period-
Wesenheit (PW) relations (e.g. marconi05 ; marconi10 ; ripepi12 ; fiorentino13
; ripepi17 ; ripepi19 . In contrast, lower-mass RR Lyrae stars do not obey PL
relations in the optical, but they are accurate Population II distance
indicators via near-IR PL relations and optical/near-IR PW relations.

In this context, the universality of PL, PLC and PW relations for both
Cepheids and RR Lyrae is uncertain. Given the many theoretical and
observational effects, there is no consensus on the metallicity dependence of
the Cepheid PL and PW relations (e.g. marconi05 ; ripepi19 ; romaniello08 ;
bono10 ). Metallicity differences between the calibrating sample and the
target galaxy could induce systematic effects of $>$10% in calibration of the
extragalactic distance scale.

RR Lyrae stars are a promising (Population II) alternative for calibration of
the cosmic distance ladder, with several authors concluding that the $K$-band
PL relation includes a metallicity term (e.g. bono01 ; bono03 ; sollima06 ;
marconi15 ) but without agreement on its magnitude. Moreover, the relative
abundance of CNO elements in metal-poor RR Lyrae stars can simulate a
different global metallicity, giving variations in the pulsation properties
and in the relation between period, magnitudes and colours. In particular, CNO
abundances can affect the intrinsic luminosity of Horizontal Branch pulsating
stars and also the coefficients of the traditionally adopted absolute visual
magnitude versus [Fe/H] relation.

The CNO abundance had long been considered as constant among stars in the same
cluster, but within the discoveries of multiple stellar populations in GCs,
the summed CNO abundances have also been seen to vary in some clusters. To
support the dramatic improvement in astrometry from Gaia and the chance to
study variables out to cosmologically-interesting distances, we need accurate
abundance determinations of individual pulsating stars in external galaxies to
quantify the impact of such abundance variations. In particular, CUBES will
allow us to investigate iron-peak elements in metal-poor stars, as well as
estimating the CNO abundances from the molecular bands in the UV range (see
Case G6).

Another interesting application at UV wavelengths is to investigate pulsating
low-mass horizontal branch (HB) stars. During the pulsational cycle their
effective temperatures undergo large variations, where they can become hotter
(cooler) than the blue (red) boundary of the instability strip. The effective
temperatures of RR Lyrae stars are usually obtained from colour relations that
have uncertainties of at least $\pm$100K. Spectroscopic determinations of
temperature would avoid using such relations. Moreover, accurate sampling of
the temperature variations during their pulsational cycles would also allow us
to investigate (via comparisons with suitable pulsation models) the impact of
hydrodynamical processes such as convective shocks on their stellar envelopes.

####  3.8.1 G8: Key requirements

  * •

Spectral resolution: $R$ $\geq$ 20,000.

  * •

Spectral coverage: Essential: 305-400 nm.

###  3.9 G9: Horizontal branch stars in Galactic GCs

Observations of HB stars provide another route to investigate the origins of
multiple populations in GCs. In addition to the spectroscopic results that
highlight star-to-star differences in GCs gratton12 , high-quality space
photometry has further reinforced the idea of multiple populations. This is
manifested by the spread in the main sequence in GCs, which cannot be
explained by photometric uncertainties alone. For instance, the HST study of
NGC 2808 was particularly important as it revealed three discrete components,
that were interpreted as due to different helium abundances of the stars in
each sequence, with the bluest stars the most helium-rich piotto07 .

These results argue that the stars with anomalous chemistry (on the blue side
of the main sequence) formed from gas processed by proton-capture
nucleosynthesis. This is supported by the finding of a significant spread in
the Al content in NGC 2808, with stars on the bluer sequence being Al-enriched
bragaglia10 . As noted in Section 3.4.2, such results have stimulated theories
regarding possible self-enrichment mechanisms. Possible sources of this
polluted material from the first generation of stars include massive AGB stars
dercole08 , rapidly-rotating massive stars decressin07 , super-massive stars
dh14  and interacting massive binaries sdm09 . Each of these scenarios have
their own difficulties when compared to the observational results (see e.g.
renzini15 ).

####  3.9.1 HB stars in the context of multiple populations

The HB feature in the colour-magnitude diagram (CMD) of GCs is populated by
stars brighter than the main-sequence turn-off and bluer than their
counterparts evolving through the red giant branch (RGB). HB stars are
currently in the core He-burning phase, which starts in the He-flash modality
once the stars reach the tip of the RGB, and continues in a quiescent mode
after the electron degeneracy in the core is removed. The CMDs of GCs obtained
so far offer a variety of situations regarding the HB extension and
morphology. Some GCs are characterised by narrow and red HBs (e.g. 47 Tuc;
sarajedini07 ), others exhibit very extended HBs, harbouring groups of stars
significantly fainter and bluer than the main HB component (e.g. NGC 2808,
bedin00 ; NGC 2419, dc15 ).

The location of HB stars in the CMD is primarily related to the mass of the
envelope on the He-burning core. This is influenced by factors such as the age
and metallicity of the cluster, and the mass-loss history during the RGB
phase. Unfortunately, some of these factors counterbalance each other, which
introduces degeneracies in the interpretation of the observed HB morphologies.
In addition to the mass distribution, other physical properties such as
chemical composition have also been suggested to explain the peculiar HBs of
some GCs.

####  3.9.2 He abundances of HB stars

To illustrate the potential contribution of CUBES, we consider the GC NGC
6441. Its CMD has a well developed blue HB, with a strong slope upwards from
the red clump to the blue of the RR Lyrae region piotto02 . The peculiar
features of the HB of NGC 6441 are not explained by conventional evolutionary
models, which is a strong indication of the co-existence of at least two
generations of stars. One suggestion is that a spread in the helium content
(of $\Delta Y$ $\sim$ 0.1) between the stars on the red and blue side of the
HB might account for the morphology dantona02 ; cda07 . Direct measurement of
the He content of the HB stars is required to test this – if confirmed then
the difference in He abundance would indicate that the most contaminated stars
formed from gas with $Y$ $\sim$ 0.35, which is consistent with contamination
from super-AGB stars.

Synthetic model spectra ($R$ $=$ 20,000) over the 300-420 nm range of stars
with log($g$) $=$ 4.0 and Teff $=$ 25 and 13 kK are shown in Fig. 8. These
examples correspond to stars on the hot and cool side of the blue HB in NGC
6441. There are multiple He I lines in this wavelength region that can be used
to estimate the He-mass fraction. Specifically, in the hotter star these
include: He I 318.8, 335.5, 344.8, 358.7, 363.4, 370.5, 382.0, 386.9 and 387.2
nm. In the cooler star only the latter three lines are useful, and the use of
He I 402.6 nm becomes more important. Thus, extension of the CUBES range to
400-420 nm is desirable for this case.

As for NGC 6441, the majority of GCs investigated so far have extended HBs
(e.g. the case of NGC 2808, see piotto07 ; bedin00 ). By obtaining estimates
of surface helium abundances CUBES will provide important information towards
understanding the multiple populations in GCs, with no other instrument
capable of achieving the same result for our targets.

To achieve the desired precision, we require S/N $\gtrsim$ 50 to estimate the
He abundance. Using the model spectra shown in Fig. 8 as input templates to
the CUBES ETC for a target with $V$ $=$ 16 mag., the predicted exposure times
to reach S/N $=$ 50 are 1 and 4 hrs for the 13 and 25 kK models, respectively.
In contrast, UVES observations would require integrations of 8 and 24 hrs for
the same stars.

The He I lines at longer wavelengths can also be used to partly address this
case, but UV observations with CUBES will take advantage of the rising
spectral energy distribution towards shorter wavelengths for the hotter stars
(e.g. upper panel of Fig. 8), enabling observations to more distant clusters
than currently possible. In addition to the He abundances, CUBES observations
of HB stars will also provide important insights into the CNO abundances and
distribution of s-process elements (see Cases G6 and G7 above).

Figure 8: Synthetic spectra of stars with effective temperature 25 kK (upper
panel) and 13 kK (lower panel) in the 300-420 nm range. The wavelength scale
in the lower panel is limited to $\lambda$ $>$ 360 nm (3600 Å), as there are
no significant helium lines at shorter wavelengths at this temperature.

####  3.9.3 G9: Key requirements

  * •

Spectral resolution: $R$ $\gtrsim$ 10,000; Goal: $R$ $\sim$ 20,000.

  * •

Wavelength coverage: Essential: 300-400 nm; Desirable: 400-420 nm.

  * •

Simultaneous observations at longer $\lambda$: Observations with UVES are a
useful complement.

###  3.10 G10: Early-type companions in binary Cepheids

Most Cepheid variables (up to 80%) are expected to be members of binary
systems, where their companions are main-sequence or slightly evolved early-
type stars, which are typically much fainter in optical bands and hardly
detectable. One of the methods to detect such companions is to look at the Ca
II H and K lines. Normally they should have similar depths, but when a hot
companion is present its strong H-epsilon line is superimposed on the Ca II H
line, which appears significantly deeper than the Ca II K line. Characterising
the orbital properties of these systems beyond the Milky Way is challenging,
with current spectroscopy limited to V $\sim$15.5 mag (e.g. pilecki21 ). CUBES
will provide the greater sensitivity needed for efficient observations of
larger samples in the Magellanic Clouds.

The motivation for this case is to better characterise Cepheid binary systems
in the Clouds down to V $=$ 17-18 mag. The improved efficiency of CUBES will
enable us to better detect main-sequence companions beyond the reach of
current facilities, either via the Ca II H and K lines, or perhaps via the
broadening function technique pilecki21 . The latter is less demanding in
terms of the required S/N given the use of multiple lines, but tailored tests
are needed for its applicability to the CUBES range.

####  3.10.1 G10: Key requirements

  * •

Spectral resolution: $R$ $\geq$ 20,000; Goal: $R$ $\sim$ 40,000.

  * •

Spectral coverage: Essential: 380-400 nm.

###  3.11 G11: Extragalactic massive stars

With masses of more than ten times that of the Sun, massive OB-type stars burn
brightly for just a few million years before exploding violently as
supernovae. During their lives they manufacture the materials from which
planets and life are made and, via their high-energy winds and explosive
deaths, they shape the chemistry and evolution of their host galaxies a82 ;
a13 ; ebm18 . Only by learning how these stars behave in the local Universe
can we attempt to interpret distant, star-forming galaxies, whose light is
dominated by vast numbers of these spectacular objects, and that are used to
chart the history of the Universe from the Big Bang through to the present.

A focus of the past 20 years has been to explore the massive-star populations
in galaxies beyond the Milky Way, to investigate their properties in the
metal-poor environments of the Magellanic Clouds and dwarf galaxies in the
Local Group and beyond.

As expanded further in an article in this Special Issue ExA_Evans , we have
been able to study the properties (including oxygen abundances) of B-type
supergiants in metal-poor galaxies in targets at 1-2 Mpc (e.g. evans07 ;
castro55 ) but we lack the sensitivity for comparable studies of O-type stars
at such distances.

The ground-UV domain is relatively unexplored in the context of massive OB-
type stars, with only a few published observations (e.g. lamers72 ; dufton ;
drissen95 ). However, there are several He I absorption lines shortwards of
400 nm, as well as the higher members of the Balmer series up to the Balmer
limit. The latter can provide constraints on the spectroscopic gravities,
particularly for low-metallicity stars where wind effects are relatively
minimal and comparisons with plane-parallel model atmospheres are sufficient
for quantitative analysis. Moreover, above effective temperatures of Teff
$\sim$ 27.5 kK, absorption from He II 302.3 nm becomes increasingly strong,
offering a potential temperature diagnostic for O-type stars when compared
with, e.g., the strength of the He I 318.8 nm line.

There are also a large number of O III absorption lines in the ground UV (plus
lines from O IV and O II at the hotter and lower temperatures, respectively,
of O-type stars). Given sufficient S/N ($\gtrsim$ 100), these O III lines can
be used to provide robust estimates of oxygen abundances (e.g. martins15 ),
that are required to test model predictions of mixing in stellar evolution
models (complementing carbon and nitrogen abundances obtained from
observations at longer wavelengths).

When coupled with the rapidly increasing spectral energy distributions of
O-type stars over the ground UV, observations with an optimised near-UV
spectrograph can potentially bring targets into reach that are otherwise too
faint at longer wavelengths. For instance, from calculations using the CUBES
ETC, it should be possible to obtain high-quality (S/N $>$ 100) spectra of
O-type stars in metal-poor systems at a distance of $\sim$1 Mpc in a total
integration of $\sim$3 hrs (for more detailed performance estimates see
ExA_Evans ). This will extend such abundance studies for O-type stars beyond
the Galaxy and Magellanic Clouds for the first time, sampling a broad range of
metallicities and galaxy types.

####  3.11.1 G11: Key requirements

  * •

Spectral resolution: Essential: $R$ $\geq$ 5,000; Desirable (for some
applications): $R$ $\sim$ 20,000.

  * •

Spectral coverage: Essential: 305-400 nm; Goal: 300-415 nm.

##  4 Extragalactic Cases

###  4.1 E1: Primordial deuterium abundance

Cosmology is in a golden age. We have established a ‘Standard Model’ of
particle physics and cosmology, and we know the content of the Universe to
within a few percent. However, we still have no accepted model of dark energy
and dark matter. Somewhat remarkably, we do not know why the Universe contains
baryons instead of antibaryons, nor why the Universe contains baryons at all
(i.e. the baryon asymmetry). We are also missing crucial properties of
neutrinos, such as their hierarchy, why they change flavour, and the number of
neutrinos that existed during the earliest phases of the Universe. These
limitations indicate there is missing physics in our understanding. Thus, new
cosmological observations may allow us to discover something truly fundamental
about the early properties of the Universe.

Some of the above questions can be studied and better understood by measuring
the nuclides that were produced a few minutes after the Big Bang – the so-
called ‘primordial elements’. The primordial deuterium abundance (D/H) is
currently our most reliable probe of Big Bang Nucleosynthesis. It can be
determined to $\sim$1% precision using near-pristine gas clouds that imprint D
I and H I absorption lines on the spectrum of a background quasar. The gas
clouds where this measurement is possible are extremely rare, with just seven
systems with precise measurements of D/H cooke18  (see Fig. 9), most of which
are at redshift $z$ $\sim$2.8-3.1, placing the weak D I Lyman series
absorption lines at an observed wavelength of $<$400 nm. The accessible
redshift range is due to a combination of the efficiency of current
spectrographs plus the atmospheric cut-off at UV wavelengths and, on the other
hand, of the increasing thickness of the H I Lyman-$\alpha$ forest (which can
contaminate potential D/H systems) at high redshift.

Figure 9: High-precision D/H measures as a function of oxygen abundance
cooke18 ). The weighted mean is shown by the red dashed and dotted lines,
which represent the 68% and 95% confidence levels, respectively. Assuming the
Standard Model of cosmology and particle physics, the right-hand axis shows
the conversion from D/H to the universal baryon density. This conversion uses
the theoretical determination of the d(p,$\gamma$)3He cross-section from m16 .
The dark and light shaded bands are the 68% and 95% confidence bounds on the
baryon density derived from the CMB planck16 .

To secure a measurement, the gas cloud needs to be near-pristine with
unusually quiescent kinematics, and not blended by unrelated contamination
from the forest of Lyman-$\alpha$ absorption lines due to the intergalactic
medium along the line-of-sight. CUBES will deliver three key improvements over
previous instruments that have been used for this measurement (specifically,
VLT-UVES and Keck-HIRES):

  * •

Twice as many quasars are known at 2.5 $<$ $z$ $<$ 2.8 than at 2.8 $<$ $z$ $<$
3.1 (the latter is the range used by current studies, the former is the range
that will be enabled by CUBES). This increases the sample size by a factor of
three.

  * •

The impact of unrelated blending from the Lyman-$\alpha$ forest is much less
significant at low redshift, leading to cleaner measurements.

  * •

The increased sensitivity of CUBES at blue wavelengths will allow fainter
quasars to be observed. Thanks to the steep shape of the quasar luminosity
function, this will increase the sample size by a factor of at least 10
(simply because of the access to more quasars). Alternatively, the brightest
quasars can be recorded in a fraction of the time compared to UVES or HIRES.

CUBES should be able to improve the sample size by one-to-two orders of
magnitude, while also providing cleaner measurements. UVES observations of the
quasar spectrum at $>$400 nm are also important to complement the CUBES data.
This is because we need to record the H I Lyman-$\alpha$ absorption line, in
combination with the narrow metal absorption lines (with Doppler widths of
$\sim$3-5 km s-1) belonging to the same redshift system, which reveal the
structure of the velocity profiles. UVES has the most appropriate wavelength
range and spectral resolution for this goal (i.e. it will resolve the key
metal lines and has the ability to efficiently record the spectra of faint
$u_{\rm AB}$ $>$ 19 quasars).

Calculations with the CUBES ETC predict that a quasar at $z=2.7$ of $u_{\rm
AB}$ $=$ 19, 19.5, 20 mag could be recorded with the required S/N $=$ 20 at
320 nm in $\sim$ 2, 5, and 13 hr of observation, respectively (with 1 hr
exposures, airmass $=$ 1, two days from new moon and seeing of 0.8′′). For
this case the continuous wavelength coverage provided by the CUBES design is
ideal, as there are 10-15 D I absorption lines that could be detected in any
given system, and these lines could fall anywhere in the 320-380 nm wavelength
range due to the redshift range of the targets.

####  4.1.1 E1: Key requirements

  * •

Wavelength range: CUBES alone will not be able to simultaneously cover the
high order D I Lyman series lines and the H I Lyman-$\alpha$ line. So the
ideal wavelength range should be 320–380 nm to cover the D I absorption lines
between $2.5<z<3.1$, together with simultaneous UVES observations redwards of
$\sim 400$ nm to cover the H I Lyman-$\alpha$ line and resolve the metal
absorption lines (the latter is needed to determine the component structure
and requires the $R\sim 40\,000$ provided by UVES).

  * •

Spectral resolution: For gas at $10^{4}$ K, the D I lines have a width of
$\sim$10 km s-1. To resolve the line widths of deuterium absorption and assess
line saturation, a spectral resolving power of $R\sim 20\,000$ is required.

  * •

Spectral sampling: The line widths to be observed will be narrow (sometimes
comparable or less than the instrument resolution). Therefore, this science
case requires the data to at least be sampled at the Nyquist frequency (2–3
pixels FWHM) over the entire wavelength range.

###  4.2 E2: Missing baryonic mass in the high-$z$ CGM

An active field of extragalactic research is the so-called ‘missing baryon’
problem, where the detected baryonic mass compared to the universal fraction
suggests that we are missing more than 60% in galaxies outside of clusters,
rising to $\geq$90% for gas-dominated galaxies and dwarfs (e.g. mcgaugh10 ).
Studies with HST of the CGM in galaxies at $z$ $\sim$ 0.2 concluded that
diffuse gas in their halos could account for at least half of the previously
missing baryonic mass werk14 . A significant contribution to the baryonic mass
is also expected from the hot diffuse gas (T $\sim$ 105.7 to 106.3 K) traced
by O VII and detectable in the soft X-rays nicastro18 .

The dispersion observed in a sample of localized Fast Radio Bursts (FRBs, e.g.
chatterjee17 ) at $z$ $\sim$0̇.1-0.5 has been used to measure the electron
column density and accounts for every ionised baryon along the line of sight
macquart20 . This independent measurement of the baryon content of the
Universe with a first sample of five FRBs is consistent with the CMB and Big
Bang Nucleosynthesis values, substantially solving the missing baryon problem
at low redshift. This result has established the total amount of baryons, but
FRB observations cannot distinguish the share of the different gaseous
components. In this respect, quasar absorption line studies will play a key
role and, in particular, they could help understand whether a substantial
fraction of the mass at high redshift is also contained in diffuse halo
components.

Several current and future instruments are being designed with the study of
the high-redshift ($z$ $>$ 2) CGM in mind. The Keck Cosmic Web Imager (KCWI)
kcwi  was partly designed to address this topic (although its blueward
coverage stops at 350 nm), and one of the cases developed for the ELT-MOSAIC
instrument (e.g. puech18 ) is to use high-redshift Lyman-break galaxies as
background sources to reconstruct the 3D density field of the IGM and to probe
the CGM of galaxies at $z$ $\sim$ 3 japelj19 .

Observations in the ground-UV can play an important role in this field
(complementing observations with HST and the ELT in the future), providing
access to the baryons in halos at $z$ $=$ 1.5 to 2, immediately after the era
of peak star-formation in the Universe (e.g. md14 ). The relatively low number
density of lines in the Lyman-$\alpha$ forest at these redshifts implies that,
even with observations at $R$ $\sim$ 20,000, the impact of blending on the
measurement of the CGM metal lines can be minimised (which is not the case at
larger redshifts). Furthermore, the proposed wavelength range for CUBES will
enable observations of the important O VI line (used to trace the warm-hot gas
(see t11 )) to redshifts $2$ $<$ $z$ $<$ 2.8. In both cases, to derive the
baryon contribution of the gas, complementary UVES (or ESPRESSO) spectra will
be required to complete the coverage of the metal transitions and of the
Lyman-$\alpha$ forest.

A significant improvement in near-UV efficiency cf. UVES will open up
background sources 1-2 magnitudes fainter than the quasars used at present,
significantly increasing the number (and the space density) of available
targets and/or the possibility to obtain very high signal-to-noise ratio
spectra to detect the faint lines (e.g. vdo16 ).

To be able to characterise absorption systems associated with galaxies at
impact parameters of up to a few hundreds of physical kpc (e.g. werk14 ) we
will need to carry out preparatory studies to identify and determine the
properties of the galaxies in the fields of the quasars forming our sample.
The recent QUBRICS survey calderone19 ; boutsia21  is being used to identify
quasars in the southern hemisphere with $i_{\rm AB}$ $<$ 18 mag. and will
provide the 30-40 quasars needed for this case.

Detailed simulations of the proposed observations are presented elsewhere in
this Special Issue ExA_DOdorico . These demonstrate that observations reaching
S/N $\sim$ 15 per pixel would be possible for a sample of 10 targets at $z$
$=$ 2, 20 at $z$ $=$ 2.5 and 10 at $z$ $=$ 3 (all with a conservative $u_{\rm
AB}$ $=$ 18 mag.) in a total programme of 13 hrs.

####  4.2.1 E2: Key requirements

  * •

Position angle: Non-parallactic observations are desirable to observe quasar
pairs or strongly-lensed quasar images (arcs) for tomographic studies.

  * •

Observations at longer $\lambda$: Required, so simultaneous observations would
give greater operational efficiency.

###  4.3 E3: Cold gas at high redshift

Understanding the physical and chemical properties of the gas in and around
galaxies is fundamental for a complete theory of star-formation along the
history of the Universe. Molecular hydrogen (H2) is a critical species to
achieve this goal since it traces the cold (T $\sim$ 100 K) component of the
neutral gas, which is inclined to gravitational collapse through Jeans
instabilities. Indeed, the cooling of the gas through atomic emission lines
depends on the abundance of metals – to which the formation of H2 is also
linked – and both the gas heating and the dissociation of H2 depend on the
strength of the UV field. In addition, H2 is vital for the formation of other
molecules that participate in the gas cooling, eventually initiating the
collapse of the cloud that gives birth to stars.

Molecular hydrogen can be detected directly through electronic absorption
lines in the Lyman and Werner bands (90-110 nm rest-frame). These have been
observed since the 1970s towards bright nearby stars thanks to space-borne UV
spectrographs and greatly contributed to the understanding of the local ISM.
At $z$ $>$ 2, H2 lines are conveniently shifted into the optical domain and
become accessible from the ground. UVES and X-Shooter have both played key
roles in the investigation of diffuse molecular and translucent gas, probed by
H2 in the distant Universe towards quasars (e.g. nd08 ; balashev19 ) and
gamma-ray burst (GRB) afterglows (e.g. bolmer19 ). More recently, it has also
been shown that H2 can serve as an independent and sensitive probe of AGN
feedback when studied in the quasar environment itself nd19 ; nd21 . In short,
H2 together with other atomic (in particular C0) and molecular (HD, CO)
species, provide very sensitive and unique tools to investigate cold gas and
the prevailing conditions such as temperature, density, UV flux and cosmic ray
ionization rate (e.g. klimenko20 ; kosenko21 , and see Fig. 10).

While good progress has been made in selecting targets for study from large
spectroscopic surveys (e.g. SDSS), the field is now reaching a bottleneck with
the vast majority of the (newly) identified systems being out of reach with
current instruments (see Fig. 11). This is due not only to the steep shape of
the quasar luminosity function, but also to reddening of the quasar light by
the absorbing (high column density and dusty) gas itself.

Only a significant improvement in sensitivity in the ground-UV domain can
provide the observations required to transform this field. The efficiencies
envisaged for CUBES will:

  * •

Increase the sample size by 1-2 orders of magnitude. In addition to the
systems that are already available from pre-selection in SDSS, other surveys
in the coming years will also contribute to the discovery of further systems
requiring follow-up.

  * •

Provide access to more extreme environments (e.g. more dusty/molecular rich).

  * •

Enable cleaner measurements. Indeed, in addition to the high blue throughput
providing observations of fainter background sources (both quasars and GRB
afterglows), it will also open-up observations of systems at lower redshift
(reducing the huge gap in cosmic history that is still poorly explored).

  * •

Cover more H2 bands for a given system or redshift, enabling robust
measurements.

The observed population of H2 (and HD) rotational levels, with proper
modelling, enable the physical conditions prevailing in the associated diffuse
molecular gas of high-$z$ galaxies to be derived. Observations of the H2 (and
HD) lines require a careful compromise between spectral resolution and system
efficiency. For instance, X-Shooter (with $R$ $\sim$ 5,000) has been used a
few times to complement high-resolution ($R$ $\sim$ 50,000) UVES observations
(e.g. nd17 ). However, this is extremely costly in observing time and the
tenfold lower resolution of X-Shooter is far less suited to deblend the H2
lines from the Lyman-$\alpha$ forest, impeding robust measurements of H2 over
various rotational levels which encode crucial information about collisional,
radiative and turbulent processes. A spectral resolution of $R$ $\sim$ 20,000
is the sweet spot between efficiency and sufficient information for
quantitative analysis.

To illustrate this huge gain, we highlight example observations of a quasar at
z = 2.3 with $V_{\rm AB}$ $=$ 20 mag. At an airmass of 1.3, with 1′′ seeing
and in a 1 hr exposure, the ETC predicts a S/N per pixel at 350 nm of: 1
(UVES, binning 2$\times$2), 5 (X-Shooter, binning 2$\times$2), and 6 (CUBES,
no binning)111We note that the estimate provided in the paper on the H2
science case ExA_Balashev  are unrealistically optimistic, due to the use of a
preliminary version of the CUBES ETC..

Finally, we note that CUBES will be highly complementary to ELT observations
with ANDES (formerly HIRES) to measure the evolution of the Cosmic Microwave
Background temperature from the excitation of CO molecules. Indeed, recording
of H2 lines will be necessary to constrain the collisional contribution to the
excitation of CO.

####  4.3.1 E3: Key requirements

  * •

Wavelength coverage: Essential: 305-380 nm; Desirable: 305-400 nm.

  * •

Simultaneous observations at longer $\lambda$: To maximise the impact of CUBES
observations of H2 in the blue, observations at $>$ 400 nm are also desirable
to include the metal lines and associated fine-structure levels, ideally at
higher spectral resolution. This will enable more precise modelling of the
absorption profiles and robust derivation of the physical and chemical
conditions. Simultaneous observations with UVES at $>$ 400nm are thus ideal
for this case.

Figure 10: An example of tight constraints on the number density and
temperature in cold H2-bearing gas in galaxies at cosmological distances, that
can be obtained from observations of molecular hydrogen (H2 and HD) and fine-
structure levels of associated metal species (adapted from balashev17 ).
Figure 11: $G$-band magnitude of quasars with intervening H2 systems observed
over the past two decades. As progress has been made in selecting targets from
large surveys (e.g. SDSS) and also those that are molecular rich, we have
reached the magnitude limits of UVES (even with 10-15 hr exposures and, even
then, without much information on the physical conditions). X-Shooter has
extended the limiting magnitude, but its resolution is not sufficient to
estimate the physical conditions. CUBES will bring these fainter targets
within our grasp, while also opening-up observations at lower redshifts and of
absorption systems with more extreme properties (chemically enriched, high
columns).

###  4.4 E4: Reionisation

The search for sources responsible for cosmic reionization, likely to have
occurred at z $>$ 6.5, is one of the most pressing questions in modern
cosmology and galaxy evolution (e.g. f19 ). Galaxies are commonly thought to
be able to produce the bulk of the UV emissivity at high redshift (e.g.
robertson15 ) but the AGN population is also proposed as a relevant or
dominant contributor (e.g. giallongo19 ; but also see fontanot12 ; hs15  for
different views), and more exotic possibilities such as decaying particles
cannot be excluded (e.g. poulin15 ). Due to the very high cosmic opacity (e.g.
worseck14 ), it is not possible to directly study the Lyman continuum (LyC)
emission at the epoch of reionization or even in the aftermath (e.g., at $z$
$\geq$ 5). As a consequence, analogues of the first sources have to be found
and their properties studied at lower redshifts.

Quasars are known to be efficient in producing UV photons with the fraction
(fesc) of ionizing photons able to escape to the IGM close to 100% cristiani16
; grazian18 . In contrast, galaxies are much more numerous, but fesc has large
uncertainties. The census of LyC-galaxies is growing fast in the local
Universe (e.g. izotov16a ; izotov16b ; izotov18 ) and more slowly at high
redshift, 1 $<$ $z$ $<$ 4 (shapley16 ; vanzella16 ; bian17 ; vanzella18 ;
saha20 ; rt19 ), given the challenging observations required. Thus, the
contribution of star-forming galaxies to the cosmic ionization budget must be
inferred from other properties related to their ability to leak ionizing
radiation into the IGM.

While the construction of a sample of LyC emitters serves as a reference for
the identification of cosmic reionizers, the physical processes that made
these galaxies transparent to LyC is interesting by itself. The way ionised
channels are carved in the interstellar medium is matter of investigation,
especially in the local Universe (e.g. bik18 ; micheva19 ) where the detection
of LyC radiation is also a tool to catch this phenomena in the act heckman11 .
A recent work based on VLT/X-Shooter observations on a gravitationally-lensed
object at $z$ $=$ 2.37 shows that the LyC radiation can escape through ionized
channels carved by massive stars hosted in young massive star clusters
vanzella20 . The enhanced spatial resolution provided by lensing magnification
allows us to observe such channels - aligned along the ionizing sources - with
a better contrast, and consequently witness high fesc values (as high as 90%;
rt19 ).

It is not fully clear if the properties of local ($z$ $<$ 0.4) LyC-galaxies
really match those at high redshift ($z$ $>$ 2), despite the fact they are
often referred to as genuine analogues. It is therefore critical to
investigate the LyC leakage at high redshift (e.g. $z$ $>$ 2, 2 Gyr after the
Epoch of Reionization), to reduce the cosmic time interval from reionization
and avoid any tantalising extrapolation from $z$ $\sim$ 0\. This needs to be
done under the more favourable observational conditions: (1) using a
spectroscopic approach closely probing the Lyman edge at 91.2 nm steidel18 ;
(2) reducing the cosmic opacity as much as possible while keeping $z$ $>$ 2;
(3) observing the enhanced signal and spatial resolution provided by strong
lensing (e.g. vanzella20 ).

Combining these factors together for strongly-lensed, star-forming regions at
$z$ $>$ 2.35 argues for high-efficiency spectroscopy in the ground UV. In
addition to ensuring high throughput, such an instrument needs to be sky
limited at intermediate and high resolution (needed for disentangling in cases
of contamination) and with a relatively long slit for background subtraction.

####  4.4.1 Lyman continuum detection of a sample of faint AGN at z $\sim$ 3-4

A component of the UV background case is to measure the escape fraction of a
sample of $\sim$30 faint AGNs ($M_{\rm 1450}$ $=$ $–$23) at $z$ $\sim$ 3-4.
These AGN have been selected from an extensive spectroscopic campaign in the
COSMOS field, where multiwavelength coverage from X-rays to radio ensures that
the sample is complete and representative of the population of active
supermassive black holes at $z$ $\sim$ 3-4. Deep CUBES spectroscopy shortwards
of 400 nm (sampling shortwards of the restframe Lyman limit) will enable
precise estimates of fesc for these faint AGN and to derive the mean free path
of H I ionizing photons. These quantities are fundamental to derive the
contribution of faint AGN to the ionizing background at $z$ $\sim$ 4\. These
empirical results will then provide a solid foundation to explore the role of
faint AGN in the epoch of reionisation.

####  4.4.2 E4: Key requirements

  * •

Spectral coverage: Essential: 320-380 nm; Goal: 300-400 nm.

  * •

Spectral resolution: Essential: $R$ $>$ 1000; Goal: $R$ $>$ 10,000. High
resolution is not critical – the ideal resolution is the highest value
possible for which the detectors remain background limited (at all
wavelengths) for 1 hr observations of faint objects.

  * •

Detector dark current: The need to be background limited informs the
requirements on the acceptable dark current of the detectors. The dark current
should not generate more counts than a typical new moon sky. The variance in
read-out noise should not exceed the dark current in a 1 hr exposure (with 2
$\times$ binning), giving a requirement on the dark current of $<$ 1e-/hr and
on the read-out noise of $<$2e-/hr.

##  5 Transients

###  5.1 T1: Gamma-ray bursts

GRBs are extremely powerful explosions of cosmological origin lasting just a
few seconds (see piran13  and references therein). Their $\gamma$-ray (or
prompt) emission is followed by afterglow radiation at longer wavelengths
(from X-rays to the radio), which can last several days (e.g. zhang07 ). Such
sources fade very quickly, so the capability to repoint narrow-field
instruments to the $\gamma$-ray position and spot the afterglow has been
critical to understand these elusive phenomena. The Swift satellite
(specifically designed to study GRBs gehrels04 ) can detect the prompt event
in $\gamma$-rays, repoint its X-ray and optical-UV instruments automatically,
and disseminate arcsecond-precision coordinates of the afterglow in tens of
seconds. This information is exploited by robotic telescopes that follow-up
and gather information on the GRB just minutes after the explosion.

To follow-up such sources ESO developed a Rapid Response Mode (RRM), which
allowed the VLT to react promptly to a GRB alert for the first time. The RRM
procedure overrides execution of the current Observing Block and automatically
repoints to the afterglow position (after checking for visibility
constraints). This enabled UVES to obtain high-resolution spectra of GRB
afterglows only a few minutes after the prompt event, when the sources were
still bright. Indeed, the quickest spectrum was obtained for GRB060607 fox08 ,
just 6.5 minutes after the trigger, with better quality spectra (S/N $=$ 50)
obtained 8 minutes after trigger for the naked-eye burst GRB080319B, observed
when its magnitude was still $R$ $=$ 12 mag. delia09 . The capability to
quickly point the VLT to such sources has uncovered a terrific wealth of
information about the explosion sites of GRBs, their host galaxies and the gas
along their line of sight (enabling detailed studies of abundances,
metallicity and chemical enrichment up to cosmological scales, through the
analysis of the hydrogen and metal absorption lines).

In the context of spectroscopy of the afterglow, the UVES+RRM combination made
dissection of the close environment of GRBs possible via separation in
velocity space of the absorbing components from the same metal feature. Even
more importantly, time-resolved spectroscopy revealed a strong variation of
fine-structure lines from excited levels. The overall decrease in time of the
absorbing column density in these features was interpreted as the decay of the
excited levels into the fundamental states as less and less flux was available
for the gas to populate the higher energy states. This detection enabled
determination of which gas component in velocity space was close to the GRB
explosion site (as the burst flux directly influenced the gas physical state)
and even determination of the GRB-absorber distance, by comparing the column
density variation of the absorbing gas with time-dependent photo-excitation
models vreeswijk07 . Remarkably, this distance is quite large, ranging from
tens to several hundred parsecs. This shows how far the GRB flux can modify
its surrounding medium and provides important constraints on radiative models
of GRBs.

The primary diagnostics to determine the GRB-absorber distance are from Fe II.
Its reddest lines are at $\sim$260 nm (rest frame), so this kind of analysis
is difficult for GRBs at $z$ $<$ 0.8 because the throughput of UVES
dramatically reduces below 400 nm. CUBES will give a S/N that is at least five
times better than UVES for the same exposure time over 300-400 nm, so will
significantly increase the redshift range for studies of GRB afterglow to
investigate the distance between the GRB progenitor and the absorbing gas
(UVES has enabled this for GRBs at $z$ $\geq$ 1, but CUBES will bring events
at 0.25 $\lesssim$ $z$ $<$ 1 within reach).

We note that such studies require a bright afterglow and a fast reaction,
combined with clear weather and (in the past) the correct instrument mounted
at Paranal. These demanding constraints mean that only 15 GRBs have been
observed with UVES to date, with less than half yielding a spectrum with high
enough S/N and with the correct set of lines falling in the instrumental range
to warrant publication. The advent of CUBES will significantly increase these
numbers, securing data for close-by GRBs, for which the flux level should be
higher at constant luminosities. We add that, while the Swift satellite is
presently in good health, there are other missions in development (e.g. SVOM)
that will provide alerts once Swift is no longer operational.

For reference, $\sim$25% of long GRBs with measured redshifts are at z $<$ 1,
the fraction grows to $\sim$75% for the short GRBs. In terms of absolute
numbers, Swift typically detects 90-100 GRBs/yr, roughly 10 of which are
short. So we expect $\sim$30 long/short GRBs at $z$ $<$ 1 per year. However,
only (approx) 10% will be promptly visible from Paranal during the night, i.e.
events asking for possible RRM observations (although before the observation
we also do not necessarily know the redshift if targetting an afterglow). In
short, roughly one RRM per month is to be expected.

####  5.1.1 Key requirements

  * •

Simultaneous observations at longer $\lambda$: Simultaneous UVES observations
are fundamental to provide:

    1. 1.

Better handling of uncertainties such as skylines and intervening absorbers
(which may only affect specific portions of the spectrum).

    2. 2.

Insights into how our understanding of the GRB environment is influenced by
different diagnostic lines and spectral/velocity resolutions. In particular,
it will be useful to determine how the different instruments separate the
interstellar features, to study the saturation problem which arises as the
instrumental resolving power goes down, and to evaluate the goodness of the
fit with Voigt profiles (via data that could have a higher resolution but
lower S/N, and vice versa).

    3. 3.

Cross-matching the flux calibration of data from the two instruments,
especially for observations at high airmass.

  * •

Acquisition images: From past experience of GRB observations acquisition
imaging has provided useful scientific information, particularly for
instruments without imaging capabilities (e.g. X-Shooter). For transients, we
often have an inaccurate idea of the target spectrum and/or are taking
observations in sub-optimal conditions. Having one or more photometric points
from the acquisition camera would be useful to support analysis of the CUBES
data.

###  5.2 T2: Kilonovae

GRBs with $\gamma$-ray durations of $<$2 s (‘short-GRBs’) are likely produced
by compact binary mergers involving either two neutron stars or a neutron star
and a solar-mass black hole. The merger remnant ejects an ultra-relativistic
jet, producing GRB emission. The GRB afterglow originates from synchrotron
radiation produced by the interaction between the ultra-relativistic jet and
the surrounding medium. The luminous, short-lived transient also ejects
$\sim$0.01M⊙ of neutron-rich matter that creates a radioactively-powered
explosion called a kilonova or mini-supernova. The transient arises from the
radioactive decay of freshly synthesised heavy r-process elements lp98 , that
may be the dominant mechanism for their production throughout the Universe.

The short-GRB field was revolutionized by a kilonova event following the
simultaneous detection of a gravitational-wave chirp (GW170817, abbott17 ;
tanvir17 ). At $z$ $\sim$ 0, future detections of associated gravitational-
wave events will allow us to constrain kilonova heavy-element yields. However,
observations of short-GRBs up to $z$ $\sim$ 1 or farther are required to
constrain the evolving rate density of mergers to determine their contribution
to the enrichment of heavy elements over cosmic history. Moreover, it will
continue to be crucial to understand short-GRBs beyond the LIGO/Virgo horizon
of $\sim$200 Mpc (for neutron star-neutron star mergers) and $\sim$800 Mpc
(for black hole-neutron star mergers). The kilonova spectra can tell us about
the nuclear yields, neutrino flux, mass, velocity, and geometric orientations
of the ejecta. Kilonova spectra are the least studied, with only one
unambiguous case identified so far, and there is no broad consensus on the
expected spectrum (e.g., bk13 ). The theoretical models will remain
unconstrained in the absence of spectroscopy, particularly of the kilonova UV
emission. For AT 2017gfo, the electromagnetic counterpart to GW 170817, a
bright UV/optical flash was detected during the first hours of the evolution
of the event abbott17 ; evans17 ; arcavi18 . The flash evolved rapidly, and
after a few hours the kilonova emission peaked in the infrared due to
absorption by the heavy r-process elements.

The Swift satellite can detect the short-GRB in $\gamma$-rays and can localise
the afterglow in a few seconds after the trigger. Short-GRB afterglows are
faint and fade fast. Prompt reaction to a GRB alert with CUBES will maximise
chances of obtaining a redshift from the transient and/or host and detection
of elements present in its spectrum. The CUBES wavelength range is also
crucial to study the early emission from kilonovae and to model the blackbody
in the UV range. Multi-epoch observations will allow determination of the
effective temperature of the source and the status of the photosphere at days
and weeks after the event. CUBES will significantly increase the number of
observations and detection of short-GRBs that essentially are at lower
redshifts and cannot be captured by the sensitivity of UVES.

####  5.2.1 Key requirements

  * •

Simultaneous observations at longer $\lambda$: Simultaneous UVES observations
are desirable to give clues on the environment of the kilonova and its host
galaxy.

###  5.3 T3: Superluminous Supernovae

Superluminous supernovae (SLSNe) are a class of extremely energetic SNe
quimby07 ; gal-yam12 ; gal-yam19 , which are at least 10 and often 100 times
more luminous than normal SNe, and evolve slowly over weeks/months decia18 ;
lunnan18 ; angus19 ; chen22a ; chen22b . Because of their brightness, they can
be observed out to high redshift (to $z\sim 4$ so far, cooke12 ). Their
explosion mechanisms and progenitor stars are highly debated, and may be
associated with very massive stars, potentially interacting with a surrounding
medium, or with pair-instability SNe heger02 ; gal-yam09  or magnetars
nicholl13 ; inserra17 . The host galaxies of SLSNe have low metallicities and
extreme properties neill11 ; vreeswijk14 ; leloudas15 ; lunnan15 ; perley16 ;
chen17 ; schulze18 . In addition, SLSNe have a potential use for cosmology
inserra13 ; scovacricchi16  out to much higher $z$ than currently done with
SNe Type Ia. Spectra of SLSNe are diverse, and fundamental to assess the
properties of these explosions quimby18 ; gal-yam19b . Current time-domain
survey facilities such as the Zwicky Transient Facility (ZTF, graham19 )
discover about 30 SLSNe per year, with $z<0.7$ chen22a ; chen22b . In the
future, the Legacy Survey of Space and Time (LSST) at the Vera Rubin Telescope
will discover about 1000 SLSNe per year scovacricchi16 , most of which will be
fainter than 20 mag. in the $r$-band and out to higher $z$.

Although SLSNe are much brighter than normal SNe, and have copious UV
emission, they are also observed at typically larger distances, and thus are
faint targets for near-UV spectroscopy. Spectral observations of SLSNe with
8-m telescopes are mostly limited to low-spectral-resolution studies, with a
few notable exceptions where the SLSN host galaxies could be studied in
absorption vreeswijk14 ; yan18 . X-Shooter spectra of one of the most luminous
H-poor SLSNe, vreeswijk14  showed that the SLSN host galaxy has extreme
properties in its absorption lines: they are very narrow, barely resolved at
the X-Shooter resolution, among the narrowest in a larger sample of galaxies
observed in absorption with UVES ledoux06 . In addition, the absorption lines
in SLSN host galaxies (also measured at lower resolutions) are extremely week,
much weaker than what is typically seen in other starforming galaxies such as
GRB hosts. The reason for these extreme properties is not yet understood.
Highly efficient observations in the near-UV regime with intermediate-high
spectral resolution ($R>10,000$) will allow us to shed light on this puzzling
result. Most metal lines (e.g. from Fe ii, Zn ii, Ni ii, etc.) associated with
the ISM of the SLSN host galaxies are in the rest-frame UV, and will be
observable in the ground UV region for targets at $z>0.5$. The study of the
kinematics and relative abundance of metals, regardless of the detection of
the Ly$\alpha$ line, will enable the characterisation of the chemical and
dynamical properties of the ISM of the exotic galaxies hosting SLSNe.

In addition, ISM atoms that are close enough to the SLSN could, in principle,
be excited by the SLSN radiation, and potentially show (variable) fine-
structure lines of excited levels of Fe ii and other ions. This has been often
observed in GRB afterglows (delia09 ; vreeswijk07 ; decia12 ), but never in
SLSNe. The detection of Fe ii fine-structure lines and their variability would
enable the study of the environment surrounding SLSNe, and potentially reveal
the nature of their progenitors. The strongest Fe ii fine-structure lines are
in the UV (around 230, 240, and 260 nm), and are shifted into the ground UV
for targets at $0.15<z<0.7$. The high near-UV efficiency of CUBES compared to
X-Shooter will make it the ideal instrument to observe the Fe ii fine-
structure lines and their variability.

Finally, the emergence of Mg ii ($\lambda\lambda$ 279.6, 280.3 nm) emission
was observed about 100 days after a SLSN from low-resolution spectroscopy
lunnan18b . This was interpreted as a line echo from a shell around the SLSN,
showing how such observations can be very powerful to constrain the nature of
the SLSN progenitors and the physics of the explosion. Thanks to its high
efficiency in the near-UV, CUBES will be ideal to observe other potential Mg
ii echoes from SLSNe at $z<0.4$, with about twice the efficiency of X-Shooter
and at higher spectral resolution. In addition, the spectra of H-poor SLSNe
show quite distinctive broad features at rest-frame 350 to 500 nm, which seem
to be due to a superposition of several individual O ii narrower components.
This may indicate relatively narrow photospheres (in velocity space) that are
nevertheless expanding rapidly, giving potentially new insights on the physics
of the explosion gal-yam19 ; gal-yam19b . While the velocity width of these
narrower components could be of hundreds of km s-1, they have yet to be
resolved with higher-resolution spectroscopy. The spectral coverage of CUBES
will enable a better characterisation of the near-UV spectra of the most
nearby SLSNe ($z<0.1$), and further constrain the physics of these energetic
and hotly debated explosions.

####  5.3.1 T3: Key requirements

  * •

Spectral coverage: Essential: 300–405 nm. Desirable: 300–510 nm. A wider
spectral coverage would enable the measurement of the ISM dust-corrected
metallicity from simultaneous measurements of Ly$\alpha$ and metal lines from
Fe ii and Zn ii, increase the $z$-range of the potential targets, and increase
the variety of metals lines that can be detected.

  * •

Spectral resolution: Essential: $R>7,000$. Goal: $R>10,000$. A higher
resolution would resolve the narrow lines associated with the ISM of the SLSN
host galaxies.

  * •

Operations: Target of Opportunities observations of SLSNe can be non-
disruptive (soft ToOs), given the slow evolution of their light curves. The
Rapid Response Mode is not essential, but advisable in case the fine-structure
lines turn out to vary on short timescales, or for rare discoveries of SLSNe
very early after their explosion.

##  6 Defining cases for the top-level requirements

To inform specification of the TLRs for the instrument, and to quantify their
tolerances, specific cases were identified where they drive a particular
requirement. These are now briefly summarised.

###  6.1 Wavelength coverage: Blue end

Several cases inform the extent of the wavelength coverage at the blue end:

  * •

Cometary Science (Case S1): Good coverage required on both sides of the OH
line at 308 nm; coverage down to 305 nm is essential.

  * •

Neutron-capture elements (Case G7): Specific absorption lines demand coverage
shortwards of 305 nm, namely: Bi I at 302.5 nm and Ge I 303.9 nm. In metal-
poor stars with ‘normal’ (solar-scaled) abundances, we note that Bi I is very
hard to detect, such that it should not drive the wavelength coverage
requirements by itself. (Although there are some extremely metal-poor stars
with an r-process excess, e.g. [Bi/Fe] $=$ $-$1.0 for CS 31082-001 Bar11 ).

  * •

Reionization (Case E4): A general argument for this and other extragalactic
cases is that by going to shorter wavelengths additional parameter space is
opened-up in terms of accessing a broader range of redshifts.

  * •

Kilonovae (Case T2): To some extent we do not know what to expect for these
(and some other) transients at the shortest wavelengths. Nevertheless, the
opacity of lanthanides in the ejecta of kilonovae increases towards the blue.
Maximising the instrument coverage to the atmospheric cut-off would therefore
provide the best opportunities for ‘discovery space’.

###  6.2 Resolving power: High resolution

Initial tests were presented of which elements had (potentially) useful
diagnostics in stellar studies at $R$ $=$ $20,000$ compared to the $R$ $=$
$40,000$ from a 1” slit with UVES ernandes20 . In short, the conclusion was
that for the large majority of the diagnostic lines the S/N was a more
critical aspect than the resolution. Initial tests were also undertaken for
$R$ $=$ $10,000$, but this resolving power is then too low to allow
quantitative measurements of e.g. the Be doublet at 313 nm.

Given the technical design constraints, a significantly larger resolving power
($R$ $>$ 30,000) would require cross-dispersion and blazed gratings, leading
to a large reduction in efficiency. As such, and reinforcing analysis from the
previous Phase A study (e.g. smiljanic14 ), $R$ $=$ 20,000 to 25,000 is an
optimum balance between sufficient spectral resolution for quantitative
analysis and maximum sensitivity.

The baseline design ExA_Zanutta  uses 9K CCDs and neatly satisfies the
requirement that $R$ $>$ 20,000, combined with good spectral sampling ($>$2.3
pixels per resolution element) and with broad wavelength coverage of 300-405
nm.

###  6.3 Resolving power: Low resolution

Some of the cases outlined above require the excellent throughput of CUBES but
spectral resolution is less critical. For example, simulations were performed
of the sky background at the shortest wavelengths in Case E4 to understand how
far we can go to short wavelengths and remain background (rather than
detector) limited. From this case (and others, e.g. S1), it was clear that a
low-resolution (LR) option would provide a powerful complement to the primary
high-resolution (HR) option. The technical team therefore investigated a LR
option in the optical design, by exchanging the slicer unit with a second unit
with physically larger slices (thus also sampling a larger on-sky area).

A couple of different possibilities were investigated using the end-to-end
simulator, with 1′′ and 0.75′′ slits. The former of these was background
limited with either 1 $\times$ 1 or 1 $\times$ 2 binning and was adopted as
the LR option in the slicer design, giving $R$ $\sim$ 7,000 (comparable to
that obtained with a 0.8′′ slit with the UVB arm of X-Shooter).

###  6.4 Position angle

The cometary science case (Case S1) requires the ability to configure the slit
to a position angle chosen by the user, without the complications of
atmospheric dispersion in the data, i.e. an ADC is required such that
observations can be obtained away from the parallactic angle. This enables
observations of different parts of the cometary coma.

Atmospheric dispersion at the CUBES wavelengths quickly becomes significant
when moving away from zenith. A removable ADC for observations near to zenith
could give a marginal 0.2-0.3% gain in end-to-end efficiency (incl. telescope
and atmosphere), but would entail additional mechanisms in the instrument,
increase the operational overhead for daytime calibration, and complicate the
data reduction (in the current design we have a fixed-format science product).
As such, we did not pursue this further.

###  6.5 Exposure times / duty cycle

Case G3 was used to define the required duty-cycle of observations. The case
is to monitor chromospheric activity in the Ca II H&K lines (potentially in
parallel to ESPRESSO observations, obtained with a separate UT). UVES has been
used for this type of monitoring programmes with a cadence of 1 min (45s
exposures with the ultrafast readout mode of the UVES arrays klocova17 ). The
goal is to achieve a similar $\sim$1 min cadence with CUBES (potentially with
shorter 30s exposures given the better efficiency of CUBES).

In general for CUBES programmes we expect exposure times to be relatively long
($>10$ mins), but some bright targets might require observations of only 1 min
or so. For instance, for a V=11 mag G2-type (solar) template at 313 nm (at
airmass $=$ 1.2, 1′′ seeing), we find S/N = 32 from the CUBES ETC in 60s,
compared to $\sim$0.5 hr to achieve the same S/N with UVES.

In short, targets with V$<$11 mag can be done with reasonable UVES exposure
times, but CUBES will still yield substantial efficiency gains in operational
time even for relatively bright targets (and accounting for the overhead of
the telescope preset etc).

###  6.6 Flux-calibration requirements

A few cases require flux calibration of the science spectra to a precision of
10% or better (e.g. S1, G1). Approximate flux calibration is also beneficial
to other cases (e.g. E2, T1, T2).

The cometary cases (S1) are among the most demanding, with a requirement of
10% precision, and a goal of 5%. From recent experience with UVES observations
of cometary targets, use of a spectrophometric standard observed during the
night, combined with the instrument master-response curve, has been sufficient
to achieve such results (assuming weather conditions remain comparable). This
has been used to the shortest wavelengths possible with UVES (and X-Shooter)
and we plan to adopt a similar strategy for CUBES.

###  6.7 Fibre-link to UVES

The primary motivation for a fibre-link to UVES is follow-up of transients
(Cases T1, T2), although note that Case G1 to study accretion and outflows in
YSOs also requires simultaneous observations. For these cases, the data must
be taken at exactly the same time for astrophysical reasons, but many of the
other cases outlined above would also benefit from simultaneous UVES
observations (in the sense of obtaining important complementary or supporting
data more efficiently than observing separately with the two instruments).

In general, at least half of the cases outlined above would benefit
significantly from the proposed UVES fibre-link, bringing a powerful addition
to CUBES, for a relatively small cost and with low risk. The fibre-link
ultimately provides efficiencies by avoiding separate observations:

  * •

Transient/YSO cases: Simultaneous UVES observations are required for these
cases to realise their scientific objectives. In absence of the fibre-link,
separate observations would have to be undertaken with CUBES and UVES at the
same time (implying that CUBES and UVES would have to be mounted on different
UTs). Provision of the fibre-link, and assuming CUBES and UVES are mounted on
the same UT, then means that only one UT is used by these programmes.

  * •

Other cases: Constraints on simultaneity are less critical, but there is a
significant efficiency gain via not needing separate observations for many
programmes.

Of course, these gains only apply if the performance of UVES with the fibre-
link is sufficiently good to deliver the observations required within the time
needed for CUBES observations. Assuming the estimated 40% throughput of the
fibre-link at $>$500 nm ExA_Zanutta , calculations for the YSO case
demonstrated the overall potential gain of CUBES$+$UVES observations compared
to separate observations ExA_Alcala .

##  7 Comparison with other instruments

Adopting the efficiency estimates for the end-to-end performance of CUBES
(from ExA_Zanutta ), we now place its sensitivity in the context of existing
instruments. In what follows we have assumed a telescope efficiency for one of
the UTs of 76%, which are then convolved by the skycalc sky transmission model
used in the ESO ETC (which averages over the year and full nights).

###  7.1 UVES

The end-to-end efficiencies of CUBES and UVES are compared in Fig. 13,
adopting values from the ESO ETC for the blue-arm 346 UVES template. In the
figure we have assumed an airmass of 1, seeing of 0.8′′ and two slit-widths
for UVES (as indicated). The dashed red line shows the potential factor of 1.5
gain in performance that is anticipated if UVES were to be upgraded in the
future.

Note that the UVES curves plotted in Fig. 13 are the peak efficiencies at the
centre of the echelle orders – the mean (or typical) efficiencies away from
the blaze peak are lower. To illustrate this point, Fig. 24 shows a similar
comparison, but now with the addition of the UVES efficiencies at the min/max
of the free spectral range of each echelle order. (In reality, the ‘saw-tooth’
in Fig. 13 is a consequence of only having three points per order; the true
response of each order is much smoother.) The comparison in Fig. 13 also
highlights another advantage of CUBES over higher-resolution (cross-dispersed)
data such as those from UVES, as no blaze correction and order merging will be
required for the final CUBES data product.

As shown in the figures, CUBES represents a genuinely transformative
difference in the end-to-end performance for high-resolution spectroscopy in
the near UV at Paranal. For example, the transmission at 313 nm for CUBES
(plus telescope and atmosphere) is 12.5%, as compared to 0.6-1.3% across the
corresponding UVES echelle order, i.e. factors of 10-20 times better with
CUBES across the order.

Figure 12: Comparison of predicted CUBES efficiency (including telescope and
atmosphere) with those for the central wavelengths of the UVES echelle orders
from the ESO ETC. The dashed red line shows the anticipated gain in
performance ($\times$1.5) from a possible UVES upgrade.

Figure 13: Comparison of predicted CUBES efficiency (including telescope and
atmosphere) with those for the central wavelengths of the UVES echelle orders
from the ESO ETC. The dashed red line shows variations of efficiency across
each of the UVES orders.

###  7.2 Other ESO instruments

In Fig. 14 we present a similar comparison of the end-to-end efficiency of
CUBES with results for the UVB arm of X-Shooter (again, where the curves are
interpolated between the peak efficiencies of the echelle orders). We include
results for a 1.0′′ slit (giving $R$ $\sim$ 5400) and a 0.5′′ slit ($R$ $\sim$
9700). CUBES is factors of three (0.5′′ X-Shooter slit) and two (1.0′′ slit)
better over the 320-340 nm range, while also providing the greater resolving
power required for quantitative spectral analysis (e.g. for the Be doublet).
For completeness, we include the efficiency curve for the 1 UT mode of
ESPRESSO in the overlapping spectral range with CUBES.

We also note possible synergies of CUBES LR observations with the LR-BLUE
spectroscopic configuration of the MAVIS instrument now in development for the
VLT mavis . The LR-BLUE setting will deliver $R$ $=$ 5,900 spectroscopy over
370-720 nm (of an AO-corrected source).

Looking further ahead, the baseline spectral range of the high-resolution
ANDES spectroscopy for the ELT is 400-1800 nm, with a goal to extend its
coverage shortwards to 350 nm hires . ANDES will provide R $\sim$ 100k
spectroscopy of individual sources, but the overall system throughput of the
ELT (with the current mirror coatings) drops-off steeply shortwards of 400 nm
(see the expected limiting magnitudes for ANDES from hires ). The broader UV
coverage and competitive system efficiency of CUBES will mean that, in spite
of the large difference in telescope apertures, it will remain unique well
into ELT operations. Indeed, one can envisage strong synergies between CUBES
and ANDES (e.g. observations of metal-poor stars).

Figure 14: Comparison of predicted CUBES efficiency (including telescope and
atmosphere) with those for the central wavelengths of the X-Shooter UVB
echelle orders and the shortward orders of the 1 UT mode of ESPRESSO from the
ESO ETC.

###  7.3 Other facilities

It is harder to compare efficiencies of non-ESO instruments directly, given
different sites (i.e. local atmospheric transmission) and the sparse data
available for some instruments. We first consider two of the high-resolution
spectrographs on Mauna Kea (Subaru-HDS, Keck-HIRES), then two of the other
blue-sensitive Keck instruments (LRIS, KCWI):

####  7.3.1 Subaru-HDS

The efficiencies for the blue set-up of the Subaru HDS instrument, including
the effects of both telescope and atmosphere, range from $\sim$3 to 7% across
the 310-400 nm range. However, the instrument provides $R$ $\sim$ 100k so
there is significant dispersion. We used the HDS
ETC222https://subarutelescope.org/cgi-bin/hds$\\_$etc.cgi, considering a G0 V
(Pickles) spectral template, for $V$ $=$ 17 mag., seeing $=$ 0.8′′, in a 1 hr
exposure using the ‘Ub’ standard set-up, which delivers $R$ $=$ 90,000 over
298-458 nm. At this high dispersion the S/N at 313 nm (with 4 $\times$ 4 pixel
binning) is $\sim$1.5 so, even with significant further binning to the
resolution of CUBES, the gains will be relatively limited. In contrast, for
the same G0 V Pickles template with $V$ $=$ 17 mag. the CUBES ETC predicts S/N
$\sim$15 in 1 hr.

####  7.3.2 Keck-HIRES

The HIRES instrument at Keck has been immensely productive over the years, and
is well placed to take advantage of higher site at Mauna Kea and larger
collecting area of Keck compared to the VLT. We used the HIRES
ETC333http://etc.ucolick.org/web$\\_$s2n/hires to estimate its performance
compared to CUBES, considering a $V$ $=$ 18 mag. flat input spectrum at
airmass $=$ 1 and seeing $=$ 0.8′′ at 313 nm in a 1 hr exposure:

  * •

HIRES: slit-width $=$ 1.15′′ ($R$ $=$ 37k): S/N $=$ 2.3/pix.

  * •

HIRES: slit-width $=$ 1.70′′ ($R$ $=$ 24k): S/N $=$ 2.7/pix.

  * •

CUBES HR: S/N $=$ 31/pix.

Similar calculations for 302.5 nm (near the Bi I line) yielded S/N $=$ 11/pix
for CUBES, compared with S/N $=$ 0.2 and 0.3/pix for 1.15′′ and 1.7′′ slits
(respectively) with HIRES. A similar gain is also obtained at longer
wavelengths with e.g. S/N $=$ 40/pix for CUBES compared to S/N $\sim$ 4/pix
with HIRES at 340 nm.

Compared with existing high-resolution instruments on Mauna Kea, our above
estimates demonstrate the significant performance gain that CUBES will
provide.

####  7.3.3 Keck-LRIS:

The publicly-available LRIS ETC does not include the two configurations that
are most relevant to be compared to the CUBES LR option, namely the 1200 l/mm
grating (blazed at 340 nm) and the 600 l/mm grating (blazed at 400 nm). With a
0.7′′ slit the 1200 grating provides R $\sim$ 3,600, and the total system
throughput (excluding atmosphere) is reported as $\sim$25% at 350 nm steidel04
. The end-to-end efficiency of CUBES at 350 nm is $\sim$25% (see Fig. 13), but
this also includes the effect of the atmosphere. Taking this into account, and
considering that the Keck telescopes have a larger collecting area than the
UTs, the performance of LRIS and CUBES are expected to be roughly comparable,
but with CUBES LR delivering greater resolving power.

####  7.3.4 Keck-KCWI:

KCWI was designed as a blue-optimised, integral-field spectrograph, with a
range of resolving powers (from $R$ $=$ 1250 up to 20,000) over the 350-560 nm
range kcwi . As with LRIS, it is hard to directly compare the on-sky
performance with CUBES. From comparison of the published KCWI efficiency
curves kcwi , while taking into account the larger collecting area of Keck
(cf. the VLT), CUBES will be very competitive with KCWI at 350-400 nm, with
the important addition that only CUBES covers the 300-350 nm range.

##  8 Summary

We have presented the scientific cases that were developed to motivate the
design and construction of the new CUBES instrument. This broad range of cases
illustrates the strong demand for a new, optimised near-UV spectroscopic
capability at ESO’s VLT, that will provide a unique facility for the community
for many years to come.

The Phase A design of CUBES ExA_Zanutta  provides unprecedented spectroscopic
sensitivity at ground UV (300-405 nm) wavelengths. This will enable
observation of much fainter sources than currently possible, opening-up new
opportunities and parameter space across Solar System, Galactic and
extragalactic targets, including rapid follow-up of several classes of
explosive transients. The cases introduced here included performance estimates
calculated with simulation tools developed during the Phase A study ExA_Genoni
; for more detailed examples of the anticipated observational samples, we
refer to the expanded papers on specific cases as indicated in Table 1.

To bolster its observational capabilities, the CUBES design includes
exchangable imager slicers ExA_Calcines  to provide two spectral resolving
powers: $R$ $\sim$ 24,000 for high-quality quantitative spectroscopy, and $R$
$\sim$ 7,000 to maximise the efficiency for observations of faint targets. The
design also includes a fibre-link to UVES to give the powerful capability of
simultaneous, high-resolution ($R$ $\sim$ 40,000) observations at longer
wavelengths. This will provide critical information for studies of transients
and YSOs, while delivering significant operational savings for many cases
where longer-wavelength observations are also required.

###### Acknowledgements.

A. R. da Silva, R. E. Giribaldi and R. Smiljanic acknowledge support by the
National Science Centre, Poland, under project 2018/31/B/ST9/01469.

## Conflict of interest

The authors declare that they have no conflict of interest.

## Data availability statement