The SDSS-V Local Volume Mapper (LVM): Scientific Motivation and Project Overview

The physical processes shaping galactic disks span orders of magnitude in physical scale. Star formation involves cooling warm atomic gas (≫1 kpc), assembling giant molecular clouds (GMCs; ∼50 pc), and collapsing/fragmenting filaments, clumps, and cores (∼10–0.1 pc) inside GMCs, linking kiloparsec to subparsec scales (e.g., E. Schinnerer et al. 2019, and references therein). The effects of feedback by stellar winds, supernovae (SNe) and supernova remnants (SNRs), photoionization, radiation pressure, and active galactic nucleus propagate from small to large scales, eventually giving rise to galaxy-wide phenomena such as chemical abundance gradients and azimuthal variations, galactic outflows, fountains, and the regulation of the physical conditions of the warm and diffuse ISM. Timescales span from the short freefall and crossing times of clumps and cores, to the long dynamical times of galaxy disks (R. C. J. Kennicutt 1989; B. G. Elmegreen 1993). It is now clear that star formation must be at least partly a feedback-driven self-regulated process (see, e.g., E. C. Ostriker et al. 2010; R. A. Crain & F. van de Voort 2023). Despite its importance, feedback is poorly understood quantitatively (see, e.g., R. S. Klessen & S. C. O. Glover 2016). How strongly, how far reaching, and for how long does it affect its environment? How does this change with stars of different masses, ages, and metal abundances? And how does it depend on the properties of the gas on which it acts (e.g., its clumpiness, ionization state, and kinematics)?

Most observational approaches adopt either the top-down view of 100 pc to kiloparsec-scale maps of external galaxies (e.g., A. K. Leroy et al. 2021; E. Emsellem et al. 2022) or the bottom-up perspective of studying detailed parsec-scale structures of local regions in the MW and the Local Group (e.g., S. Kong et al. 2018; C. H. M. Pabst et al. 2020; R. Orozco-Duarte et al. 2023). However, large and small scales are closely coupled. For example, turbulence is driven at the 100–500 pc scale by various dynamical processes and cascades down to subparsec structures, while energy is injected at small 1–50 pc scales via stellar feedback that propagates to drive large scale bulk motions (even galactic outflows in the most extreme cases) that provide turbulent pressure support to the ISM.

The warm (10³–10⁴ K) ionized phase of the ISM is an ideal tracer to study the effects of feedback at optical wavelengths (see L. M. Haffner et al. 2009, for a review of the low density component). Emission line diagnostics are at the core of measuring the physical condition of the ISM (e.g., densities, temperatures, dynamics, chemical abundances, dust extinction, ionization, and excitation, etc.). On scales of ≈10–50 pc, relatively dense H ii regions reprocess about half of the emitted ionizing stellar radiation, while the rest escapes and gets absorbed by DIG. On scales up to kiloparsecs, momentum, energy, and metals get deposited, dissipated, and distributed in the ISM.

LVM is designed to provide maps of the ISM resolving individual sources of feedback while at the same time covering whole galactic systems. The aim of the LVM is to map both the microscopic and macroscopic processes that regulate the behavior, structure, and longevity of the ionized ISM, starting with clouds in the MW and Local Group and their interface to the diffuse ISM to understanding the emergence of kiloparsec-scale patterns. The overarching goal is to contribute to a synthesized multiwavelength picture of the local and global physics of galactic disks, including molecular, atomic, and ionized gas, the stellar populations, star formation, feedback and resulting gas flows, dust content, element abundances and enrichment, galactic environment, and local and global gravitational potential and kinematics. At the same time, LVM will provide important constraints to models of massive-star evolution and atmospheres.

LVM provides spatially resolved spectroscopic observations of dense ionized nebulae, mainly H ii regions but also planetary nebuale, SNR, and superbubbles carved by the collective winds and SN explosions of massive stars, for reliable instantaneous measurements of their physical conditions, with full galactic coverage to obtain a population sample that enables us to understand the time evolution of similar regions as well as the variability of conditions among coeval regions. Separating the spatial and temporal dependencies in such a way for different physical processes involved in star formation and the injection of feedback is crucial to overcome the inherent degeneracies associated with observing instantaneous snapshots of a dynamic ISM at any given time.

Clearly, optical emission lines trace only the warm ionized phase of the ISM. However, in recent years, a number of large-area surveys of the MW and Magellanic System have been conducted at many wavelengths and increasingly at spatial resolutions very similar to the LVM's. This rich set of data will be invaluable when addressing the physical state of the ISM. In the IR, the MW has been imaged at 33'' by BOLOCAM (J. E. Aguirre et al. 2011), by the Spitzer Space Telescope (E. Churchwell et al. 2009), and by Herschel (S. Molinari et al. 2016). Planck has provided maps of the dust distribution and polarization (Planck Collaboration 2015). Large portions of the galactic plane have been recently imaged in the submillimeter with a 19.2'' beam by F. Schuller et al. (2009), in CO with 48'' beam size by J. M. Jackson et al. (2006), in HI at 1' resolution by J. M. Stil et al. (2006), at 20'' by S. Bihr et al. (2015), and at 2' by N. M. McClure-Griffiths et al. (2005). In the radio continuum, among others, recent surveys were carried out by the Giant Metrewave Radio Telescope at 150 MHz (H. T. Intema et al. 2017) and MeerKAT at 1.3 GHz (S. Goedhart et al. 2024), as well as by the THOR project at the Very Large Array (S. Bihr et al. 2015).

Higher density gas at warm temperatures has been observed with a 2'' beam through the stacking of several radio recombination lines (RL) by the Green Bank Telescope (GBT) in the GBT Diffuse Ionized Gas Survey (M. Luisi et al. 2020). Additionally, the Wisconsin H-Alpha Mapper (WHAM) has an all-sky map of Hα emission and limited sky coverage of other optical emission lines at very high sensitivity and velocity resolution, but poor spatial resolution (a 1° beam; L. M. Haffner et al. 2003; D. Krishnarao 2019). X-ray coverage, tracing the hot ionized gas, was mapped earlier with ROSAT (C. Motch et al. 1991) and more recently expanded with eRosita (P. Predehl et al. 2021).

In the Magellanic Clouds, the Spizter and Herschel Space Telescopes surveyed the LMC and SMC (M. Meixner et al. 2006; K. D. Gordon et al. 2011; M. Meixner et al. 2013) in the IR, while T. Wong et al. (2011) and E. Muller et al. (2010) provide CO maps at 45'' beam size. In HI, S. Kim et al. (2003) provide combined ATCA and Parks data in the LMC at 1' resolution, while S. Stanimirovic et al. (1999) provides data on the SMC with 15 resolution. The SMC has also been observed by N. M. McClure-Griffiths et al. (2018) with ASKAP on 30'' spatial scales. In the optical, Hα maps from WHAM provide a large scale picture of the low density warm ionized gas in the LMC (B. M. Smart et al. 2023), SMC (B. M. Smart et al. 2019), and Magellanic Bridge (K. A. Barger et al. 2013). Higher-resolution narrow-band imaging has also been undertaken for a limited set of strong emission lines (SELs; R. C. Smith & MCELS Team 1999). Finally, with regards to low-metallicity massive stars, the recently completed Hubble Space Telescope Director's Discretionary program, Ultraviolet Legacy Library of Young Stars as Essential Standards (ULLYSES; J. Roman-Duval et al. 2020), and optical follow-up, XShooting-ULLYSES (J. S. Vink et al. 2023), are providing UV to near-IR spectra, stellar and stellar-wind parameters, and ionizing fluxes, which are crucial for some comprehensive LVM studies. These two highly complementary spectroscopic surveys targeted ∼250 stars in the LMC and SMC (mostly) and a couple of stars in NGC 3109 and Sextans A. In the LMC and SMC, the surveys uniformly sample the fundamental astrophysical parameter space of high-mass stars while in the other two galaxies they provide important clues about stars at even lower (one tenth solar) metallicity.

The availability of these diverse data sets, as well as many others not listed here, makes a wide-area optical spectroscopic survey with LVM's scope particularly timely.

We now define the survey to be conducted within the LVM project that drives the design of the hardware and survey infrastructure. Figure 1 shows the planned footprint of the survey and visualizes the spatial resolution achieved in the MW and in the Magellanic System. An overview of the LVM targets is given in Table 1. LVM cannot avoid the trade-off between sky coverage and spatial resolution. A reasonable balance is reached by aiming for resolving the internal structure of individual star-forming H ii regions, separating the gas affected by individual sources of ionization (massive stars) in the MW (<1 pc scales), while reaching sufficient resolution in the Magellanic Clouds to study gradients in physical conditions across H ii regions, and to resolve the largest young stellar clusters (10 pc scales).

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Milky Way Survey. This covers the bulk of the MW disk observable from the southern hemisphere, including the vast majority of known optically selected H ii regions and their interface with the diffuse ionized ISM. Our target line depth will allow us detect SELs to map the metallicity and ionization structure of the ISM over the full surveyed area, characterize the feedback-induced kinematics of gas in and around H ii regions, and study electron temperature fluctuations (via auroral lines) and their impact on chemical abundance measurements for a significant fraction of optically visible H ii regions in the sky. A subparsec resolution while covering kiloparsec to global galactic scales allows for feedback effects to be traced starting from the individual sources through ionization fronts and shocks to global effects such as gas flows and galactic fountains.

Magellanic System Survey. This covers the nearby metal-poor SMC and LMC to a line sensitivity limit allowing for the detection of SELs over the full area and auroral lines in the majority of H ii regions, sampling a lower metal abundance regime than is available in the MW while still resolving the inner structures of H ii regions. At the same time, the measurement of the stellar continuum will enable an unprecedented picture of metal-poor stellar populations. Assuming past average weather conditions, we can reach 100% of D₂₅ in the SMC and LMC within 4 yr. At 10 pc resolution, we will distinguish shocked versus photoionized emission, resolve large scale ionization gradients in star-forming regions, and fully sample all sites of star formation and large scale diffuse gas globally.

Nearby Galaxies Survey. This covers nearby (<20 Mpc) galaxies when other targets are not observable on a best-effort basis. This sample of galaxies provides data on many galaxies with a large angular extent in the nearby Universe that cannot be covered with existing IFUs. It provides a data set bridging the LVM and more traditional extragalactic IFU surveys. Some of these galaxies have metallicities below that of the SMC. They are not part of the core survey and hence do not drive requirements.

The LVM will map the physical conditions (kinematics, temperature, density, and ionization state) of ionized gas in the vicinity of young, high-mass stars and young clusters. Quality stellar spectroscopy is now widely available, providing quantitative estimates of hot star luminosities, effective temperatures, and metallicities, and hence measures of ionizing spectra and wind properties (see, e.g., E. Zari et al. 2021). The combination of gas and stellar data will allow us to study quantitatively how energy and momentum are deposited into the surrounding ISM, driving turbulence, and bulk motions of gas, and heating the gas via photoionization and shocks. Even in cases when individual star spectroscopy is not possible in the Magellanic Clouds, young cluster properties can be derived from color–magnitude diagrams (CMDs; T. Bitsakis et al. 2017), adding an outside, nearly face-on perspective of feedback in a lower-metallicity, HI-dominated environment.

3.1.1. Maps of the Line-of-sight Velocity of Ionized Gas

3.1.2. Maps of Emission Line Velocity Dispersion

3.1.3. Maps of Strong Emission Line Ionization Parameter and Shock Diagnostics

The LVM will produce the ideal data set to solve the "abundance discrepancy problem" currently plaguing the derivation of nebular chemical abundances using different methods (J. García-Rojas & C. Esteban 2007; L. J. Kewley & S. L. Ellison 2008; G. A. Blanc et al. 2015, 2019; J. E. Méndez-Delgado et al. 2023a). It is well established that nebular abundances derived using the direct T_e method, SEL methods calibrated against theoretical photoionization models, and methods based on the direct measurement of metal RL show systematic discrepancies at the 0.2–0.3 dex level and sometimes larger. The inability to measure gas-phase chemical abundances to better than a factor of 2 is a severe limiting factor in our understanding of the chemical evolution of galaxies and the Universe (e.g., Y. Lu et al. 2015). Several origins have been proposed for these discrepancies, including temperature and density inhomogeneities in ionized nebulae (M. Peimbert 1967; J. E. Méndez-Delgado et al. 2023a, 2023b), non-Maxwellian electron temperature distributions (D. C. Nicholls et al. 2012), local metallicity variations (S. Torres-Peimbert et al. 1990; X. -W. Liu et al. 2000), and assumptions about the ionization structure and geometry of ionized nebulae. For bright H ii regions in the MW, the LVM will map both strong nebular metal lines and temperature-sensitive auroral lines, producing spatially resolved maps of chemical abundances using different methods, while simultaneously mapping the geometry and ionization state of the gas. For the same regions, the MWM will produce photospheric stellar abundances of the same elements (O, N, S, Ar, see, e.g., Y. I. Izotov et al. 2006) for several of cospatial young massive stars, which provides a unique reference to evaluate and solve the systematics affecting different abundance diagnostics.

3.2.1. Maps of Abundance and Electron Density Diagnostic SEL Ratios across Nearby Bright MW H ii Regions

3.2.2. Maps of Temperature-sensitive Auroral Lines of Different Ions across Bright Optically Detected MW H ii Regions

The LVM will provide measurements of H ii region chemical abundances across our targets. While the level of spatial detail will be 10–100 times lower than in the MW, the SMC and LMC provide a large dynamic range in metallicity when combined with the MW. Also, the external view of these targets allows us to map the chemical composition of ionized nebulae across their disks, without the challenges that dust extinction and distance determinations impose on MW studies. To do this, we need to measure SEL abundance diagnostics across these targets, and auroral lines where possible. Also, to trace the dispersal of chemical elements into the ISM, we push the survey depth to map these emission lines out to the transition into the WIM, beyond the boundary of the discrete ionized nebulae.

3.3.1. Maps of SEL Abundance Diagnostics across the LMC and SMC

3.3.2. Maps of Auroral Lines across the LMC and SMC

Existing wide-field maps of the MW ionized ISM, like WHAM (L. M. Haffner et al. 2003), or the combined MW Hα map presented in D. P. Finkbeiner (2003), show that faint MW Hα emission from the DIG is present everywhere on the celestial sphere, and that it contains significant substructure. WHAM has a beam of ∼1°; and the D. P. Finkbeiner (2003) map, which combines WHAM data with Virginia Tech Spectral line Survey (B. Dennison et al. 1998) data in the north, and the Southern Hα Sky Survey Atlas (SHASSA; J. E. Gaustad et al. 2001) in the south, resolves between 6' and 1°.

LVM is designed to map large portions of the DIG in the MW for bright emission lines in the full optical range (from [O ii]λ3727 Å to [S iii]λ9532 Å) at <1' scales. Aided by spatial coadding, fainter MW DIG emission lines like HeIλ5876 and [S iii]λ λ 9069,9531 can be detected on larger spatial scales (∼10'–100'). These fainter lines allow better determinations of temperature and constrain the amount of leaky photons from H ii regions needed to photoionize the DIG (e.g., C. G. Hoopes & R. A. M. Walterbos 2003; L. M. Haffner et al. 2009). At the outer-edge of star-forming regions, the nebular lines from low-ionization species reach a surface brightness of the order of the Hα intensity. These are enhanced over canonical values of integrated nebulae because the volume of gas where low-ionization emission is emitted in optically thick H ii regions is significantly smaller than that of the fully ionized gas, making it easier to detect. This transition is characterized by large changes in both the absolute surface brightness and also the relative intensities of emission lines.

Studies of the DIG in the MW and other galaxies (e.g., A. Jones et al. 2017; K. Zhang et al. 2017) revealed a possible need for a heating source other than photoionization from H ii regions to consistently explain the emission lines (e.g., L. M. Haffner et al. 2009). The increase in [N ii]/Hα [S ii]/Hα and [O iii]/Hβ with decreasing Hα surface brightness, while [N ii]/[S ii] remains constant, is difficult to reproduce with models of photoionization from leaky H ii regions alone (C. G. Hoopes & R. A. M. Walterbos 2003; J. E. Barnes et al. 2014; F. Belfiore et al. 2022), although 3D photoionization (K. Wood et al. 2005, 2010, 2013) or self-consistent hydrosimulations with ray tracing (E. Kado-Fong et al. 2020; C.-G. Kim et al. 2023) show it is possible. The spatial resolution and wealth of spectral information of the LVM data will allow us to correlate the observed diffuse structures and their emission line ratios with nearby H ii regions and other possible additional heating sources.

The distinction between DIG and dense H ii regions also delineates the sphere of influence of feedback from stars and massive clusters. The same line ratios that probe DIG ionization reveal leaky H ii regions. On scales of ∼10–50 pc, relatively dense H ii regions surrounding massive compact clusters reprocess much of the emitted ionizing stellar radiation, as evidenced by their large contribution to the integrated Hα luminosity of typical galaxies. These dense nebulae absorb stellar photons and prevent them from heating the large scale diffuse ISM in galaxies. Theoretical studies suggest the barrier is porous (C. S. Howard et al. 2017), which can explain observations indicating high leakage (e.g., K.-I. Seon 2009; L. M. Haffner et al. 2009), which vary with environment or with cluster luminosity (e.g., J. E. Beckman et al. 2000). LVM will spatially resolve the transition from the interiors of star-forming regions into their surrounding ISM. This will provide two critical diagnostics of the escape fraction of UV continuum: the extinction-corrected Hα luminosity and an ionization parameter map (E. W. Pellegrini et al. 2012). These two constraints allow for the determination of the covering fraction of escaping radiation.

3.4.1. Maps of SELs in High Galactic Latitude Fields

3.4.2. Extinction-corrected Hα Luminosities of MW and LMC/SMC H ii Regions

The LVM will be particularly powerful at obtaining integrated spectroscopy of young stellar clusters in the Magellanic Clouds, and of MW globular clusters. The external view of the former and compactness of the later minimize the effects of foreground and background contamination from field stellar populations. Of particular interest are clusters with available CMDs from, e.g., the Hubble Space Telescope (A. Sarajedini et al. 2007; A. P. Milone et al. 2023), the VISCACHA survey (F. F. S. Maia et al. 2019; B. Dias et al. 2020), and others. Many degeneracies associated with the stellar population modeling of mixed stellar populations are broken when analyzing individual stellar clusters, providing further opportunities for refining and calibration of stellar population synthesis (SPS) models (e.g., R. P. Schiavon et al. 2004; R. Asa'd & P. Goudfrooij 2020; R. Asa'd et al. 2022). Furthermore, the presence of massive stars in young clusters allows for the opportunity to constrain models of high-mass stellar evolution and atmospheres by demanding simultaneous and self-consistent modeling of both the integrated spectra and the CMDs of these systems (e.g., C. Leitherer 2005; P. R. T. Coelho et al. 2020).

3.5.1. Integrated Continuum Spectra of Young Clusters in the LMC, SMC

Here, we summarize the final science requirements defined as the most stringent requirements derived from the union of the key science cases described above. Table 1 provides an overview of the depth and resolution requirements by survey target, and Section 5 describes the survey implementation in accordance with the requirements defined here.

Survey footprint. The LVM survey footprint should encompass the bulk of all known D < 3 kpc, optically selected H ii regions in the MW visible from LCO. The footprint should also cover the full disks of the SMC and the LMC out to their optical radius. The surveyed areas should be based on a simple geometric selection function that allows straight-forward comparisons with theoretical models and numerical simulations.

Emission line flux and continuum depth. The survey spectra must reach an emission line 5σ depth of 6 × 10⁻¹⁸ erg s⁻¹ cm⁻² arcsec⁻² at 6563 Å in the MW, and 2 × 10⁻¹⁸ erg s⁻¹ cm⁻² arcsec⁻² at 6563 Å in the Magellanic Clouds. Additionally, the SMC/LMC survey must reach a 5σ V-band continuum depth of 23.0 AB mag arcsec⁻².

Spatial resolution. The LVM data must have a spatial resolution better than 1 pc out to a distance of 3 kpc in the MW (i.e., <69''), and sufficient to resolve typical H ii and young cluster scales of <20 pc out to the ∼63 kpc distance of the SMC (i.e., <65''). The LVM system optical design yields 37 arcsec spaxels yielding a resolution of ∼10 pc in LMC/SMC, and <1 pc out to 6 kpc within the MW.

Spectral resolution. The system spectral resolution is largely set by the design of the adopted DESI spectrographs (R ≈ 4000 at 6563 Å). This resolution is sufficient to satisfy all spectral resolution requirements for all the measurements described above.

Relative flux calibration. We require a relative spectrophotometric calibration accuracy of better than 5% over the wavelength range from (3700–9600 Å), which covers all emission line and continuum features of interest. This allows us to satisfy all relative calibration requirements for all the measurements described above, including subdominant biases in Balmer decrement extinction corrections, electron temperature, and abundance determinations based on strong and auroral line ratios.

Absolute flux calibration. We require an absolute spectrophotometric calibration accuracy of better than 5%, which satisfies the requirements of all the science cases outlined above, and keeps calibration errors as a subdominant source of systematic error when comparing the LVM data to models of nebular photoionization and SPS.

Sky subtraction. The goal for sky subtraction accuracy is Poisson limited performance everywhere. The requirement is to ensure systematic (i.e., non-Poissonian) sky subtraction residuals that are small and random enough as to not significantly bias emission line measurements, and allow for S/N improvements when stacking spaxels. Based on the experience of the SDSS-IV MaNGA survey, we establish this as a requirement to achieve a ratio between sky subtraction residuals and expected Poisson residuals of <1.3 over the full wavelength range and <1.1 outside of bright sky emission lines (D. R. Law et al. 2016).

The LVM-I facility has a nonconventional, innovative design, that is driven by the requirements of the LVM program. In particular, the ultra-wide-field nature of the survey (thousands of square degrees) pushes for an order of magnitude increase in IFU field of view (FOV) with respect to previous similar instruments. Mapping this large area in a reasonable amount of time requires moving from arcminute sized IFUs (e.g., MUSE, MaNGA, PPAK, VIRUS-P) to fields of view spanning a significant fraction of a degree. Physical and budget limitations on the number and size of spectrographs and detectors that can be deployed imposes a limit on the number of spaxels whose spectra can be imaged simultaneously. This limit in the number of spaxels pushes for large arcmin scale spaxels in order to reach the desired ultrawide IFU FOV regime (instead of the arcsecond sized spaxels of the instruments mentioned above). Doing this, while at the same time reaching the ambitious physical spatial resolution scales required by the LVM science program (<1–20 pc), is only possible thanks to the very nearby nature of our targets.

Key to understanding the LVM-I system is the fact that optical fibers allow for a very narrow range of input f-ratio to minimize focal ratio degradation, and that we are interested in observing emission that is extended on scales that are larger than the angular size of the IFU spaxels. The sensitivity to extended emission of the instrument scales as the etendue, AΩ, of the telescope, with A being the telescope aperture and Ω as the telescope FOV. For a fixed f-ratio (as required by the fibers), Ω scales as A⁻¹, and therefore, the etendue and the surface brightness sensitivity of the system is independent of the telescope aperture. Put simply, for an IFU with a fixed number of spaxels, fed at a fixed f-ratio, the gains associated with a larger telescope collecting area are canceled by the losses associated with a smaller FOV and a smaller spaxel solid angle on-sky. The choice of telescope aperture merely selects the desired plate scale of the system. We therefore set the telescope aperture at a value that gives us the desired physical resolution at the distance of our targets rather than going for the largest possible aperture as astronomers might be used to.

A second major design driver for the LVM-I is the fact that our main target, the MW, spans the full sky, requiring unusual solutions to the problem of sky subtraction. Also, isolated spectrophotometric calibration stars of sufficient brightness to dominate the flux in our spaxels (which are much larger than the seeing) have a low on-sky surface density of ∼1 deg⁻² (see Figure 7), which makes it difficult to observe a sufficient number of standards inside the FOV of an individual telescope. These two factors push the design toward having separate dedicated telescopes feeding the same spectrographs. Different telescopes are used to obtain sky spectra at large angular distances from the science target pointings, and to observe spectrophotometric standards beyond the reaches of the science IFU telescope FOV.

We choose an existing replicable spectrograph design (DESI) rather than develop our own, motivated by the need to shorten the project time frame to go on-sky and to keep the project costs within the available SDSS budget. While we are aware that, for the purpose of studying ISM kinematics, a resolution of R > 10,000 would have been desirable, this would have required either compromising on wavelength range (and hence on the science) or developing custom-made cross-dispersed spectrographs, which would have exceeded any realistic budget and timeline given the scale of the survey. We therefore opted for the lower-resolution (R ∼ 4,000) DESI spectrographs, and have put a focus on driving the system design to deliver very good LSF stability. This permits a precise characterization of the LSF for individual science frames, which permits pushing the measurement of line broadening further below the instrumental spectral resolution limit (D. R. Law et al. 2021). LSF stability is also key for minimizing sky subtraction residuals. Experience with fiber-fed IFU spectrographs like MaNGA, PPAK, and VIRUS-P has shown that the main driver of LSF variability during observations is flexure, and the focal ratio degradation effects caused by the movement of optical fibers as the telescope points in different directions. We therefore favor a static focal plane design in which the fibers are always stationary. These high-level design drivers lead us to the following solution.

A detailed description of the LVM telescope system is given in T. M. Herbst et al. (2024). LVM-I uses four separate 16.1 cm aperture refractive telescopes at f/11.42 (A. E. Lanz et al. 2022) delivering a 14 field. The first telescope (called "Sci" for "science") feeds the main science IFU, a second and a third telescope ("SkyE" and "SkyW" for sky East and West respectively) target sky fields with very faint MW emission, and are used to obtain sky spectra (A. Jones 2024, in preparation), and a fourth telescope ("Spec" for "spectrophotometric") targets bright and very isolated stars used to flux calibrate the spectra and to calculate atmospheric telluric-absorption corrections (K. Kreckel et al. 2024).

A major advantage of the LVM is that it is using a stationary fiber system that does not move with the telescopes. In fact, all powered optics in LVM-I are stationary. (see, e.g., N. Drory et al. 2015; D. R. Law et al. 2021; and K. Bundy et al. 2022 for discussions of issues of fiber motion and LSF effects in the context of the MaNGA survey). Our design of choice uses siderostats mounted on commercial PlaneWave L-600 direct drive mounts. Each siderostat (consisting of two flat mirrors) delivers images of the sky to a bench-mounted refractive telescope objective (A. E. Lanz et al. 2022) fitted with K-mirror field derotators. The anatomy of an LVM-I telescope is shown in Figure 3. All four telescopes are identical in their optical design, with the exception of the calibration telescope, which observes individual point sources and therefore does not require a field derotator.

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The LVM-I fiber system consists of four fiber-coupled microlens array IFUs that are bench-mounted at the fixed focal planes of the telescopes (Figure 4). After leaving the four IFU assemblies, the 1944 optical fibers meet in a sorting box (the "sorting hat") where they are reorganized into three separate output bundles (T. Feger et al. 2020). Each of these bundles carries 648 fibers consisting of a mix of science, sky, and spectrophotometric calibration fibers from different telescopes, which terminate in a pseudoslit unit that feeds their light into each of the three LVM-I spectrographs. The fibers are 107 μm core Molex FBP fed at f/3.7. All fibers and powered optics are stationary, which together with the temperature stabilized spectrograph chamber (see below) contributes to minimizing LSF variability.

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The science IFU consists of 1801 fibers fed by a 25 ring hexagonal microlens array (see Figure 4) that images the f/11.42 telescope beam down to f/3.7. To prevent diffraction effects at the edge of the hexagonal lenslets, a chrome mask is used to produce an array of circular 353 diameter apertures that follows the hexagonal tiling pattern of the microlens array, which has a 37'' pitch. This translates into an 83% fill factor after the mask is installed. The FOV of the science IFU is a hexagon with a 302 outer diameter, 154 long sides, and a total on-sky area of 0.165 deg². This half a degree IFU FOV pushes the LVM-I into a new regime of ultra-wide-field integral-field spectroscopy.

The two sky IFUs have the same optical prescription as the science IFU, but their hexagonal microlens arrays only have five rings (61 lenslets). The spectrophotometric calibration IFU is identical to the science IFU, but it is sparsely populated with only 24 fibers, distributed along two rings of spaxels (12 in the 19th ring and 12 in the 24th ring, as seen in Figure 4). These individual calibration fibers are exposed by actuating a rotating shutter in front of the IFU, allowing the LVM-I to observe typically 12 bright spectrophotometric standard stars around the direction of each science field during an exposure.

Each telescope is fitted with a pair of CMOS acquisition and guiding cameras of $26\buildrel{\prime\prime}\over{.}8\times 18\buildrel{\prime\prime}\over{.}4$ FOV (M. Häberle et al. 2022), fed by two pick-off mirrors flanking the central IFU (except for the Spec telescope, which is only fitted with one camera).

The LVM-I images the fiber spectra using three spectrographs, which are identical to the Dark Energy Spectroscopic Instrument spectrographs (S. Perruchot et al. 2018; DESI Collaboration 2022) except for certain modifications that are described in detail by N. P. Konidaris et al. (2020). The main differences between the LVM-I and the DESI units are that in our system the slits are denser, carrying 648 fibers instead of 500, and that we adopt a different design for the detectors and dewars. Unlike DESI, the LVM-I detectors are LN2 cooled and are also designed and fabricated by a different vendor.

The spectrographs have a nested dichroic design that splits the light of the fiber spectra into three wavelength channels: blue (b: 3600–5800 Å), red (r: 5750–7570 Å), and infrared (z: 7520–9800 Å). Along each channel, volume-phase holographic gratings are used to disperse the light and image the fiber spectra through fast f/1.7 cameras onto 4 × 4k CCDs with 15 μm pixels. The detectors are Semiconductor Technology Associates, Inc, 4850 devices, with a 30 μm epitaxial variant used for b and r channels, and the 100 μm deep depletion variant used for the z channel. We employ Archon controllers for readout (G. Bredthauer 2014). The dewars are custom designs manufactured by Universal Cryogenics, Inc. The CCD/dewar system is contributed by the Imaging Technology Lab, University of Arizona. The average spectral resolutions across the b, r, and z channels are R = 2700, 4000, and 4600 respectively.

The spectrographs are housed in a temperature stabilized insulated room with an HVAC system keeping the ambient temperature at 20°C ± 0.1°C. As mentioned above, the detectors are cooled by LN2, which is fed by an automatic filling system connected to an external storage tank.

We divide the MW midplane survey into a high-priority 8° wide band centered on b = 0, surrounded by a lower-priority 18° wide extension. Within the MW, the survey will prioritize survey tiles covering known H ii regions to facilitate early science. However, survey simulations predict near complete coverage of the survey area within the planned survey duration.

The Orion Nebula and the Gum Nebula are two of the closest and best studied nebulae in the MW. Due to their proximity, they extend far from the disk midplane, and we therefore target them individually. We add a circular region centered on R.A. 20642 and decl. −1774 with a radius of 65 to densely sample the core of the Orion Nebula, and a larger region of radius 195 sampled with a fill factor of 1/5 to cover the entire ionized nebula and its interface to the surrounding ISM. The Orion region is of high priority, and within the region, we observe outwards from the center. Similarly to the core of Orion, we observe the shell and the southwestern limb of the Gum Nebula. We add two densely sampled circular regions of radius 55 centered on R.A. 11975, decl. −5090, and R.A. 10953, decl. −4154, respectively.

All observations in the MW have a nominal exposure time of 900 s, with the goal of reaching a line sensitivity of 5σ = 6 × 10⁻¹⁸ erg s⁻¹ cm⁻² arcsec⁻² at 6563 Å. The MW is observed at any lunation, with observing constraints of a distance to the Moon larger than 60°, airmass below 1.75, and Earth shadow height ³³ of >1000 km to lower the impact of geocoronal Hα on the MW signal. Due to the relative brightness of Hα in the Orion Nebula and Gum Nebula, the shadow height constraint is relaxed to >500 km, with all other parameters the same as for the rest of the Galaxy. For the fiducial exposure time, a small number of survey tiles show saturation in the strong nebular lines. We identify these cases and perform repeat short 10 s exposures on them to recover the spectra of the bright saturated features.

Finally, we target a sparse grid of all-sky pointings with a 4° spacing through the out-of-disk ISM to be observed when other targets are not observable given the scheduling constraints. Table 1 lists the survey geometry, and Figure 5 shows the survey area on the sky.

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All observations in the Magellanic System have a nominal exposure time of 8100 s broken up into nine exposures of 900 s each. The exposures are dithered on a nine point grid similar to the one employed in MaNGA (see D. R. Law et al. 2016). The Magellanic System observations aim to reach an emission line depth of 5σ = 2 × 10⁻¹⁸ erg s⁻¹ cm⁻² arcsec⁻² at 6563 Å, and a continuum surface brightness depth at 5σ of 23.0 AB mag arcsec⁻² in the V band.

Unlike the MW survey, we restrict the Magellanic System observations to dark/gray time, with a lunation constraint of <0.25. We relax the airmass constraint to $\sec z\lt 2.0$ and observe at distance from the Moon of 45° or more. Because of the velocity offset of the Magellanic Clouds relative to the Earth, geocoronal emission is less of a concern, and we do not impose a shadow height limit.

The LMC surveys cover the galaxy to roughly its optical radius, D₂₅. We survey a circular region of radius 5° centered on R.A. 7951, decl. −6853. The SMC survey area is comprised of the union of two ellipses to capture the main part of the SMC disk and the star-forming complexes on the southeastern side of the galaxy. The main galaxy area is defined as R.A. 1316, decl. −7280, with semimajor axes a 28, b 16, and position angle (PA) of 45°. The southeastern extension is defined in terms of R.A. 2000, decl. −7320, with semimajor axes a 15, b 10, and PA of 155°. Figure 6 shows an overlay of the survey area in the Magellanic system on top of Hα images, and Table 1 lists the survey geometry.

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The remarkably large FOV of the LVM is designed to efficiently observe our core survey targets; however, it is naturally well suited for other extended nearby galaxies as well. Moving to larger distances, the 353 LVM fiber size samples increasingly larger physical scales. At distances of ∼20 Mpc, this corresponds to 3–4 kpc scales, equivalent to what has been achieved for more distant galaxies in surveys like MaNGA (K. Bundy et al. 2015) and SAMI (S. M. Croom et al. 2021).

LVM enables the unique opportunity to resolve complete nearby galactic systems out to ∼20 Mpc distances, often in a single LVM footprint. Modern IFU instruments remain inefficient when mapping the large apparent sizes (R₂₅ ∼ 5'–10') of such nearby galaxies, with the most ambitious programs (e.g., PHANGS-MUSE; E. Emsellem et al. 2022) typically covering out to only 0.5 R₂₅. With LVM, we are able to spectroscopically map full galaxy disks, with the unique advantage of having multiwavelength data at matched (or significantly better) resolution easily achievable, and often already available in the literature. With ancillary 1'' ∼ 100 pc resolution observations able to isolate individual star-forming regions and ionizing sources, LVM observations will constrain the kiloparsec-scale ionized gas and stellar populations. This enables us to link ancillary observations at small scales with bulk energy injection on large scales, placing direct constraints on how these physical processes relate to global gas inflow and outflow processes. These targets are also ideally suited for matched-scale multiphase observations, linking ionized, atomic, and molecular gas properties. All targets will be observed with dithering to achieve improved image fidelity and recover the complete flux. The addition of these targets to the survey plan also makes effective use of unused telescope time, when our higher-priority core science targets are unavailable.

We have compiled a master list of nearly 500 galaxies within 20 Mpc that are well suited for LVM observations, and will be observed on a best-effort basis over the course of the survey (K. Kreckel 2024, in preparation). The full completion of this target list is not guaranteed, but survey simulations suggest near complete coverage is feasible. Due to the incompleteness of galaxy catalogs out to 20 Mpc distances, and the unique opportunities provided by LVM to observe the closest star-forming dwarf galaxies, we have developed our target list by combining two strategies. The first is tailored to the Local Group, and the second focuses on massive star-forming galaxies out to ∼20 Mpc.

5.3.1. Local Group

5.3.2. Local 20 Mpc Sample

The MW fills the sky. Subtracting the night sky spectrum requires us to find places in the sky where the MW emission is faint enough to be considered zero at our survey's depth. However, night sky lines are also highly variable spatially, leading us to construct a grid of optimal sky fields for any 5°–10° position on the sky. Our sky field identification and selection strategy balances both of these concerns. Figure 7 (left) shows the distribution of our sky field candidates on the sky.

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WHAM darkest sky fields. We use the WHAM data (L. M. Haffner et al. 2003) to identify regions of exceptionally low MW emission, with Hα below 0.5 R ³⁴ (red points in Figure 7, left). This corresponds to 1/3 of the MW survey 5σ depth and is comparable to the depth of the Magellanic System survey. There are only a small number (14) of these, to ensure frequent repeat observations for geocoronal monitoring.

Gridded sky fields. A healpix ³⁵ division of the sky with nside = 8 results in 768 grid locations with a separation of 7.5° between grid centers. For each position, we optimize the exact field center within the healpix to minimize the MW Hα foreground, as measured from WHAM, and the integrated stellar light, as measured from summing individual Gaia sources. This results in maximum separations of 15° between grid locations. When possible, we preferentially select positions corresponding to infrared dark clouds (IRDCs) closer to the midplane of the disk (blue points in Figure 7, left), selected from K. Dobashi (2011). For these IRDCs, we require a large area >71 arcmin² (twice the size of the sky IFU), A_V > 5 mag, and a separation of at least 5' (the sky IFU size) from sources in the Wide-field Infrared Survey Explorer H ii region catalog (L. D. Anderson et al. 2014).

Our initial strategy is to always point one sky telescope to one of the WHAM darkest sky fields, as these are critical for the geocoronal monitoring and subsequent subtraction (see detailed discussion of the sky subtraction strategy in A. Jones 2024, in preparation). However, these will be at larger separations from the science fields. In parallel, we use the other sky telescope to observe the closest gridded sky field to our science tile. During commissioning and early observations, we validate and calibrate the sky fields, particularly the IRDC fields, against the WHAM darkest regions, and remove fields with high MW foreground emission or unacceptable stellar contamination from the catalog.

The LVM spectrophotometric calibration procedure is a direct descendant of the calibration strategy employed for the MaNGA survey, described in detail in R. Yan et al. (2016), which achieved better than 5% spectrophotometry across the MaNGA wavelength range of 3600–10000 Å. As in MaNGA, we rely on F-type stars for spectrophotometry as these are abundant, bright, and have relatively well understood stellar atmospheres (R. L. Kurucz 1992). Unlike MaNGA, we now have GAIA DR3 (Gaia Collaboration 2023) available to us, including BP/RP and RV spectra, to select suitable stars.

To refine the selection of F-type stars, we select based on T_eff from low-res spectroscopy as well as high-res RV spectroscopy. We select bright stars with 5 < G < 9 mag, parallax error less than 20%, no known variability, 6000 K < T_eff < 7500 K, and publicly available BP/RP spectra.

Based on this first selection, we further select only those stars that are isolated enough to dominate the flux across the BP/RP spectrum by at least a factor of 100 in our 35'' diameter spaxel, i.e., the sum of the flux of all other stars down to the GAIA limit is less than 1% of the flux of the bright F-type star.

This final catalog of stars provides us with a sufficient density of bright calibration stars for all science targets (Figure 7, right). A limited number of sources are found directly toward the inner Galaxy midplane, as crowding makes it challenging to find isolated stars, as well as toward the ecliptic poles as they do not yet have BP/RP spectra available due to the Gaia scanning laws and hence small number of visits in DR3 (F. De Angeli et al. 2023). Overall, the median separation between calibration sources in our bright isolated star catalog is a few degrees, and we can find ensembles of 12 calibration stars within 5°–8° of any science field (12 stars is a sufficient number to achieve our calibration goals, as demonstrated by R. Yan et al. 2016 and verified during LVM commissioning).

Spectrophotometric calibration uses the GAIA BP/RP (XP) spectra for the overall calibration, and stellar atmosphere fitting for the high-frequency features as well as for telluric-absorption correction. A detailed description of the calibration procedure and performance is given by K. Kreckel et al. (2024). More than 90% of exposures cycle through 12 standard stars, and 100% of the exposures have at least 10 stars. Given the exposure time of 900 s, each standard gets exposed for about 50 s, with the remainder of the time spent slewing and acquiring the stars. Further optimization of the system will allow for faster acquisitions and hence longer exposures for each standard star (see J. Sanchez-Gallego 2024, in preparation).

The LVM-I began commissioning in 2023 July, and the LVM survey started in 2023 November. As of the writing of this paper, the survey operations are funded through 2027 May. The facility is starting operations with a remote observer present at all times to open and close the enclosure and monitor the weather and the system for any critical issues. Target selection, scheduling, target acquisition, guiding, and cycling through the spectrophotometric standards during an exposure are fully automated and do not require any human intervention. The plan is for the facility to transition to a fully autonomous operation during 2024. An in-depth description of the operations procedures and software is given in J. Sanchez-Gallego (2024, in preparation).

One of our highest priority survey targets is the Orion nebula, where LVM is able to provide remarkable <0.1 pc resolution views of the nearest massive star-forming region (Figure 8; K. Kreckel et al. 2024). Our inside-out survey strategy has already uncovered delicate filamentary ionized gas struc- ture and eroding dense clumps associated with the HII region IC 434. To the east, we observe the iconic Horsehead Nebula, which is associated with the Orion B molecular cloud (L1630) and in the process of being photodisassociated. At each position in this map, we obtain a full optical spectrum, for a total of ∼182,000 spectra. These 101 tiles were observed on 14 different nights over the course of more than 2 months, with varying observing conditions. The remarkable uniformity of this mosaic speaks to the high level of precision we are able to achieve in our flux calibration and astrometry.

Our early survey operations have also marked the start of our coverage of the LMC, and our survey simulations suggest we will achieve full coverage (see Table 1) within the duration of the project. Figure 9 shows 58 tiles, each observed in nine dither positions, to create detailed maps of the filamentary nebular emission that threads the ISM of this galaxy. Individual filaments can be identified at 10 pc scales, and are traced in multiple lines (e.g., [O ii], Hβ, [N ii], Hα, [S ii]). Hα bright photoionized Hii regions can also be identified and resolved, such as the well-studied star-forming complex N44 (M. S. Oey & P. Massey 1995). Our IFU spectroscopy will provide detailed kinematic constraints on the energetic input from the well-studied powering stellar sources (A. F. McLeod et al. 2019) for a comprehensive sample of star-forming regions across this low-metallicity galaxy. The combination of multiple emission lines, along with fainter auroral line detections, will also provide detailed abundances of these Hii regions, allowing us to directly map the diffusion and mixing of metals into the disk. In this way, we will build up a comprehensive picture of where energy and material are deposited into the ISM at the energy injection scale across this entire galaxy.

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The survey was designed from the very beginning with the goal of providing not only the raw and reduced data but also high-order data products that allow the community to make use of the observations easily and homogeneously. We follow the path of most recent IFS surveys such as CALIFA (S. F. Sánchez et al. 2012), SAMI (S. M. Croom et al. 2012), and in particular MaNGA (K. Bundy et al. 2015). All of them have provided high-level data products including observational and/or physical parameters extracted from the observed spectra (S. F. Sánchez et al. 2016a; I. T. Ho et al. 2016; K. B. Westfall et al. 2019).

Following the DRP that produces reduced and calibrated row-stacked spectra (RSS), we are developing a DAP that decomposes the observed spectra into the stellar (and nebular) continuum, and the ionized gas emission lines, and then extracts a set of observational and physical quantities that characterize each component.

Inside the Local Group, and in particular, in the MW, individual LVM fibers do not sample scales large enough to contain a representative sample of the local stellar population. In fact, fibers often contain single stars to few stars, or parts of single star clusters. This requires changing the way in which the stellar continuum is typically modeled in IFS observations of nearby galaxies. For the LVM-DAP, we develop a strategy that we name "resolved stellar populations," described in A. Mejia (2024b, in preparation). We do not model the stellar spectra using linear combinations of single stellar populations, but a library of individual stellar spectra covering a wide range of stellar properties (T_e , log(g), [Fe/H], and [α/Fe]).

The current version of the LVM-DAP (pilot DAP) uses the pyFIT3D algorithms included in the pyPipe3D package (E. A. D. Lacerda et al. 2022). A complete description of the pipeline will be presented in S. Sanchez (2024, in preparation). The pDAP analyzes the LVM-data pointing-by-pointing, using as inputs the spectra and errors provided by the data-reduction pipeline. For each individual spectrum, it performs the following steps:

We first estimate the systemic velocity, v_⋆, velocity dispersion, σ_⋆, and dust attenuation, A_V,⋆, of the stellar component by fitting the continuum with a linear combination of only four stellar spectra that are shifted in velocity, broadened, and dust attenuated, while masking the strongest emission lines. Once these nonlinear parameters are determined, a first model of the stellar component is subtracted, and a set of SELs are fit using Gaussians. We subtract these emission lines, and then again fit the remaining continuum, this time with a larger stellar template library covering a wide range of stellar properties, yielding our model of the stellar spectrum. This stellar spectrum is then subtracted from the original spectrum. The remaining "gas-only" spectrum is then analyzed using a weighted-moment nonparametric procedure to estimate the integrated flux, velocity, velocity dispersion, and EW of a predefined set of 192 emission lines. The parameters derived for the stellar and ionized gas components in each of the previous steps are stored in a set of FITS tables that are packed into a single multiextension FITS-file, and the model spectra are stored in an RSS format. For the details of the algorithms and their implementation, we refer the reader to S. Sanchez (2024, in preparation) and E. A. D. Lacerda et al. (2022). ³⁶

Figure 10 illustrates the pDAP analysis using data in the Rossette Nebula. The best-fit final model generated by the pDAP is shown in the upper panel. The middle panels show a set of zoom-ins highlighting the fits to the strongest emission lines using multi-Gaussian models, including [O II]3726,28, Hδ, Hγ, Hβ, [O III]5007, Hα+[N II]6548,83, [S II]3717,31, [S III]9069, and [S III]9531. These emission lines are frequently used to determine the properties of the ionized gas, including the ionizing source, the dust attenuation, the electron density, the ionization parameter, and the strength of the ionization (e.g., J. M. Vilchez & B. E. J. Pagel 1988; L. J. Kewley et al. 2001). The bottom row comprises two panels illustrating the quality of the stellar template fit (3800–4100 Å and 4800–5100 Å), and four additional ones showing the detection of weak but relevant emission lines: (i) HeI4471 and HeI6678, tracers of the He abundance (e.g., M. Valerdi et al. 2021), (ii) [O II]7318, a tracer of SNR (e.g., R. Cid Fernandes et al. 2021), and (iii) [N II]5755 and [S III]6312 (e.g., R. A. Marino et al. 2013), sensitive to the electron temperature. There are computational and accuracy reasons why not all the emission lines are fit using parametric models, extensively discussed in previous studies (e.g., S. F. Sánchez et al. 2016b; E. A. D. Lacerda et al. 2022; S. F. Sánchez et al. 2022). A typical example of such lines is auroral lines, a particularly interesting set of emission lines for the science goals of the LVM, as they are frequently used to trace the electron temperature of different species (e.g., [N II]5755 and [S III]6312, shown in bottom panels in Figure 10). For these weak lines, the weighted-moment analysis provides reliable estimations of their properties instead. The wavelengths of the full set of emission lines analyzed are shown as vertical orange lines in all panels of Figure 10. We are currently exploring the use of both Gaussian fits and nonparametric estimates in our analysis pipeline.

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Figures 11 and 12 showcase spatially resolved properties of the ionized gas in the LVM data. Each panel shows the spatial distribution of a particular parameter extracted from a seven-pointing mosaic of the Rosette Nebula, 12,607 spectra in total, obtained during early science observations in 2023 October. Figure 11 shows the flux intensity maps for a subset of analyzed emission lines, including the [O II]3728, Hγ, Hβ, [O III]5007, Hα, [N II]6583, [S II]6717, [S II]6731, and [S III]9531. The changes in the ionizing conditions and physical properties of the nebula are clearly seen as different spatial distributions. It is important to highlight that, despite the considerable dynamic range of the fluxes, we detect emission in all the emission lines in this figure at any location within the nebula.

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Figure 12 shows the spatial distribution of some additional properties of the ionized gas emission lines. The first two panels, from top left to bottom right, show the velocity and velocity dispersion derived for Hα. Both figures highlight the accuracy achieved by the LVM data to measure low-velocity patterns. We are measuring velocities with an accuracy of a fraction of 1 km s⁻¹, and velocity dispersions with an accuracy of the order of a few kilometers per second. The remaining panels of this figure show the spatial distribution of a set of line ratios frequently utilized in the analysis of ionized gas. The third panel shows the Hα to Hβ ratio, a tracer of the dust attenuation (e.g., D. E. Osterbrock 1989; V. Wild et al. 2011; S. Salim & D. Narayanan 2020). A ring structure and a shell to the southeast are clearly observed. The fourth and fifth panels show the [O III]/Hβ and [N II]/Hα line ratios, frequently used to trace the ionizing source (for unresolved nebulae) and the changes of the ionizing conditions (for resolved ones). There is a clear and expected anticorrelation between these parameters, with regions of high [O III]/Hβ corresponding to values of low [N II]/Hα. Finally, the sixth panel shows the [S II]6716/[S II]6731 line ratio, a tracer of the varying electron density (e.g., D. E. Osterbrock 1989).

As has been the case for two decades, all SDSS-V data will be made public. LVM is no exception, and the first public data release of LVM raw, reduced, and science-ready data and products will happen with SDSS DR20 in late 2025 and yearly after that, with the full data set becoming public 1 yr after survey operations end.

Among the data we will release are the raw data; fully reduced, sky subtracted, and flux calibrated RSS (RSS files) including fluxes, uncertainties, LSF information, point-spread function information, fiber R.A. and decl., and sky spectra; fluxes of more than 200 emission lines or upper limits; kinematic data for subsets of those emission lines; as well as all code used to reduce and analyze the data, map and cube making routines; and finally all template spectra used in reduction and analysis. Model-dependent quantities, such as densities, temperatures, abundances, fits to nebular emission models, etc., will be made available as value added catalogs.