Introduction
Introduction
Building a census of baryons in galaxy halos is critical for connecting feedback on the smallest, star-forming scales to structure formation on cosmological scales. The stellar population and interstellar gas in external galaxies have been well characterized across various cosmic epochs using photometric and spectroscopic observations (see C. Conroy [30 ] ; L. J. Kewley et al. [74 ] ). Although these data reveal a variety of physical processes driving galaxy formation, stars and the interstellar medium (ISM) constitute less than 25% of the baryonic matter in galaxy halos. The remaining baryons reside in the circumgalactic medium (CGM, see J. Tumlinson et al. [126 ] ; H.-W. Chen & F. S. Zahedy [24 ] ), which is the interface between the star-forming region and the intergalactic medium.
The CGM hosts gas with a vast range of thermodynamic properties. The hot CGM with volume density n H < 10 − 4 c m − 3 n_{\mathrm{H}} < 10^{-4}\ \mathrm{cm}^{-3} n H < 1 0 − 4 cm − 3 T ≳ 10 6 K T \gtrsim10^{6}\ \mathrm{K} T ≳ 1 0 6 K [63 ] ; J. N. Bregman et al. [12 ] ), although insights about lower-mass halos can be obtained through stacking efforts (see e.g., P. Singh et al. [118 ] ; Y. Zhang et al. [141 ] ; P. Popesso et al. [98 ] ; J. C. Siegel et al. [117 ] ).
Absorption line spectroscopy of background quasars (QSOs) remains the most common technique for probing gaseous reservoirs in and around individual galaxies. Absorption line measurements of the CGM typically trace the cool- and warm-hot phases (n H ≈ 10 − 4 n_{\mathrm{H}} \approx10^{-4} n H ≈ 1 0 − 4 0.01 c m − 3 0.01\ \mathrm{cm}^{-3} 0.01 cm − 3 T ≈ 10 4 T \approx10^{4} T ≈ 1 0 4 10 5 K 10^{5}\ \mathrm{K} 1 0 5 K [134 ] ; M. Rauch [107 ] ; H.-W. Chen [19 ] ; J. Tumlinson et al. [126 ] ; S. Muzahid et al. [87 ] ; S. Dutta et al. [39 ] ; H.-W. Chen & F. S. Zahedy [24 ] ). An important feature of circumgalactic gas is its clumpy nature, which is revealed by the presence of multiple kinematic components observed in high-resolution spectra (see e.g., M. Rauch et al. [109 ] ; N. H. M. Crichton et al. 2015; G. C. Rudie et al. [111 , 112 ] ).
Explaining the formation, growth, and survival of these clouds remains a significant challenge (see e.g., H. J. Mo & J. Miralda-Escude 1996; A. H. Maller & J. S. Bullock [83 ] ; J. Schaye et al. [116 ] ; A. Afruni et al. [1 ] ; C.-A. Faucher-Giguère & S. P. Oh [44 ] ). The pressure balance between the cool- and hot-phases is a critical ingredient in theoretical treatments of this problem (see also Y. Faerman & J. K. Werk [41 ] ). In addition to thermal motions, the pressure support in cool gas may be attributed to a variety of non-thermal sources, including turbulence (see e.g., H.-W. Chen et al. 2023), cosmic rays (see e.g., I. S. Butsky et al. [13 ] ; S. Ji et al. 2020), and magnetic fields (see e.g., F. van de Voort et al. 2021).
Existing absorption line data can provide unique insights into the pressure support in circumgalactic clouds. The varying absorption strengths of ionic tracers across kinematic components reflect the ionizing conditions encountered along the QSO line of sight. An estimate of the underlying ionization parameter, or equivalently gas density, can be obtained using simple photoionization models that assume ionization equilibrium with the metagalactic radiation background (see e.g., V. Khaire & R. Srianand [75 ] ; C.-A. Faucher-Giguère [42 ] ). Such models have been successful at predicting the observed ionic column densities within individual absorption components and providing quantitative estimates of the underlying gas density (see e.g., N. H. M. Crichton et al. 2015; F. S. Zahedy et al. [138 ] ). Additionally, the observed line broadening for individual ions parameterized by the Doppler parameter b b b b I 2 = b T , I 2 + b N T 2 b_{\mathrm{I}}^{2} = b_{\mathrm{T,I}}^{2} + b_{\mathrm{NT}}^{2} b I 2 = b T , I 2 + b NT 2 b T , I = 2 k B T / m I b_{\mathrm{T,I}} = \sqrt{2 k_{\mathrm{B}} T/m_{\mathrm{I}}} b T , I = 2 k B T / m I b N T = 2 σ v , N T b_{\mathrm{NT}} = \sqrt{2}\sigma_{\mathrm{v,NT}} b NT = 2 σ v , NT k B k_{\mathrm{B}} k B T T T m I m_{\mathrm{I}} m I σ v , N T \sigma_{\mathrm{v,NT}} σ v , NT [109 ] ; G. C. Rudie et al. [112 ] ; Z. Qu et al. 2022). Recovering the gas density, temperature, and non-thermal broadening for individual absorption components provides direct estimates of the total kinetic pressure within the clumpy CGM.
While a photoionization origin is accepted for gas hosting low- and intermediate-ionization species (e.g., Mg+ {}^{+} + 2 + {}^{2+} 2 + 5 + {}^{5+} 5 +
Email: suyashk@uchicago.edu
sorption line surveys find a diversity in the properties of high-ionization absorbers. While broad absorption (b > 30 k m s − 1 b > 30\ \mathrm{km\ s^{-1}} b > 30 km s − 1 [50 ] ; N. Lehner et al. [81 ] ; B. D. Savage et al. [115 ] ; J. K. Werk et al. [133 ] ; B. D. Oppenheimer et al. [90 ] ), narrower high-ionization components can be accommodated by low-density photoionized gas (see e.g., B. Carswell et al. [17 ] ; C. Thom & H.-W. Chen [123 ] ; J. Stern et al. [119 ] ; F. S. Zahedy et al. [140 ] ; T. J. Cooper et al. [31 ] ). Accordingly, kinematically aligned low- and high-ionization species may indicate large fluctuations in local gas pressure (see e.g., K. J. Haislmaier et al. [57 ] ; Z. Qu et al. 2022; S. Kumar et al. [80 ] ; Sameer et al. [114 ] ). While projection effects can be important (see e.g., M. Ticoras et al. [124 ] ), a larger incidence of components hosting both low- and high-ions in the inner halo suggests that the corresponding gas phases are likely co-spatial (see e.g., Z. Qu et al. 2023). Therefore, investigating the physical conditions underlying kinematically aligned absorption from low- and high-ionization species is critical for understanding the presence of multiphase gas in galaxy halos (see e.g., M. Donahue & G. M. Voit [36 ] ).
Despite the accuracy with which physical conditions in the absorbing gas can be studied, constructing density and pressure profiles for cool gas in galaxies spanning a range of halo masses remains a significant observational challenge for two reasons. First, sufficiently deep galaxy surveys are required to uncover low-mass galaxies in QSO fields (see e.g., H.-W. Chen et al. [27 ] ; S. Muzahid et al. [87 ] ). Second, because galaxies tend to reside in overdense cosmological environments, the survey also needs to be wide for uncovering satellite galaxies associated with the host galaxy of an absorber (see e.g., Z. Qu et al. 2023). The Cosmic Ultraviolet Baryon Survey (CUBS, see H.-W. Chen et al. [28 ] ) is uniquely posed for addressing this challenge. With high-quality far- and near-UV spectra of 15 z ≈ 1 z \approx1 z ≈ 1 Hubble Space Telescope (HST ) and accompanying deep galaxy surveys extending out to projected distances of d p r o j ≲ 1 M p c d_{\mathrm{proj}} \lesssim1\ \mathrm{Mpc} d proj ≲ 1 Mpc
Constraints on the gas density and pressure are presented for absorbing gas around 23 galaxies at z = 0.4 z = 0.4 z = 0.4 R v i r R_{\mathrm{vir}} R vir
tion spectra. Recently, echelle-resolution near-UV data were also acquired for the CUBS QSOs, enabling more robust absorption profile decomposition, subsequently improving constraints on the gas ionization state. Combined with a previously conducted study for passive luminous red galaxies at z = 0.2 z = 0.2 z = 0.2 [25 ] ; F. S. Zahedy et al. [138 ] , this effort results in a characterization of the cool gas density and pressure profile around galaxies with dark matter halo masses of M h ≈ 10 12 M ⊙ M_{\mathrm{h}} \approx10^{12}M_{\odot} M h ≈ 1 0 12 M ⊙ 10 13 M ⊙ 10^{13}M_{\odot} 1 0 13 M ⊙
The paper is organized as follows. Details of the QSO spectroscopy and accompanying galaxy surveys are presented in § 3 . The methodology for absorption line fitting and ionization modeling of cool clumps is described in § 4 . Finally, key results regarding the physical properties of cool gas in galaxy halos are discussed in § 5 . For brevity, details regarding individual systems are deferred to the Appendix. A flat Λ \Lambda Λ Ω m = 0.3 \Omega_{\mathrm{m}} = 0.3 Ω m = 0.3 Ω Λ = 0.7 \Omega_{\Lambda} = 0.7 Ω Λ = 0.7 H 0 = 70 k m s − 1 M p c − 1 H_{0} = 70\ \mathrm{km\ s^{-1}\ Mpc^{-1}} H 0 = 70 km s − 1 Mp c − 1 d p r o j d_{\mathrm{proj}} d proj
Observations and Data
Probing the ionizing conditions and thermal state in the diffuse CGM requires both galaxy survey data and quasar absorption spectra. The available QSO absorption spectra and the accompanying galaxy imaging and spectroscopic data for the study are summarized here.
QSO spectroscopy
The Cosmic Ultraviolet Baryon Survey (CUBS, H.-W. Chen et al. [28 ] ) is designed to map diffuse baryons around galaxies at z ≲ 1 z \lesssim1 z ≲ 1 A B ≲ 18.5 AB \lesssim18.5 A B ≲ 18.5 [54 ] ) on board the HST as part of a Cycle 25 General Observer Program (GO-CUBS, PID=15163, PI: Chen). The data presented in this work are available on the Mikulski Archive for Space Telescopes: doi:10.17909/qasf-4k87 . The COS G130M and G160M spectra cover a wavelength range from λ ≈ 1140 \lambda\approx1140 λ ≈ 1140 ≈ 20 k m s − 1 \approx20\ \mathrm{km\ s^{-1}} ≈ 20 km s − 1 b ≈ 12 k m s − 1 b \approx12\ \mathrm{km\ s^{-1}} b ≈ 12 km s − 1 z < 1 z < 1 z < 1 z > 0.2 z > 0.2 z > 0.2 N ( H I ) N(\mathrm{H\,I}) N ( H I ) [7 ] ), covering transitions from low-ions (e.g., Mg+ ^{+} + + ^{+} + [138 ] ; Z. Qu et al. 2022). The COS and MIKE spectra are co-added using custom software described in H.-W. Chen et al. [25 ] and H.-W. Chen et al. [20 ] ; F. S. Zahedy et al. [139 ] , respectively, and the details of this procedure for CUBS sightlines can be found in H.-W. Chen et al. [28 ] .
Recently, near-UV (NUV) echelle spectra were also obtained using the Space Telescope Imaging Spectrograph (STIS; B. E. Woodgate et al. [135 ] ) with the E230M grating for all CUBS sightlines under a Cycle 31 General Observer Program (GO-CONTACT, PID=17517, PI: Chen). These new STIS NUV spectra provide additional spectral coverage from λ ≈ 2000 \lambda\approx2000 λ ≈ 2000 λ ≈ 2800 \lambda\approx2800 λ ≈ 2800 ≈ 10 \approx10 ≈ 10 − 1 ^{-1} − 1 b b b − 1 ^{-1} − 1 3 + ^{3+} 3 + + ^{+} + 5 + ^{5+} 5 + [80 ] ). NUV observations are complete for 14 out of the 15 sightlines with a mean signal-to-noise per resolution element of S/N ≳ 7 \gtrsim7 ≳ 7
All STIS exposures were processed and extracted using stistools (version 1.4.5) with dark current correction omitted. Applying default dark correction consistently leads to negative background counts, while omitting dark correction leads to consistency between the observed fluctuations in the continuum and the error spectrum. A custom software was developed to co-add multiple exposures to achieve optimal spectral resolution and S/N. A finer bin size of 2 km s− 1 ^{-1} − 1 − 1 ^{-1} − 1 λ 1526 \lambda1526 λ 1526
For each spectral bin, the effective exposure time, defined as the total contribution from all unique exposure-order pairs, is tracked and subsequently used for flux calibration. Uncertainties are derived from the gross counts using two-sided Poisson confidence intervals corresponding to the 68% (1 σ 1\sigma 1 σ [51 ] and propagated through the stacking procedure. To perform flux calibration, the combined net count spectrum is compared with the pipeline-generated flux spectrum to obtain a wavelength-dependent scaling factor based on a running median across the full wavelength range. This scaling function is then applied to the net count rates and associated errors to obtain the final flux spectra for each sightline.
Galaxy survey
The CUBS program design includes a deep galaxy survey along all 15 QSO fields to associate gas properties observed along a sightline to the galaxy host. The survey design is summarized in H.-W. Chen et al. [28 ] . Briefly, the survey includes a shallow and wide component carried out using the IMACS multi-object spectrograph on the Magellan Baade telescope (A. Dressler et al. [37 ] ) to map the large-scale environment at d p r o j ≲ 3 d_{\mathrm{proj}} \lesssim3 d proj ≲ 3 L ∗ L_\ast L ∗ z g a l ≲ 0.8 z_{\mathrm{gal}} \lesssim0.8 z gal ≲ 0.8 0.1 L ∗ 0.1L_\ast 0.1 L ∗ z g a l ≃ 1 z_{\mathrm{gal}} \simeq1 z gal ≃ 1 d p r o j < 300 d_{\mathrm{proj}} < 300 d proj < 300 [4 ] ) is designed to uncover galaxies as faint as 0.01 L ∗ 0.01L_\ast 0.01 L ∗ z g a l ≲ 1 z_{\mathrm{gal}} \lesssim1 z gal ≲ 1 d p r o j < 250 d_{\mathrm{proj}} < 250 d proj < 250
In addition to the galaxy spectroscopy campaign, the CUBS design includes deep optical and near-infrared imaging for all 15 fields. The optical imaging performed using LDSS3C and IMACS covers g , r , i , z g,r,i,z g , r , i , z H H H [93 ] ). The available photometry for spectroscopically confirmed galaxies provides constraints on their spectral energy distributions (SEDs). Specifically, the photometric measurements for individual galaxies are fit using the Bayesian Analysis of Galaxies for Physical Inference and Parameter Estimation code (BAGPIPES; A. C. Carnall et al. [16 ] ) with an adopted initial mass function (IMF) from P. Kroupa [78 ] . The inferred stellar masses (M s t a r M_{\mathrm{star}} M star [34 ] ). The corresponding dark matter halo masses M h M_{\mathrm{h}} M h [77 ] and P. Behroozi et al. [5 ] , which differ primarily in their attribution of intrahalo light to the central galaxy or satellite mergers. Recently, high-resolution images of the CUBS fields were obtained for the F814W band of Advanced Camera for Surveys (ACS; [47 ] ) on the H S T HST H S T H S T HST H S T
The CUBS-midz sample at z g a l = 0.4 z_{\rm gal}=0.4 z gal = 0.4 0.7 0.7 0.7 λ λ 1031 \lambda\lambda1031 λλ 1031 λ λ 1548 \lambda\lambda1548 λλ 1548 [see 72 to be submitted] . The steps for identifying likely neighbors for each galaxy in the CUBS-midz sample are summarized below. An initial galaxy sample is identified based on d p r o j d_{\rm proj} d proj R v i r R_{\rm vir} R vir − 1 ^{-1} − 1
CUBS-midz galaxy systems are listed in Table 1 . For each system, the galaxy host is designated as the member with the smallest R v i r R_{\rm vir} R vir d p r o j / R v i r d_{\rm proj}/R_{\rm vir} d proj / R vir d p r o j < R v i r d_{\rm proj}<R_{\rm vir} d proj < R vir Table 1 follow an increasing order in d p r o j d_{\rm proj} d proj Table 1 . In 11 out of 23 CUBS-midz systems, at least one satellite galaxy can be identified following the above mentioned criteria for neighbor identification. The stellar mass distribution of the entire galaxy sample (including satellites) is shown in the left panel of Figure 1 . Although the stellar mass range for the full sample is M s t a r ≈ 10 8 M ⊙ M_{\rm star}\approx10^{8}M_\odot M star ≈ 1 0 8 M ⊙ 3 × 10 11 M ⊙ 3\times10^{11}M_\odot 3 × 1 0 11 M ⊙ λ λ 3727 \lambda\lambda3727 λλ 3727 β \beta β Table 1 . To identify CGM absorption associated with these galaxies, QSO spectra are searched in a wide window of ± 1000 k m s − 1 \pm1000\ \mathrm{km\ s^{-1}} ± 1000 km s − 1 Table 1 , and the corresponding velocity distribution is presented in the right panel of Figure 1 . All but one of the absorbers are found to be within 500 k m s − 1 500\ \mathrm{km\ s^{-1}} 500 km s − 1 ∣ v a b s − v g a l ∣ > 200 k m s − 1 \lvert v_{\mathrm{abs}} - v_{\mathrm{gal}}\rvert> 200\ \mathrm{km\ s^{-1}} ∣ v abs − v gal ∣ > 200 km s − 1 σ v ≈ 88 k m s − 1 \sigma_v \approx88\ \mathrm{km\ s^{-1}} σ v ≈ 88 km s − 1 Figure 1 ). For a dark matter halo of 5 × 10 11 M ⊙ 5 \times10^{11}\ M_\odot 5 × 1 0 11 M ⊙ σ v i r ≈ 69 k m s − 1 \sigma_{\mathrm{vir}} \approx69\ \mathrm{km\ s^{-1}} σ vir ≈ 69 km s − 1 log M s t a r / M ⊙ > 9 \log M_{\mathrm{star}}/M_\odot> 9 log M star / M ⊙ > 9 ≈ 10 12 M ⊙ \approx10^{12}\ M_\odot ≈ 1 0 12 M ⊙ σ v i r ≈ 92 k m s − 1 \sigma_{\mathrm{vir}} \approx92\ \mathrm{km\ s^{-1}} σ vir ≈ 92 km s − 1 σ v ≈ 96 k m s − 1 \sigma_v \approx96\ \mathrm{km\ s^{-1}} σ v ≈ 96 km s − 1 [66 ] ), in contrast to sub-virial motions noted in massive quiescent halos (see e.g., Y.-H. Huang et al. [65 ] ).
Table 1. Galaxy properties for the CUBS-midz sample. Table 1 continued.
Figure 1. (Left): Stellar mass distributions of the CUBS-midz galaxies. The full sample is shown as open histograms, the most massive galaxy in each group as blue shaded histograms, and the galaxies with the smallest halo-radius-normalized projected distance, d p r o j / R v i r d_{\mathrm{proj}}/R_{\mathrm{vir}} d proj / R vir (Right): Velocity distribution of the absorbing gas, characterized by the strongest H I absorption component associated with the CUBS-midz host galaxies. The distribution is well described by a Gaussian with a centroid of μ ≈ − 3 k m s − 1 \mu\approx-3\ \mathrm{km\ s^{-1}} μ ≈ − 3 km s − 1 σ v ≈ 88 k m s − 1 \sigma_v \approx88\ \mathrm{km\ s^{-1}} σ v ≈ 88 km s − 1 M s t a r > 10 9 M ⊙ M_{\mathrm{star}} > 10^9 M_\odot M star > 1 0 9 M ⊙ § 3.2 for details.
In examining the contribution of correlated halos to observed metal-line absorption through line-of-sight projection, S. H. Ho et al. (2020) report a false-positive rate of 15–30% for Mg II within ∣ v c − v g a l ∣ < 500 k m s − 1 \lvert v_c - v_{\mathrm{gal}}\rvert< 500\ \mathrm{km\ s^{-1}} ∣ v c − v gal ∣ < 500 km s − 1 d p r o j ≈ 180 k p c d_{\mathrm{proj}} \approx180\ \mathrm{kpc} d proj ≈ 180 kpc 10 10 M ⊙ ≲ M s t a r ≲ 10 11 M ⊙ 10^{10} M_\odot\lesssim M_{\mathrm{star}} \lesssim10^{11} M_\odot 1 0 10 M ⊙ ≲ M star ≲ 1 0 11 M ⊙ ± 300 k m s − 1 \pm300\ \mathrm{km\ s^{-1}} ± 300 km s − 1 σ v ≈ 90 k m s − 1 \sigma_v \approx90\ \mathrm{km\ s^{-1}} σ v ≈ 90 km s − 1 Figure 1 ; see also F. S. Zahedy et al. [138 ] ), a velocity regime not explicitly tested in these simulations. Second, metal-line absorption declines steeply with increasing d p r o j d_{\mathrm{proj}} d proj [21 ] ; C. J. Liang & H.-W. Chen [82 ] ; Z. Qu et al. 2023), whereas contamination by correlated halos would tend to dilute such a radial trend. Taken together, these observations suggest that correlated-halo contamination is unlikely to significantly affect the presented results.
Methods
Estimates of ionizing conditions from the observed ionic ratios in the diffuse CGM rely on accurate measurements of column densities from the absorption line profiles. Additionally, the measured line broadening for different metals can be used to constrain thermal and non-thermal motions, providing insights into pressure support in the gas. In this section, the steps for the Voigt profile fitting, photoionization modeling, and line width analysis are presented.
Voigt profile fitting
The line profiles of ionic transitions are described using the minimum number of kinematic components necessary to reproduce the observed spectrum. Each component is modeled using a Voigt function with an associated velocity centroid v c − v g a l v_c-v_{\mathrm{gal}} v c − v gal N c N_c N c b c b_c b c [80 ] . Kinematically matched ions are fit simultaneously to obtain an accurate description of the component structure. COS transitions may require wavelength-dependent velocity offsets of δ v ≲ 7.5 k m s − 1 \delta v \lesssim7.5\ \mathrm{km\ s^{-1}} δ v ≲ 7.5 km s − 1 [70 , 82 , 129 , 96 ] ), but the wavelength calibration can be particularly uncertain at the segment edges, requiring offsets of 10 10 10 15 k m s − 1 15\ \mathrm{km\ s^{-1}} 15 km s − 1 [48 ] ). Each MCMC run consists of 200 walkers seeded around the maximum likelihood solution, with 2500 steps assigned per walker. The median value and uncertainties corresponding to the 68% (1 σ 1\sigma 1 σ 3 σ 3\sigma 3 σ 3 σ 3\sigma 3 σ b b b b = 12 k m s − 1 b=12\ \mathrm{km\ s^{-1}} b = 12 km s − 1 6 k m s − 1 6\ \mathrm{km\ s^{-1}} 6 km s − 1 4 k m s − 1 4\ \mathrm{km\ s^{-1}} 4 km s − 1 b = 30 k m s − 1 b=30\ \mathrm{km\ s^{-1}} b = 30 km s − 1
Photoionization modeling and line width analysis
The ionizing conditions of individual gas clumps are reflected in the relative ionic column density ratios. Photoionization equilibrium (PIE) models assuming ionization balance with the metagalactic background radiation have been shown to reproduce ionic columns measured in circumgalactic absorbers (see e.g., [6 , 23 , 132 , 138 , 114 ] (Z. Qu et al. 2022)). The photoionized origin of these clumps is also corroborated by the observed line widths, which suggest the presence of cool T ∼ 10 4 K T\sim10^4\ \mathrm{K} T ∼ 1 0 4 K [109 , 112 ] (H.-W. Chen et al. 2023)), consistent with the expected thermal balance between photoheating from the extragalactic ultraviolet background (UVB) and radiative cooling (see e.g., [89 , 52 ] ). Note that ionizing radiation escaping from star-forming regions can be important out to ≈ 50 k p c \approx50\ \mathrm{kpc} ≈ 50 kpc [10 ] ). Consequently, UVB-only models may underestimate the gas density for such components (e.g., system s1 in the CUBS-midz sample).
Photoionization modeling and line width analysis for CUBS-midz absorption components is performed using a custom software described in S. Kumar et al. [80 ] . Briefly, a grid of photoionization models is compiled using CLOUDY (version 22.01 G. J. Ferland et al. [46 ] ; M. Chatzikos et al. [18 ] ), assuming plane-parallel geometry and irradiation by the extragalactic ultraviolet background (UVB) from C.-A. Faucher-Giguère [42 ] . For reference, a clump of density n H = 0.001 c m − 3 n_{\mathrm{H}}=0.001\ \mathrm{cm^{-3}} n H = 0.001 c m − 3 z = 0.55 z=0.55 z = 0.55 log U = − 2.48 \log U=-2.48 log U = − 2.48 N ( H I ) N(\mathrm{H\ I}) N ( H I ) n H n_{\mathrm{H}} n H [2 ] ) is assumed during photoionization calculations because the ion fractions depend only weakly on the chemical pattern of the gas (see [79 ] ). However, the likelihood function incorporates departures from the solar pattern when fitting column density measurements from carbon, nitrogen, and iron ions (see S. Kumar et al. 2026b, to be submitted, for further discussion).
Following the ionization analysis, the thermal and non-thermal motions in individual kinematic components can be separated. Line width measurements from different metal ions can be modeled using a single gas temperature T T T b N T b_{\mathrm{NT}} b NT [109 , 112 , 80 ] (Z. Qu et al. 2022)), and a likelihood function comparing the predicted values with measurements from the Voigt profile analysis yields constraints on T T T b N T b_{\rm NT} b NT
As an illustrative example, the photoionization and line width analysis of component c3 at v c − v g a l = − 70 k m s − 1 v_{\rm c}-v_{\rm gal}=-70\ {\rm km\ s}^{-1} v c − v gal = − 70 km s − 1 M s t a r ≈ 6 × 10 10 M ⊙ M_{\rm star}\approx6\times10^{10}\ M_{\odot} M star ≈ 6 × 1 0 10 M ⊙ d p r o j ≈ 66 d_{\rm proj}\approx66 d proj ≈ 66 Figure 2 . The low- and intermediate-ionization species can be reproduced by a high-density phase, denoted by the suffix h, with a hydrogen volume density of n H , c 3 h ≈ 0.003 c m − 3 n_{\rm H,c3h}\approx0.003\ {\rm cm}^{-3} n H , c3h ≈ 0.003 cm − 3 Figure 2 ). This phase is constrained using stringent 3 σ 3\sigma 3 σ T c 3 h ≈ 2 × 10 4 T_{\rm c3h}\approx2\times10^{4} T c3h ≈ 2 × 1 0 4 b N T , c 3 h ≈ 6 k m s − 1 b_{\rm NT,c3h}\approx6\ {\rm km\ s}^{-1} b NT , c3h ≈ 6 km s − 1
Figure 2. Column density measurements and detection limits in component c3 (at v c − v g a l = − 70 k m s − 1 v_{\rm c}-v_{\rm gal}=-70\ {\rm km\ s}^{-1} v c − v gal = − 70 km s − 1 top-left panel. A two-phase photoionization model is necessary to reproduce both low- and high-ions detected in this absorption component. The high-density phase (n c 3 h ≈ 0.003 c m − 3 n_{\rm c3h}\approx0.003\ {\rm cm}^{-3} n c3h ≈ 0.003 cm − 3 n c 3 l ≈ 10 − 4 c m − 3 n_{\rm c3l}\approx10^{-4}\ {\rm cm}^{-3} n c3l ≈ 1 0 − 4 cm − 3 bottom-left panel. Line width measurements for various ionic species in the component are included in the right panel. The distinct physical origins of low- and high-ionization species are apparent from their significantly different line widths.
While the absorption from low- and intermediate-species is well-explained by the high-density phase, it underpredicts the amount of associated high-ionization species observed in this component. An additional low-density phase, denoted by the suffix l, with n H , c 3 l ≈ 10 − 4 c m − 3 n_{\rm H,c3l}\approx10^{-4}\ {\rm cm}^{-3} n H , c3l ≈ 1 0 − 4 cm − 3 Figure 2 ) is necessary to explain the full range of ionic absorption observed for this component, particularly the observed O VI/O IV ratio. The associated H I content for the low-density phase is uncertain because most of the measured H I is expected to arise in the high-density phase. Among absorption components that are well-described by a single-phase solution, those with 3 × 10 15 c m − 2 < N ( H I ) < 3 × 10 18 c m − 3 3\times10^{15}\ {\rm cm}^{-2}<N({\rm H\ I})<3\times10^{18}\ {\rm cm}^{-3} 3 × 1 0 15 cm − 2 < N ( H I ) < 3 × 1 0 18 cm − 3 n H ≈ 0.001 c m − 3 n_{\rm H}\approx0.001\ {\rm cm}^{-3} n H ≈ 0.001 cm − 3 0.03 c m − 3 0.03\ {\rm cm}^{-3} 0.03 cm − 3 N ( H I ) < 3 × 10 15 c m − 3 N({\rm H\ I})<3\times10^{15}\ {\rm cm}^{-3} N ( H I ) < 3 × 1 0 15 cm − 3 n H < 0.001 c m − 3 n_{\rm H}<0.001\ {\rm cm}^{-3} n H < 0.001 cm − 3 N ( H I ) ≈ 10 16 c m − 2 N({\rm H\ I})\approx10^{16}\ {\rm cm}^{-2} N ( H I ) ≈ 1 0 16 cm − 2 l c 3 h ≈ 40 l_{\rm c3h}\approx40 l c3h ≈ 40 3 σ 3\sigma 3 σ N ( H I ) = 4 × 10 13 N({\rm H\ I})=4\times10^{13} N ( H I ) = 4 × 1 0 13 5 × 10 15 c m − 2 5\times10^{15}\ {\rm cm}^{-2} 5 × 1 0 15 cm − 2 N ( H I ) N({\rm H\ I}) N ( H I ) [80 ] ). The two-phase structure is also reflected in the measured absorption line widths, where low-ionization species exhibit comparatively narrower profiles than the high-ions (right panel of Figure 2, see also, T. J. Cooper et al. [31 ] ), although the thermal and non-thermal motions cannot be separated in the low-density phase due to similar masses of C IV and O IV.
Ionization modeling and line width analysis results for the remaining CUBS-midz components are presented in Table 2 . Ionization modeling constraints are available for 31 CUBS-midz components identified near 12 unique CUBS-midz galaxies. The median values and 1 σ 1\sigma 1 σ 3 σ 3\sigma 3 σ
Table 2. Summary of absorption properties for CUBS-midz galaxy samplea ^{a} a
Physical conditions of the CGM at z0.4–0.7
Simultaneous wavelength coverage of multiple ionic species in QSO absorption systems provides powerful diagnostics of the ionization conditions in the absorbing gas. Comparing observed ionic ratios in individual absorption components with predictions from photoionization models yields constraints on the gas ionization parameter U U U n H n_{\mathrm{H}} n H b T b_{\mathrm{T}} b T b N T b_{\mathrm{NT}} b NT § 4.2 and are summarized in Table 2 .
Derived constraints on the thermodynamic state and internal motions within cool clumps provide a unique opportunity to examine the physical processes shaping the CGM at intermediate redshifts. Combining absorption-line measurements and ionization modeling results for individual components with the available galaxy data, the physical origin of the absorbing gas and its connection to the host galaxies is now investigated.
A joint CUBS-midz and literature galaxy sample
CUBS-midz galaxies are typically star-forming, with typical halo masses between M h ≈ 10 11 M_{\mathrm{h}} \approx10^{11} M h ≈ 1 0 11 10 12 M ⊙ 10^{12}\mathrm{M}_{\odot} 1 0 12 M ⊙ M h ≈ 10 12 M_{\mathrm{h}} \approx10^{12} M h ≈ 1 0 12 10 13.5 M ⊙ 10^{13.5}\mathrm{M}_{\odot} 1 0 13.5 M ⊙ z = 0.2 z = 0.2 z = 0.2 0.6 0.6 0.6 [25 ] and are reported in Table 3 . F. S. Zahedy et al. [138 ] assumed an updated version of the UVB from F. Haardt & P. Madau [55 ] for the photoionization analysis of COS-LRG absorbers. Since the CUBS-midz sample is analyzed using the FG20 UVB (§ 4.2 ), the photoionization analysis for COS-LRG components is re-performed using FG20 to enable comparison with CUBS-midz components. On average, inferred gas densities assuming the FG20 UVB are found to be lower by 0.2 dex compared to the HM05 estimates. Another important feature of the CUBS-midz sample is the spectroscopic identification of satellite galaxies (§ 2.2). For COS-LRGs, companion galaxies are crossmatched within the Sloan Digital Sky Survey (SDSS, D. G. York et al. [137 ] ), Data Release 1 of the Dark Energy Spectroscopic Instrument (DESI DR1, DESI Collaboration et al. [35 ] ), and the COS-Halos spectroscopic survey (J. K. Werk et al. [131 ] ). Due to the limited survey depth in several of these fields, galaxies with photometric redshifts z p h o t z_{\mathrm{phot}} z phot σ ( z p h o t ) < 0.1 \sigma(z_{\mathrm{phot}}) < 0.1 σ ( z phot ) < 0.1
Table 3. Ionization modeling results for COS-LRG absorbers
Thermodynamic properties of the clumpy CGM
The photoionization modeling outlined in § 4 results in constraints on the clump density n H n_{\mathrm{H}} n H T T T § 5.1 is shown in the left panel of Figure 3 . Components are grouped according to whether their galaxy host is star-forming or quiescent based on the optical spectrum as described in § 3.2 . As a reference, the expected relation between the gas density and temperature assuming photoionization equilibrium (PIE) evaluated using CLOUDY is also shown. The PIE temperature is defined by the balance between UVB heating and radiative cooling in the gas.
Figure 3. The density-temperature distribution of absorption components in CUBS-midz galaxies (z = 0.4 z = 0.4 z = 0.4 0.7 0.7 0.7 z = 0.2 z = 0.2 z = 0.2 0.6 0.6 0.6 [27 ] ; F. S. Zahedy et al. [138 ] ) are presented in the left panel. The panel includes 25 components identified near nine star-forming galaxies (blue stars) and 26 components near 12 quiescent galaxies (red circles). Clumps around galaxy hosts with known neighbors are indicated with open markers. The expected relation from photoionization equilibrium (PIE) adopting the FG20 UVB at z = 0.55 z = 0.55 z = 0.55 right panel. For components without a constraint on the non-thermal broadening, corresponding 99.7% (3 σ 3\sigma 3 σ
Due to the enhancement of cooling efficiency with metal content in the gas (see e.g., O. Gnat & A. Sternberg [53 ] ; B. D. Oppenheimer & J. Schaye [89 ] ), metal-rich clumps are expected to be cooler at a fixed gas density. Curves for the PIE density-temperature relation are shown in Figure 3 for gas metallicities of 0.01 Z ⊙ 0.01Z_{\odot} 0.01 Z ⊙ Z ⊙ Z_{\odot} Z ⊙ [138 ] , S. Kumar et al. 2026b, to be submitted). The measured gas temperatures for most components are broadly consistent with the PIE relation within the measurement uncertainties, although the observed scatter at a fixed density is larger compared to the expectation given variations in clump metallicity. This may reflect fluctuations in the ionizing background radiation, both in terms of its amplitude and spectral slope. particular, significant radiation flux from local sources in the EUV to soft X-ray regime can substantially reduce gas cooling rates by altering the ionization states of the dominant coolants, while also providing additional heating (see, e.g., S. Cantalupo [15 ] ; J. K. Werk et al. [133 ] ; Z. Qu et al. 2023; C. Cadiou et al. [14 ] ). These effects could explain why some systems exhibit temperatures higher than expected under PIE with the UVB alone, suggesting that local radiation sources may need to be included to fully model their thermal and ionization states. A detailed treatment of these effects is deferred to future studies.
The observed line broadening in metal ions cannot be explained with thermal motions alone, indicating the presence of non-thermal motions within the gas. Constraints for the non-thermal motions, parametrized by b N T b_{\mathrm{NT}} b NT Figure 3 , with the observed range between 2–30 km s− 1 ^{-1} − 1 Figure 3 that hints towards the turbulent nature of the gas is the apparent lack of points with low b N T b_{\mathrm{NT}} b NT 3 σ 3\sigma 3 σ b N T ≲ 10 b_{\mathrm{NT}} \lesssim10 b NT ≲ 10 − 1 ^{-1} − 1 n H ≈ 10 − 4 n_{\mathrm{H}} \approx10^{-4} n H ≈ 1 0 − 4 − 3 ^{-3} − 3 b N T ≳ 20 b_{\mathrm{NT}} \gtrsim20 b NT ≳ 20 − 1 ^{-1} − 1 l ≳ 1 l \gtrsim1 l ≳ 1 [80 ] ), compared to the characteristic l ∼ 100 l \sim100 l ∼ 100 [140 ] ). Therefore, the observed n H n_{\mathrm{H}} n H b N T b_{\mathrm{NT}} b NT [108 ] ).
To investigate the role of turbulent pressure support (and equivalently heating) in the CGM, contributions from thermal and non-thermal motions can be evaluated within individual clumps as follows –
P T = ( ρ μ m H ) k B T , P N T = ρ σ N T 2 , (1) P_{\mathrm{T}} = \left(\frac{\rho}{\mu m_{\mathrm{H}}}\right) k_{\mathrm{B}}T,\quad P_{\mathrm{NT}} = \rho\sigma_{\mathrm{NT}}^2, \tag*{(1)} P T = ( μ m H ρ ) k B T , P NT = ρ σ NT 2 , (1) where ρ = 1.34 m H n H \rho= 1.34m_{\mathrm{H}}n_{\mathrm{H}} ρ = 1.34 m H n H μ = 0.6 \mu= 0.6 μ = 0.6 σ N T = b N T / 2 \sigma_{\mathrm{NT}} = b_{\mathrm{NT}}/\sqrt{2} σ NT = b NT / 2 M ≡ 3 σ N T / c s \mathcal{M} \equiv\sqrt{3}\sigma_{\mathrm{NT}}/c_{\mathrm{s}} M ≡ 3 σ NT / c s
P N T P T = m H σ N T 2 k B T / μ = γ M 2 3 (2) \frac{P_{\mathrm{NT}}}{P_{\mathrm{T}}} = \frac{m_{\mathrm{H}}\sigma_{\mathrm{NT}}^2}{k_{\mathrm{B}}T/\mu} = \frac{\gamma\mathcal{M}^2}{3} \tag*{(2)} P T P NT = k B T / μ m H σ NT 2 = 3 γ M 2 (2) where c s = γ k B T / μ m H c_{\mathrm{s}} = \sqrt{\gamma k_{\mathrm{B}}T/\mu m_{\mathrm{H}}} c s = γ k B T / μ m H γ = 5 / 3 \gamma= 5/3 γ = 5/3
The relative contributions of thermal and non-thermal pressure are now compared. To account for upper limits in P N T P_{\mathrm{NT}} P NT [26 ] is adopted to construct the cumulative fraction of cool clumps with P N T / P T P_{\mathrm{NT}}/P_{\mathrm{T}} P NT / P T [45 ] . Separate CDFs are presented for isolated and non-isolated galaxies to investigate the role of galaxy environment on the thermodynamic state of the CGM. To gauge the underlying uncertainty in the CDFs, a bootstrap method is used, resampling each distribution with replacement. The resulting 1 σ 1 \sigma 1 σ Figure 4 . A key conclusion from Figure 4 is that most clumps exhibit subsonic turbulence across different galaxy environments, with ≈ 20 \approx20 ≈ 20 35 % 35\% 35% M > 1 \mathcal{M} > 1 M > 1 z ≈ 1 z \approx1 z ≈ 1 z ≈ 2 z \approx2 z ≈ 2 3 3 3 [112 ] . Isolated quiescent halos also appear to host more turbulent clumps than corresponding star-forming systems. This may be due to the presence of a hot halo in massive ellipticals (Z. Qu et al. 2022, see also § 5.3 ). Nevertheless, the median pressure ratio in both isolated star-forming (corresponding to M i s o , S F = 0.5 ± 0.1 \mathcal{M}_{\mathrm{iso,SF}} = 0.5 \pm0.1 M iso , SF = 0.5 ± 0.1 M i s o , Q = 0.8 ± 0.3 \mathcal{M}_{\mathrm{iso,Q}} = 0.8 \pm0.3 M iso , Q = 0.8 ± 0.3
Figure 4. Cumulative distribution functions (CDFs) for the non-thermal to thermal pressure ratio in cool clumps identified near star-forming and quiescent galaxies are shown in these panels. Measurements in the left panel correspond to 14 absorption components in five isolated star-forming galaxies (blue) and 11 components in four isolated quiescent galaxies (red). CDFs in the right panel are constructed with 11 clumps in the vicinity of four non-isolated star-forming galaxies and 15 clumps associated with eight non-isolated quiescent galaxies. The 1 σ \sigma σ (2) are indicated on the top axis.
Cool gas density and pressure profiles
Spatial variations in cool gas density derived in § 4.2 and total gas pressure derived in § 5.2 are now investigated. The cool gas density profiles for the combined galaxy sample are presented in the top panels of Figure 5 . For non-isolated systems, the projected distance depicted is for the galaxy host at minimum d p r o j / R v i r d_{\mathrm{proj}}/R_{\mathrm{vir}} d proj / R vir M h ≲ 10 12 M ⊙ M_{\mathrm{h}} \lesssim10^{12}\ M_{\odot} M h ≲ 1 0 12 M ⊙ Table 2 ) are omitted here. Large density variations of up to 2 dex are observed within individual halos. Despite the large scatter, there is a striking dependence of the cool gas density on the halo mass at fixed projected distances. This amplitude is distinctly higher for COS-LRGs than for CUBS-midz galaxies. Even within the CUBS-midz sample, massive star-forming galaxies (squares) exhibit denser clumps than lower-mass systems (diamonds). This trend is consistent with the expectations that these cool CGM clouds are pressure confined in a hot ambient medium which reduces in density and pressure support with decreasing halo mass (see e.g., F. S. Zahedy et al. [138 ] ; Z. Qu et al. 2023).
Figure 5. The cool gas density profiles around CUBS-midz galaxies and COS-LRGs are shown in the top panels. Different symbols indicate distinct stellar mass ranges for the galaxy host. Multiple cool clumps resolved in individual halos are connected by a vertical dotted line. Non-isolated systems are indicated using open markers. The hot gas density profile from P. Singh et al. [118 ] for typical halo masses in both samples is also shown. A scaled-up density profile from the hot halo, 100 × n e h o t 100 \times n_{\mathrm{e}}^{\mathrm{hot}} 100 × n e hot T c o o l ≈ 10 4 K T_{\mathrm{cool}} \approx10^{4}\ \mathrm{K} T cool ≈ 1 0 4 K T h o t ≈ 10 6 K T_{\mathrm{hot}} \approx10^{6}\ \mathrm{K} T hot ≈ 1 0 6 K (1) , is shown in the bottom panels. The indicated hot gas pressure profiles adopt the density profiles shown in the top panels with a fiducial gas temperature of T h o t = 10 6 K T_{\mathrm{hot}} = 10^{6}\ \mathrm{K} T hot = 1 0 6 K
To further investigate this scenario, density profiles for hot gas are included in Figure 5 . Key observational probes of the hot CGM are the X-ray luminosity and thermal Sunyaev-Zel’dovich (tSZ) signal. Direct measurements for these probes have only been made in a small number of external massive galaxies, although stacking efforts can provide insights for lower-mass halos. P. Singh et al. [118 ] jointly model stacked X-ray and tSZ observations from galaxies at z ≈ 0.1 z \approx0.1 z ≈ 0.1 M h ≈ 10 12 M_{\mathrm{h}} \approx10^{12} M h ≈ 1 0 12 10 13 M ⊙ 10^{13}\mathrm{M}_{\odot} 1 0 13 M ⊙ n e h o t ( r ) ∝ r − 3 β n_{\mathrm{e}}^{\mathrm{hot}}(r) \propto r^{-3\beta} n e hot ( r ) ∝ r − 3 β r r r β = 0.4 \beta= 0.4 β = 0.4
The corresponding density profile for both halo mass regimes are depicted in Figure 5 . Note that the hot gas profile is evaluated at r = d p r o j r = d_{\mathrm{proj}} r = d proj T c o o l ∼ 10 4 K T_{\mathrm{cool}} \sim10^{4}\ \mathrm{K} T cool ∼ 1 0 4 K Figure 3 ) and the hot medium, expected to be at T h o t ∼ 10 6 K T_{\mathrm{hot}} \sim10^{6}\ \mathrm{K} T hot ∼ 1 0 6 K [83 ] ). To evalute pressure balance between the cool and hot phases, an elevated hot halo profile corresponding to 100 × n e h o t 100 \times n_{\mathrm{e}}^{\mathrm{hot}} 100 × n e hot Figure 5 . There is reasonable agreement between the elevated hot halo profile and the densest clumps along massive CUBS-midz galaxies (square symbols) and COS-LRGs. This would suggest that the densest clumps along the sightline are situated at a physical distance of r ≈ d p r o j r \approx d_{\mathrm{proj}} r ≈ d proj [138 ] ; Z. Qu et al. 2023). Note that cool clumps in lower-mass halos (diamonds in Figure 5 ) exhibit systematically lower densities compared to more massive systems (see e.g., D. R. Piacitelli et al. [94 ] ; A. Polzin et al. [97 ] ). The inferred density of n H ∼ 10 − 4 c m − 3 n_{\mathrm{H}} \sim10^{-4}\,\mathrm{cm}^{-3} n H ∼ 1 0 − 4 cm − 3 [72 ] , N. Mishra et al. 2026, to be submitted). This may reflect reduced pressure support in the hot CGM of low-mass galaxies, although the hot phase is not expected to be stable in such halos (see e.g., J. Stern et al. [120 ] ; A. Kakoly et al. [73 ] ) and observational constraints are challenging to obtain due to low X-ray emissivity (see e.g., J. Ott et al. [91 ] ; L. He et al. 2025).
If the densest clumps along individual sightlines are indeed situated at r ≈ d p r o j r \approx d_{\mathrm{proj}} r ≈ d proj r > d p r o j r > d_{\mathrm{proj}} r > d proj n H ≲ 0.001 c m − 3 n_{\mathrm{H}} \lesssim0.001\,\mathrm{cm}^{-3} n H ≲ 0.001 cm − 3 r ≳ 300 k p c r \gtrsim300\,\mathrm{kpc} r ≳ 300 kpc Figure 5 . However, the presence of low-density components is noted in isolated and non-isolated systems alike. To enable a more precise comparison between cool and hot gas pressures, the total pressure constraints derived in § 5.3 for individual cool clumps are presented in the bottom panels of Figure 5 . Clumps without a constraint on the non-thermal broadening b N T b_{\mathrm{NT}} b NT P T h o t ≡ ρ h o t k B T h o t / μ P_{\mathrm{T}}^{\mathrm{hot}} \equiv\rho_{\mathrm{hot}} k_{\mathrm{B}} T_{\mathrm{hot}}/\mu P T hot ≡ ρ hot k B T hot / μ ρ h o t = 1.34 m H n H h o t \rho_{\mathrm{hot}} = 1.34m_{\mathrm{H}}n_{\mathrm{H}}^{\mathrm{hot}} ρ hot = 1.34 m H n H hot μ = 0.6 \mu= 0.6 μ = 0.6 T h o t = 10 6 K T_{\mathrm{hot}} = 10^{6}\,\mathrm{K} T hot = 1 0 6 K Figure 5 .
Reasonable agreement is found between cool clump pressure and the expected hot gas profile for both CUBS-midz and LRG samples, supporting the need of accounting for warmer gas temperatures and turbulence when considering pressure support in diffuse components (see Figure 3 ). Nevertheless, a large scatter around the expected hot gas pressure profile is observed. It is worth acknowledging that the adopted hot gas temperatures bear significant uncertainties. Direct measurements of hot gas temperature are restricted to nearby galaxies, typically within the inner 20 kpc (see e.g., E. J. Hodges-Kluck et al. [63 ] ), with limited constraints for the outer halo. In the Milky Way, O VII and O VIII emission profiles also suggest the presence of temperature gradients (see e.g., M. J. Miller & J. N. Bregman [85 ] ). Finally, contaminating signals from satellite galaxies and unresolved sources may further skew the hot gas density and temperature measurements (see e.g., Y. Zhang et al. [141 ] ; L. He & Z. Li [59 ] ). Accounting for these systematic effects in the hot gas density may be able to accommodate higher temperatures for the hot gas.
Despite the detailed accounting of the pressure budget in individual absorption components, there are still clumps that exhibit a lower pressure support than the hot halo expectation at r = d p r o j r = d_{\mathrm{proj}} r = d proj Figure 5 . Escaping ionizing flux from local stellar populations can be significant out to r ≲ 50 k p c r \lesssim50\,\mathrm{kpc} r ≲ 50 kpc [22 ] ; J. Bland-Hawthorn et al. [10 ] ), causing UVB-only models to underestimate gas pressure (§ 4.2 ). However, low-density clumps with n H < 0.001 c m − 3 n_{\mathrm{H}} < 0.001\,\mathrm{cm}^{-3} n H < 0.001 cm − 3 d p r o j > 50 k p c d_{\mathrm{proj}} > 50\,\mathrm{kpc} d proj > 50 kpc d p r o j − 2 d_{\mathrm{proj}}^{-2} d proj − 2 d p r o j d_{\mathrm{proj}} d proj r > d p r o j r > d_{\mathrm{proj}} r > d proj [136 ] ). Additional sources of non-thermal pressure like magnetic fields (see e.g., E. Fan de Voort et al. 2021; R. Ramesh & D. Nelson [106 ] ) and cosmic rays (CRs) can also be important. The role of CR electrons has been invoked to explain the observed X-ray surface brightness profiles of Milky-Way like galaxies (see e.g., P. F. Hopkins et al. [64 ] ). Cosmic ray pressure can also promote isochoric thermal instabilities in the cool gas, leading to the prevalence of low-density clumps (see e.g., I. S. Butsky et al. [13 ] ). However, theoretical predictions regarding the effect of CRs on the CGM are sensitive to resolution effects (see e.g., M. Weber et al. [130 ] ) and assumptions about the transport mechanism (see e.g., E. Quataert et al. [104 ] ).
Baryon census in galaxy halos
Independent of uncertainties in detailed 3D spatial profiles of the cool CGM, a robust measurement of the surface gas mass density can be made (e.g., J. K. Werk et al. [132 ] ; F. S. Zahedy et al. [138 ] ). Specifically, the best-fit photoionization models provide the necessary ionization fraction for different species, including H I, which in turn enable estimates of the total hydrogen column density, N H , c N_{\rm H,c} N H , c Table 2 and Table 3 ). The integrated total hydrogen column density N H N_{\rm H} N H [138 ] ).
These constraints are presented for M h ≈ 10 12 M ⊙ M_h \approx10^{12}M_\odot M h ≈ 1 0 12 M ⊙ 10 13 M ⊙ 10^{13}M_\odot 1 0 13 M ⊙ Figure 6 , following the symbol legend from Figure 5 . The hydrogen column density is converted to cool gas surface mass density following the relation Σ c o o l = 1.34 m H N H \Sigma_{\rm cool}=1.34\,m_{\rm H}\,N_{\rm H} Σ cool = 1.34 m H N H Σ c o o l = Σ 0 exp ( − d p r o j / d s ) \Sigma_{\rm cool}=\Sigma_0\exp(-d_{\rm proj}/d_s) Σ cool = Σ 0 exp ( − d proj / d s ) [138 ] ). Random draws from the posterior distribution for individual models are also depicted in Figure 6 . The best-fit profile for a M h ≈ 10 12 M ⊙ M_h \approx10^{12}M_\odot M h ≈ 1 0 12 M ⊙ Σ 0 ≈ 1.5 M ⊙ p c − 2 \Sigma_0 \approx1.5\,M_\odot\,{\rm pc}^{-2} Σ 0 ≈ 1.5 M ⊙ pc − 2 d s ≈ 50 k p c d_s \approx50\,{\rm kpc} d s ≈ 50 kpc M h ≈ 10 13 M ⊙ M_h \approx10^{13}M_\odot M h ≈ 1 0 13 M ⊙ Σ 0 ≈ 0.1 M ⊙ p c − 2 \Sigma_0 \approx0.1\,M_\odot\,{\rm pc}^{-2} Σ 0 ≈ 0.1 M ⊙ pc − 2 d s ≈ 200 k p c d_s \approx200\,{\rm kpc} d s ≈ 200 kpc d p r o j = 364 k p c d_{\rm proj}=364\,{\rm kpc} d proj = 364 kpc Σ 0 ≈ 1 M ⊙ p c − 2 \Sigma_0 \approx1\,M_\odot\,{\rm pc}^{-2} Σ 0 ≈ 1 M ⊙ pc − 2 d s ≈ 40 k p c d_s \approx40\,{\rm kpc} d s ≈ 40 kpc
Figure 6. Surface mass density constraints for cool gas in M h ≈ 10 12 M ⊙ M_h \approx10^{12}M_\odot M h ≈ 1 0 12 M ⊙ M h ≈ 10 13 M ⊙ M_h \approx10^{13}M_\odot M h ≈ 1 0 13 M ⊙ Figure 5 . The corresponding total hydrogen column densities are indicated on the right. Randomly drawn exponential profile fits for Σ c o o l \Sigma_{\rm cool} Σ cool
In contrast, the total hot gas mass differs substantially between massive LRG halos and lower-mass star-forming halos. This difference is evident in their respective surface mass density profiles, constructed from the three-dimensional density distributions of the hot gas (see e.g., [118 ] ). Following previous studies, the expected hot gas surface mass density is computed by integrating along the line of sight out to 2 × R v i r 2\times R_{\rm vir} 2 × R vir [141 ] ). The strong halo-mass dependence of the hot gas profile, contrasted with the comparable cool gas surface mass density profiles, therefore implies a clear decline in the fraction of baryons residing in the cool phase from M h ≈ 10 12 M ⊙ M_h \approx10^{12}M_\odot M h ≈ 1 0 12 M ⊙ 10 13 M ⊙ 10^{13}M_\odot 1 0 13 M ⊙ [122 , 33 ] ).
The baryon census between stars, the ISM, the cool CGM, and the hot halo is evaluated more precisely for the CUBS-midz and COS-LRG samples. Specifically, the exponential profile model for the surface mass density profiles of the cool gas, Σ \Sigma Σ M h ≈ 10 12 M ⊙ M_h \approx10^{12}M_\odot M h ≈ 1 0 12 M ⊙ 1 σ 1\sigma 1 σ ≈ 190 k p c \approx190\,{\rm kpc} ≈ 190 kpc M c o o l , 12 ≈ ( 1.0 M_{\rm cool,12}\approx(1.0 M cool , 12 ≈ ( 1.0 3.2 ) × 10 10 M ⊙ 3.2)\times10^{10}M_\odot 3.2 ) × 1 0 10 M ⊙ ≈ 75 % \approx75\% ≈ 75% d p r o j = 177 k p c d_{\rm proj}=177\,{\rm kpc} d proj = 177 kpc v c − v g a l = − 309 k m s − 1 v_c-v_{\rm gal}=-309\,{\rm km\,s}^{-1} v c − v gal = − 309 km s − 1 Figure 1 ). A comparable cool gas mass range, M c o o l , 12 ≈ ( 1.1 M_{\rm cool,12}\approx(1.1 M cool , 12 ≈ ( 1.1 3.2 ) × 10 10 M ⊙ 3.2)\times10^{10}M_\odot 3.2 ) × 1 0 10 M ⊙ 1 σ 1\sigma 1 σ M h o t , 12 ≈ ( 2.5 M_{\rm hot,12}\approx(2.5 M hot , 12 ≈ ( 2.5 5.0 ) × 10 10 M ⊙ 5.0)\times10^{10}M_\odot 5.0 ) × 1 0 10 M ⊙ M h ≈ 10 13 M ⊙ M_h\approx10^{13}M_\odot M h ≈ 1 0 13 M ⊙ R v i r ≈ 440 k p c R_{\rm vir}\approx440\,{\rm kpc} R vir ≈ 440 kpc M c o o l , 13 ≈ ( 1.0 M_{\rm cool,13}\approx(1.0 M cool , 13 ≈ ( 1.0 2.5 ) × 10 10 M ⊙ 2.5)\times10^{10}M_\odot 2.5 ) × 1 0 10 M ⊙ ≈ 80 % \approx80\% ≈ 80% 1 σ 1\sigma 1 σ M h o t , 13 ≈ ( 4.6 M_{\rm hot,13}\approx(4.6 M hot , 13 ≈ ( 4.6 9.1 ) × 10 11 M ⊙ 9.1)\times10^{11}M_\odot 9.1 ) × 1 0 11 M ⊙ [76 ] ).
These estimates are presented in Figure 7 , alongside corresponding baryonic mass present in stars and ISM. For the stellar mass, reported values for CUBS-midz and COS LRG hosts are used to determine the 84% upper bound, corresponding to the hatched region in Figure 7 . For the ISM of M h ≈ 10 12 M ⊙ M_h\approx10^{12}M_\odot M h ≈ 1 0 12 M ⊙ [34 ] ). Although an H I survey is not available for COS-LRGs, results from A. Bianchetti et al. [8 ] constraining the H I mass in quiescent M h ≈ 10 13 M ⊙ M_h \approx10^{13}\mathrm{M}_{\odot} M h ≈ 1 0 13 M ⊙ [9 ] ) agrees with the H I mass in CUBS-midz galaxies.
Figure 7. Census of baryonic mass in different reservoirs within M h ≈ 10 12 M ⊙ M_h \approx10^{12}\mathrm{M}_{\odot} M h ≈ 1 0 12 M ⊙ 10 13 M ⊙ 10^{13}\mathrm{M}_{\odot} 1 0 13 M ⊙ [95 ] ). The stellar mass range captures the observed variation in CUBS-midz galaxies and COS-LRGs. The ISM mass is evaluated using H I surveys from D. DePalma et al. [34 ] ; A. Bianchetti et al. [8 ] . The cool CGM mass is evaluated using the surface mass density profiles in Figure 6 , integrated out to the virial radius. The hot CGM mass is evaluated using hot gas profiles from P. Singh et al. [118 ] . Note that contributions from warm-hot gas are not accounted for here (e.g., Z. Qu et al. (2023); H.-W. Chen & F. S. Zahedy [24 ] ). Consequently, the combined mass in stars, the ISM, the cool CGM, and the hot halo represents a lower limit to the total baryon content of galaxy halos.
The cool CGM plays an important role in completing the baryon census, with the anticipated total baryonic mass M b = 0.17 M h M_b = 0.17M_h M b = 0.17 M h [95 ] ). The near constant cool gas mass observed across a broad mass range directly translates to a declining baryon mass fraction from
X c o o l ≡ M c o o l / M b ≈ 6 X_{\mathrm{cool}} \equiv M_{\mathrm{cool}}/M_b \approx6 X cool ≡ M cool / M b ≈ 6 19 % 19\% 19% M h ≈ 10 12 M ⊙ M_h \approx10^{12}\,\mathrm{M}_{\odot} M h ≈ 1 0 12 M ⊙ X c o o l ≈ 0.6 X_{\mathrm{cool}} \approx0.6 X cool ≈ 0.6 1.5 % 1.5\% 1.5% M h ≈ 10 13 M ⊙ M_h \approx10^{13}\,\mathrm{M}_{\odot} M h ≈ 1 0 13 M ⊙ [24 ] ). Therefore, the sum of all four components in stars, ISM, cool CGM, and the hot halo shown in Figure 7 constitutes a lower limit to the total baryon content in galaxy halos.
Kinematically aligned low- and high-ionization gas
QSO absorption line studies frequently reveal the presence of highly ionized species (e.g., C 3 + \mathrm{C}^{3+} C 3 + O 5 + \mathrm{O}^{5+} O 5 + [24 ] ). A collisionally ionized origin for O VI absorbers has been proposed in previous studies (see e.g., T. M. Tripp et al. [125 ] ; J. D. Meiring et al. [84 ] ; T. Hussain et al. [68 ] ; S. Pachat et al. [92 ] ; J. Nevalainen et al. [88 ] ; B. Rosenwasser et al. [110 ] ). However, a significant population of high-ionization absorbers can also be decomposed into a complex of narrow components with b ≲ 20 b \lesssim20 b ≲ 20 − 1 ^{-1} − 1 [123 ] ; B. D. Savage et al. [115 ] ). The observed high-ion column densities in such components can be reproduced with photoionization equilibrium models at densities n H ≲ 10 − 4 n_{\mathrm{H}} \lesssim10^{-4} n H ≲ 1 0 − 4 − 3 ^{-3} − 3 [17 ] ; F. S. Zahedy et al. [140 ] ; S. Kumar et al. [80 ] ). However, it is also possible that narrow high-ionization components originate from rapidly cooling gas (see e.g., O. Gnat & A. Sternberg [53 ] ; B. D. Oppenheimer & J. Schaye [89 ] ; S. Kumar & H.-W. Chen [79 ] ). An empirical characterization of warm-hot absorbers around CUBS-midz galaxies can be found in Z. Qu et al. 2024. Meanwhile, photoionized gas of densities n H ≈ 0.001 n_{\mathrm{H}} \approx0.001 n H ≈ 0.001 0.01 0.01 0.01 − 3 ^{-3} − 3 § 5.2 ) cannot reproduce the observed column densities of kinematically aligned high-ions (see also F. S. Zahedy et al. [140 ] ; T. J. Cooper et al. [31 ] ; Sameer et al. [113 ] ; Z. Qu et al. 2022). In the CUBS-midz sample, 22 out of 31 components with available ionization models show kinematically aligned O VI absorption. In eight of these components, the O VI absorption is too broad (b > 30 b > 30 b > 30 − 1 ^{-1} − 1 b b b N N N Figure 2 . There are two additional non-O VI-bearing components (s21c2 and s22c1) that require a two-phase solution. Note that most COS-LRG absorbers do not have the spectroscopic coverage for O IV, which is required to constrain the low-density phase, but this does not rule out a photoionized origin for the observed O VI.
Although some high-ionization components coincident in velocity space with the low-ionization gas, it is unclear whether the two gas phases are also co-spatial. It is possible that the two phases may only be superimposed in velocity space but reside at two distinct locations within the halo (see e.g., M. Ticorcas et al. 2026). Assuming pressure equilibrium with the hot halo (Figure 5 ), the implied physical distance for the low-density phases (n H ≲ 10 − 4 n_{\mathrm{H}} \lesssim10^{-4} n H ≲ 1 0 − 4 − 3 ^{-3} − 3 r ≳ 1 r \gtrsim1 r ≳ 1 ≲ 10 % \lesssim10\% ≲ 10% [82 ] ; S. D. Johnson et al. [71 ] ; S. Dutta et al. [38 ] ). Additionally, in an effort to derive physical conditions of circumgalactic gas around CUBS galaxies at z ≈ 1 z \approx1 z ≈ 1 ≈ 80 % \approx80\% ≈ 80% d p r o j < 100 d_{\mathrm{proj}} < 100 d proj < 100 ≲ 15 % \lesssim15\% ≲ 15% d p r o j = 100 d_{\mathrm{proj}} = 100 d proj = 100 200 200 200
Figure 8 depicts the density contrast between the low- and high-density phases in such components as a function of projected distance. The observed range of density fluctuations across systems is between one and two orders of magnitude. The density contrast has a direct consequence on the pressure balance between the two phases in such components. Although the thermal- and non-thermal motions in the low-density phase cannot be distinguished (see § 4.2 ), the available 3 σ 3\sigma 3 σ § 5.3 with limiting pressure for the low-density phase, pressure balance is ruled out for 8 out of 12 components in Figure 8 . Therefore, the two phases are likely not in dynamical equilibrium if they are co-spatial [see also 57 ] .
Figure 8. In this panel, the contrast in gas density n H n_{\mathrm{H}} n H Figure 5 , and non-isolated systems are marked using open symbols.
The observed density fluctuations in CUBS-midz multiphase components [see also 114 ] (Z. Qu et al. 2022) have been noted in simulations. In an analysis of CGM virialization, J. Stern et al. [121 ] predict that gas in the inner 50 kpc of simulated halos can show over two orders-of-magnitude in local density fluctuations [see also 73 ] . Disk settling is also achieved as a result of virialization in these simulations. The galaxy at d p r o j ≈ 66 d_{\rm proj} \approx66 d proj ≈ 66 Figure 8 with an extended rotating disk (system s10, M s t a r ≈ 6 × 10 10 M ⊙ M_{\rm star} \approx6 \times10^{10} M_{\odot} M star ≈ 6 × 1 0 10 M ⊙
Conclusions
In this paper, thermal and turbulent properties of circumgalactic absorbers are presented for 23 CUBS-midz galaxies at z = 0.4 z = 0.4 z = 0.4 0.7 0.7 0.7 HST COS and STIS enable accurate decomposition of absorption profiles across multiple ionic transitions. Ionizing conditions within individual clumps are derived using photoionization models, while thermal and non-thermal motions are separated using absorption-line width measurements. Deep galaxy surveys carried out in CUBS QSO fields using VLT MUSE alongside LDSS3C and IMACS instruments on the Magellan telescopes reveal the galaxy environment of absorber hosts. Combining constraints from the CUBS-midz sample with CGM measurements around intermediate redshift galaxies in the literature, the following key results arise –
Cool clumps around star-forming and quiescent galaxies are consistent with a photoionization equilibrium (PIE) scenario, as evident from the agreement between measured gas temperatures and the PIE expectation at the inferred gas density (Figure 3 ). The presence of outliers, however, may point to variations in the intensity and spectral shape of the ionizing radiation field due to local sources.
The non-thermal to thermal pressure ratio in circumgalactic clumps is connected to the turbulent Mach number. The median pressure ratio suggests sub-sonic turbulence in cool clumps, with M ≈ 0.5 \mathcal{M} \approx0.5 M ≈ 0.5 M ≈ 0.8 \mathcal{M} \approx0.8 M ≈ 0.8 Figure 4 ).
Cool gas densities exhibit a large variations of up to 2 dex in individual halos, but the density profile distinctly depends on the halo mass M h M_{\rm h} M h Figure 5 ). The cool gas pressure profile (including thermal and turbulent contributions) around halos with M h ≈ 10 12 M ⊙ M_{\rm h} \approx10^{12} M_{\odot} M h ≈ 1 0 12 M ⊙ 10 13 M ⊙ 10^{13} M_{\odot} 1 0 13 M ⊙
A near constant total cool halo gas mass is found across a broad halo mass range, translating to a declining baryon mass fraction in the cool phase from ≈ 6 \approx6 ≈ 6 19 % 19\% 19% M h ≈ 10 12 M ⊙ M_{\rm h} \approx10^{12} M_{\odot} M h ≈ 1 0 12 M ⊙ ≈ 0.6 \approx0.6 ≈ 0.6 1.5 % 1.5\% 1.5% M h ≈ 10 13 M ⊙ M_{\rm h} \approx10^{13} M_{\odot} M h ≈ 1 0 13 M ⊙ Figure 7 ).
Kinematically aligned absorption in low- and high-ionization species is identified in 12 out 31 CUBS-midz components, occurring in seven unique galaxies. Narrow line widths for high-ions (e.g., C 3 + \mathrm{C}^{3+} C 3 + Figure 8 ).
Acknowledgements
The authors thank an anonymous referee for a swift and constructive review process. The authors thank M. Voit for illuminating discussions regarding the contribution of the X-ray-detected hot phase to the halo baryon budget. S.K. and H.W.C. acknowledge partial support from HST-GO-17517.001A and from NASA ADAP 80NSSC23K0479. G.C.R. acknowledges the support from Grant 63667 from the John Templeton Foundation. Any opinions, findings, and conclusions or recommendations expressed in this material are those of the authors and do not necessarily reflect the views of the John Templeton Foundation. S.C. gratefully acknowledges support from the European Research Council (ERC) under the European Union’s Horizon 2020 Research and Innovation program grant agreement number 864361. M.C.C. is supported by the Brinson Foundation through the Brinson Prize Fellowship Program. D.D. acknowledges the support of the John A. Lyons Fellowship from MIT’s Office of Graduate Education. CAFG was supported by NSF through grants AST-2108230 and AST-2307327; by NASA through grants 80NSSC22k0809, 80NSSC22K1124 and 80NSSC24K1224; by STScI through grant JWST-AR03252.001-A; and by BSF through grant #2024262. S.L. acknowledges support by FONDECYT grant 1231187. This work is based on observations made with ESO Telescopes at the Paranal Observatory under program ID 0104.A-0147(A), observations made with the 6.5m Magellan Telescopes located at Las Campanas Observatory, and spectroscopic and imaging data gathered under the HST-GO-15163.01A, HST-GO-17517.001A, and HST-GO-17815-001A programs using the NASA/ESA Hubble Space Telescope operated by the Space Telescope Science Institute and the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-26555. This research has made use of the services of the ESO Science Archive Facility and the Astrophysics Data Service (ADS)3 [58 ] ), Scipy ([128 ] ), Astropy ([3 , 99 ] ), Matplotlib ([67 ] ), and emcee ([49 ] ).
Facilities: HST (COS/STIS), Magellan telescopes, VLT MUSE
Software: Cloudy ([46 ] )
Notes on individual systems
Details for absorption line analysis and photoionization modeling of individual CUBS-midz systems is described below. Note that all available transitions for ionic species within the spectral coverage of the survey are considered during the profile fitting, but only the strongest transitions have been indicated in Figures A.1–A.23 for brevity.
CUBS J230837.40-525844.94 at d p r o j = 34 d_{\mathrm{proj}} = 34 d proj = 34
This is an isolated, star-forming galaxy at a redshift z g a l = 0.5425 z_{\mathrm{gal}} = 0.5425 z gal = 0.5425 M s t a r ≈ 4 × 10 10 M ⊙ M_{\mathrm{star}} \approx4 \times10^{10}\ M_{\odot} M star ≈ 4 × 1 0 10 M ⊙ [28 ] ; F. S. Zahedy et al. [140 ] . The associated absorption is a Lyman Limit System (LLS) with a total H I column density of N ( H I ) ≈ 5 × 10 17 c m − 2 N(\mathrm{H\,I}) \approx5 \times10^{17}\ \mathrm{cm}^{-2} N ( H I ) ≈ 5 × 1 0 17 cm − 2 v c − v g a l ≈ − 64 k m s − 1 v_{\mathrm{c}} - v_{\mathrm{gal}} \approx-64\ \mathrm{km\ s}^{-1} v c − v gal ≈ − 64 km s − 1 β \beta β γ \gamma γ λ 858 \lambda858 λ 858 λ 834.4 \lambda834.4 λ 834.4 3 σ 3\sigma 3 σ λ 763 \lambda763 λ 763 λ 1083 \lambda1083 λ 1083
For deriving the gas density using photoionization models, the N III/N II ratios are well-constrained in components c1 and c6 and suggest an underlying gas density of log ( n H / c m − 3 ) ≈ − 2.3 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-2.3 log ( n H / cm − 3 ) ≈ − 2.3 log ( n H / c m − 3 ) ≈ − 2.1 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-2.1 log ( n H / cm − 3 ) ≈ − 2.1 λ 1036 \lambda1036 λ 1036 λ 858 \lambda858 λ 858 3 σ 3\sigma 3 σ log ( n H / c m − 3 ) < − 1.6 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) < -1.6 log ( n H / cm − 3 ) < − 1.6 log ( n H / c m − 3 ) > − 2.8 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) > -2.8 log ( n H / cm − 3 ) > − 2.8 log ( n H / c m − 3 ) ≈ − 2.3 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-2.3 log ( n H / cm − 3 ) ≈ − 2.3 log ( n H / c m − 3 ) ≈ − 1.7 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-1.7 log ( n H / cm − 3 ) ≈ − 1.7 b ≈ 30 k m s − 1 b \approx30\ \mathrm{km\ s}^{-1} b ≈ 30 km s − 1 Table 2 .
To infer the underlying gas temperature and non-thermal motions, the measured line widths for H I and Mg II are used. Derived gas temperatures range between log ( T / K ) ≈ 3.8 − 4.2 \log(T/\mathrm{K}) \approx3.8 - 4.2 log ( T / K ) ≈ 3.8 − 4.2 b N T ≲ 6 k m s − 1 b_{\mathrm{NT}} \lesssim6\ \mathrm{km\ s}^{-1} b NT ≲ 6 km s − 1 b ≲ 12 k m s − 1 b \lesssim12\ \mathrm{km\ s}^{-1} b ≲ 12 km s − 1 T T T b N T b_{\mathrm{NT}} b NT Table 2 that are constrained using H I and Mg II absorption.
CUBS J213553.27-531700.70 at d p r o j = 35 d_{\mathrm{proj}} = 35 d proj = 35
This is an isolated, star-forming dwarf galaxy at z g a l = 0.6768 z_{\mathrm{gal}} = 0.6768 z gal = 0.6768 M s t a r ≈ 10 8 M ⊙ M_{\mathrm{star}} \approx10^{8}\ M_{\odot} M star ≈ 1 0 8 M ⊙ n H < 10 − 4 c m − 3 n_{\mathrm{H}} < 10^{-4}\,\mathrm{cm}^{-3} n H < 1 0 − 4 cm − 3
Figure A.1. The HST F814W image for the host galaxy of system s1 is shown in the top-left panel, with the arrow pointing towards the background QSO. The optical MUSE spectrum is shown in the top-right panel, with prominent emission and absorption features marked using dotted lines. The associated QSO absorption spectra are shown in the bottom panel. The velocity axis is defined relative to the galaxy redshift. The continuum-normalized QSO flux is shown in black, with associated 68% (1 σ 1\sigma 1 σ
CUBS J011139.36-031605.19 at d p r o j = 42 d_{\mathrm{proj}}=42 d proj = 42
This is a quiescent galaxy at z g a l = 0.5766 z_{\mathrm{gal}}=0.5766 z gal = 0.5766 M s t a r ≈ 10 11 M ⊙ M_{\mathrm{star}} \approx10^{11}M_{\odot} M star ≈ 1 0 11 M ⊙ [11 ] . Additional satellite galaxies are identified near this host, extended out to d p r o j = 220 d_{\mathrm{proj}}=220 d proj = 220 α \alpha α N ( H I ) ≈ 10 20 c m − 2 N(\mathrm{H\,I}) \approx10^{20}\,\mathrm{cm}^{-2} N ( H I ) ≈ 1 0 20 cm − 2 2 _2 2
Figure A.2. Galaxy image, optical spectrum, and QSO absorption spectra for system s2.
Figure A.3. Galaxy image, optical spectrum, and QSO absorption spectra for system s3.
Table A.1. Absorption properties along the sightline at d p r o j = 34 d_{\mathrm{proj}}=34 d proj = 34
Table A.2. Absorption properties along the sightline at d p r o j = 35 d_{\mathrm{proj}}=35 d proj = 35
Table A.3. Integrated absorption properties along the sightline at d p r o j = 42 d_{\mathrm{proj}}=42 d proj = 42
Table A.4. Absorption properties along the sightline at d p r o j = 43 d_{\mathrm{proj}}=43 d proj = 43
setting the gas ionization state, UVB-only photoionization models used in this work are not suitable for this system, and the reader is referred to E. Boettcher et al. [11 ] for a detailed examination.
CUBS J015455.18−071222.15 at d p r o j = 43 d_{\mathrm{proj}}=43 d proj = 43
This is an isolated, star-forming dwarf galaxy at z g a l = 0.4550 z_{\mathrm{gal}}=0.4550 z gal = 0.4550 M s t a r ≈ 10 8 M ⊙ M_{\mathrm{star}}\approx10^{8}M_{\odot} M star ≈ 1 0 8 M ⊙ α \alpha α β \beta β
CUBS J011035.01−164833.32 at d p r o j = 54 d_{\mathrm{proj}}=54 d proj = 54
This is a non-isolated, star-forming galaxy at z g a l = 0.4724 z_{\mathrm{gal}}=0.4724 z gal = 0.4724 M s t a r ≈ 10 10 M ⊙ M_{\mathrm{star}}\approx10^{10}M_{\odot} M star ≈ 1 0 10 M ⊙ [31 ] . The associated absorption is a partial Lyman Limit System (pLLS) with a neutral hydrogen column density of N ( H I ) ≈ 10 16 c m − 2 N(\mathrm{H\ I})\approx10^{16}\ \mathrm{cm}^{-2} N ( H I ) ≈ 1 0 16 cm − 2 λ \lambda λ λ \lambda λ
Figure A.4. Galaxy image, optical spectrum, and QSO absorption spectra for system s4.
Si III λ 1206 \lambda1206 λ 1206 λ 977 \lambda977 λ 977 λ 832 \lambda832 λ 832 λ 833.3 \lambda833.3 λ 833.3 λ 832.7 \lambda832.7 λ 832.7 λ 989 \lambda989 λ 989 λ λ 1548 , 1550 \lambda\lambda1548, 1550 λλ 1548 , 1550 b ≈ 20 b \approx20 b ≈ 20 − 1 ^{-1} − 1 λ 765 \lambda765 λ 765 λ 787 \lambda787 λ 787 λ λ 1031 , 1037 \lambda\lambda1031, 1037 λλ 1031 , 1037 b ≈ 30 b \approx30 b ≈ 30 − 1 ^{-1} − 1 α \alpha α β \beta β v c − v g a l ≈ − 84 v_c-v_{\rm gal} \approx-84 v c − v gal ≈ − 84 − 1 ^{-1} − 1 − 1 ^{-1} − 1
For ionization modeling, the measured O III/O II ratios in components c2 through c4 suggest an underlying gas density of log ( n H / c m − 3 ) ≈ − 2.6 \log(n_{\rm H}/{\rm cm}^{-3}) \approx-2.6 log ( n H / cm − 3 ) ≈ − 2.6 log ( n H / c m − 3 ) ≈ − 4.4 \log(n_{\rm H}/{\rm cm}^{-3}) \approx-4.4 log ( n H / cm − 3 ) ≈ − 4.4
CUBS J233912.33−552354.27 at d p r o j = 55 d_{\rm proj}=55 d proj = 55
This is an isolated, star-forming dwarf galaxy at z g a l = 0.5726 z_{\rm gal}=0.5726 z gal = 0.5726 M s t a r ≈ 3 × 10 8 M ⊙ M_{\rm star} \approx3 \times10^{8}M_{\odot} M star ≈ 3 × 1 0 8 M ⊙ λ 1550 \lambda1550 λ 1550 λ 765 \lambda765 λ 765 λ 832 \lambda832 λ 832 b ≈ 10 b \approx10 b ≈ 10 − 1 ^{-1} − 1 b ≈ 25 b \approx25 b ≈ 25 − 1 ^{-1} − 1 λ 1548 \lambda1548 λ 1548 λ 787 \lambda787 λ 787 λ 1031 \lambda1031 λ 1031 b ≈ 20 b \approx20 b ≈ 20 − 1 ^{-1} − 1 − 1 ^{-1} − 1 α \alpha α
For ionization modeling of c1, the C IV/C III ratio suggests a density of log ( n H / c m − 3 ) ≈ − 3.9 \log(n_{\rm H}/{\rm cm}^{-3}) \approx-3.9 log ( n H / cm − 3 ) ≈ − 3.9 log ( n H / c m − 3 ) ≈ − 3.6 \log(n_{\rm H}/{\rm cm}^{-3}) \approx-3.6 log ( n H / cm − 3 ) ≈ − 3.6
CUBS J224500.93−493154.29 at d p r o j = 58 d_{\rm proj}=58 d proj = 58
This is an isolated, star-forming dwarf galaxy at z g a l = 0.5338 z_{\rm gal}=0.5338 z gal = 0.5338 M s t a r = 10 8 M ⊙ M_{\rm star}=10^{8}M_{\odot} M star = 1 0 8 M ⊙
Figure A.5. Galaxy image, optical spectrum, and QSO absorption spectra for system s5.
CUBS J011138.57−031609.64 at d p r o j = 63 d_{\mathrm{proj}}=63 d proj = 63
This is an isolated, quiescent galaxy at z g a l = 0.6482 z_{\mathrm{gal}}=0.6482 z gal = 0.6482 M s t a r ≈ 10 11 M ⊙ M_{\mathrm{star}}\approx10^{11}M_{\odot} M star ≈ 1 0 11 M ⊙
CUBS J213552.98−531645.51 at d p r o j = 65 d_{\mathrm{proj}}=65 d proj = 65
This is an isolated, star-forming dwarf galaxy at z g a l = 0.5082 z_{\mathrm{gal}}=0.5082 z gal = 0.5082 M s t a r ≈ 10 8 M ⊙ M_{\mathrm{star}}\approx10^{8}M_{\odot} M star ≈ 1 0 8 M ⊙
CUBS J035723.09−481215.82 at d p r o j = 66 d_{\mathrm{proj}}=66 d proj = 66
This is an isolated, star-forming galaxy at z g a l = 0.4357 z_{\mathrm{gal}}=0.4357 z gal = 0.4357 M s t a r ≈ 5 × 10 10 M ⊙ M_{\mathrm{star}}\approx5\times10^{10}M_{\odot} M star ≈ 5 × 1 0 10 M ⊙ [28 ] ; F. S. Zahedy et al. [140 ] . The associated absorption is a LLS with a neutral hydrogen column density of N ( H I ) ≈ 2 × 10 17 c m − 2 N(\mathrm{H\,I})\approx2\times10^{17}\,\mathrm{cm}^{-2} N ( H I ) ≈ 2 × 1 0 17 cm − 2 λ 834 \lambda834 λ 834 λ 1036 \lambda1036 λ 1036 v c − v g a l ≈ − 22 k m s − 1 v_{\mathrm{c}}-v_{\mathrm{gal}}\approx-22\,\mathrm{km\,s}^{-1} v c − v gal ≈ − 22 km s − 1 λ λ 903.9 , 903.6 \lambda\lambda903.9,903.6 λλ 903.9 , 903.6 λ 1036 \lambda1036 λ 1036 λ 834 \lambda834 λ 834 λ 977 \lambda977 λ 977 λ 989 \lambda989 λ 989 λ 832 \lambda832 λ 832 λ λ 833.3 , 832.7 \lambda\lambda833.3,832.7 λλ 833.3 , 832.7 λ 834 \lambda834 λ 834 λ λ 1548 \lambda\lambda1548 λλ 1548 λ 787 \lambda787 λ 787 b ≲ 20 b \lesssim20 b ≲ 20 − 1 ^{-1} − 1 λ λ 1031 \lambda\lambda1031 λλ 1031 b ≈ 35 b \approx35 b ≈ 35 − 1 ^{-1} − 1 λ 765 \lambda765 λ 765 λ λ 1238 \lambda\lambda1238 λλ 1238 α \alpha α β \beta β v c − v g a l ≈ 37 v_{\mathrm{c}} - v_{\mathrm{gal}} \approx37 v c − v gal ≈ 37 − 1 ^{-1} − 1 v c − v g a l ≈ 154 v_{\mathrm{c}} - v_{\mathrm{gal}} \approx154 v c − v gal ≈ 154 − 1 ^{-1} − 1 log ( n H / c m − 3 ) ≈ − 2.4 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-2.4 log ( n H / cm − 3 ) ≈ − 2.4 log ( n H / c m − 3 ) ≈ − 2.7 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-2.7 log ( n H / cm − 3 ) ≈ − 2.7 log ( n H / c m − 3 ) ≈ − 4.3 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-4.3 log ( n H / cm − 3 ) ≈ − 4.3
Figure A.6. Galaxy image, optical spectrum, and QSO absorption spectra for system s6.
Figure A.7. Galaxy image, optical spectrum, and QSO absorption spectra for system s7.
Figure A.8. Galaxy image, optical spectrum, and QSO absorption spectra for system s8.
Figure A.9. Galaxy image, optical spectrum, and QSO absorption spectra for system s9.
Figure A.10. Galaxy image, optical spectrum, and QSO absorption spectra for system s10.
strained by the sensitive O III/Mg II limit (assuming [Mg/O] = 0), in addition to requiring the low-density phase to produce more O IV relative to the high-density phase (effectively treating the O IV/O III ratio as an upper limit for the denser phase).
CUBS J035721.44−481225.86 at d p r o j = 67 d_{\mathrm{proj}} = 67 d proj = 67
This is an isolated, star-forming dwarf galaxy at z g a l = 0.4489 z_{\mathrm{gal}} = 0.4489 z gal = 0.4489 M s t a r ≈ 2 × 10 8 M ⊙ M_{\mathrm{star}} \approx2 \times10^{8}M_{\odot} M star ≈ 2 × 1 0 8 M ⊙ λ \lambda λ λ λ \lambda\lambda λλ
component along c2. The C III absorption can be resolved into two components along c1 and c2. The O VI doublet is contaminated, although the blue wing of the O VI λ \lambda λ log ( n H / c m − 3 ) ≈ − 4.0 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-4.0 log ( n H / cm − 3 ) ≈ − 4.0
CUBS J042052.57−565038.97 at d p r o j = 80 d_{\mathrm{proj}} = 80 d proj = 80
This is a non-isolated, star-forming galaxy at z g a l = 0.5839 z_{\mathrm{gal}} = 0.5839 z gal = 0.5839 M s t a r ≈ 3 × 10 9 M ⊙ M_{\mathrm{star}} \approx3 \times10^{9}M_{\odot} M star ≈ 3 × 1 0 9 M ⊙ λ 977 \lambda977 λ 977 λ 832 \lambda832 λ 832 λ 765 \lambda765 λ 765 b ≈ 36 b \approx36 b ≈ 36 − 1 ^{-1} − 1 b ≈ 18 b \approx18 b ≈ 18 − 1 ^{-1} − 1 b ≈ 40 b \approx40 b ≈ 40 − 1 ^{-1} − 1 b ≈ 13 b \approx13 b ≈ 13 − 1 ^{-1} − 1
Figure A.11. Galaxy image, optical spectrum, and QSO absorption spectra for system s11.
For ionization modeling of c1, a low-density phase of log ( n H / c m − 3 ) ≈ − 4.3 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-4.3 log ( n H / cm − 3 ) ≈ − 4.3 log ( n H / c m − 3 ) ≈ − 2.9 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-2.9 log ( n H / cm − 3 ) ≈ − 2.9 log ( n H / c m − 3 ) ≈ − 4.4 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-4.4 log ( n H / cm − 3 ) ≈ − 4.4 λ 702 \lambda702 λ 702 λ 832 \lambda832 λ 832
CUBS J045416.38-611640.90 at d p r o j = 82 d_{\mathrm{proj}}=82 d proj = 82
This is a non-isolated, quiescent galaxy at z g a l = 0.4315 z_{\mathrm{gal}} = 0.4315 z gal = 0.4315 M s t a r ≈ 10 10 M ⊙ M_{\mathrm{star}} \approx10^{10}M_{\odot} M star ≈ 1 0 10 M ⊙
CUBS J035723.55-481215.38 at d p r o j = 97 d_{\mathrm{proj}}=97 d proj = 97
This is a non-isolated, star-forming galaxy at z g a l = 0.4807 z_{\mathrm{gal}} = 0.4807 z gal = 0.4807 M s t a r ≈ 10 10 M ⊙ M_{\mathrm{star}} \approx10^{10}M_{\odot} M star ≈ 1 0 10 M ⊙
CUBS J015453.61-071216.79 at d p r o j = 99 d_{\mathrm{proj}}=99 d proj = 99
This is a non-isolated, star-forming galaxy with stellar mass M s t a r ≈ 10 10 M ⊙ M_{\mathrm{star}} \approx10^{10}M_{\odot} M star ≈ 1 0 10 M ⊙ b ≈ 16 b \approx16 b ≈ 16 − 1 ^{-1} − 1 v c − v g a l ≈ 78 v_{\mathrm{c}} - v_{\mathrm{gal}} \approx78 v c − v gal ≈ 78 − 1 ^{-1} − 1 λ 832 \lambda832 λ 832 λ 765 \lambda765 λ 765 b ≈ 34 b \approx34 b ≈ 34 − 1 ^{-1} − 1 λ 977 \lambda977 λ 977 λ 787 \lambda787 λ 787 v c − v g a l ≈ 26 v_{\mathrm{c}} - v_{\mathrm{gal}} \approx26 v c − v gal ≈ 26 − 1 ^{-1} − 1 − 1 ^{-1} − 1 λ λ 770 , 780 \lambda\lambda770, 780 λλ 770 , 780 b ≈ 69 b \approx69 b ≈ 69 − 1 ^{-1} − 1 β \beta β γ \gamma γ
Figure A.12. Galaxy image, optical spectrum, and QSO absorption spectra for system s12.
For ionization modeling of c1, a single phase of density log ( n H / c m − 3 ) ≈ − 4.4 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-4.4 log ( n H / cm − 3 ) ≈ − 4.4 log ( n H / c m − 3 ) ≈ − 4.3 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-4.3 log ( n H / cm − 3 ) ≈ − 4.3 log ( n H / c m − 3 ) ≈ − 3.3 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-3.3 log ( n H / cm − 3 ) ≈ − 3.3 log ( n H / c m − 3 ) ≈ − 3.7 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-3.7 log ( n H / cm − 3 ) ≈ − 3.7 log ( n H / c m − 3 ) ≈ − 3.4 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-3.4 log ( n H / cm − 3 ) ≈ − 3.4
CUBS J224458.73-493158.12 at dproj = 114 kpc
This is an isolated, star-forming galaxy at z g a l = 0.5863 z_{\mathrm{gal}} = 0.5863 z gal = 0.5863 M s t a r ≈ 2 × 10 10 M ⊙ M_{\mathrm{star}} \approx2 \times10^{10}M_{\odot} M star ≈ 2 × 1 0 10 M ⊙ λ 832 \lambda832 λ 832 λ 977 \lambda977 λ 977 λ 787 \lambda787 λ 787
Figure A.13. Galaxy image, optical spectrum, and QSO absorption spectra for system s13.
Figure A.14. Galaxy image, optical spectrum, and QSO absorption spectra for system s14.
Figure A.15. Galaxy image, optical spectrum, and QSO absorption spectra for system s15.
Figure A.16. Galaxy image, optical spectrum, and QSO absorption spectra for system s16.
density of log ( n H / c m − 3 ) ≈ − 2.7 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-2.7 log ( n H / cm − 3 ) ≈ − 2.7
CUBS J213553.48−531714.50 at d p r o j = 128 d_{\mathrm{proj}}=128 d proj = 128
This is a non-isolated, quiescent galaxy at z g a l = 0.6219 z_{\mathrm{gal}}=0.6219 z gal = 0.6219 M s t a r ≈ 5 × 10 10 M ⊙ M_{\mathrm{star}} \approx5 \times10^{10}M_{\odot} M star ≈ 5 × 1 0 10 M ⊙ [28 ] ; F. S. Zahedy et al. [140 ] . The corresponding absorption is yet another LLS with N ( H I ) ≈ 2 × 10 18 c m − 2 N(\mathrm{H\,I}) \approx2 \times10^{18}\ \mathrm{cm}^{-2} N ( H I ) ≈ 2 × 1 0 18 cm − 2 λ 1334 \lambda1334 λ 1334 λ 1260 \lambda1260 λ 1260 λ 977 \lambda977 λ 977 v c − v g a l ≈ 8 k m s − 1 v_{\mathrm{c}}-v_{\mathrm{gal}} \approx8\ \mathrm{km\ s}^{-1} v c − v gal ≈ 8 km s − 1 v c − v g a l ≈ 174 k m s − 1 v_{\mathrm{c}}-v_{\mathrm{gal}} \approx174\ \mathrm{km\ s}^{-1} v c − v gal ≈ 174 km s − 1 v c − v g a l ≈ − 141 k m s − 1 v_{\mathrm{c}}-v_{\mathrm{gal}} \approx-141\ \mathrm{km\ s}^{-1} v c − v gal ≈ − 141 km s − 1 β \beta β b ≈ 10 k m s − 1 b \approx10\ \mathrm{km\ s}^{-1} b ≈ 10 km s − 1 b ≈ 35 k m s − 1 b \approx35\ \mathrm{km\ s}^{-1} b ≈ 35 km s − 1 N c 3 ( C I I I ) > N c 5 ( C I I I ) N_{\mathrm{c3}}(\mathrm{C\,III}) > N_{\mathrm{c5}}(\mathrm{C\,III}) N c3 ( C III ) > N c5 ( C III ) N c 4 ( C I I I ) > N c 5 ( C I I I ) N_{\mathrm{c4}}(\mathrm{C\,III}) > N_{\mathrm{c5}}(\mathrm{C\,III}) N c4 ( C III ) > N c5 ( C III ) 3 σ 3\sigma 3 σ
Ionization modeling results are available for components c2, c3, and c4. For component c2, the C IV/C III ratio suggests low-density gas of log ( n H / c m − 3 ) ≈ − 4.4 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-4.4 log ( n H / cm − 3 ) ≈ − 4.4 log ( n H / c m − 3 ) ≈ − 2.1 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-2.1 log ( n H / cm − 3 ) ≈ − 2.1
CUBS J224458.93−493129.29 at d p r o j = 138 d_{\mathrm{proj}}=138 d proj = 138
This is an isolated, star-forming galaxy at z g a l = 0.4893 z_{\mathrm{gal}}=0.4893 z gal = 0.4893 M s t a r ≈ 10 10 M ⊙ M_{\mathrm{star}} \approx10^{10}\ M_{\odot} M star ≈ 1 0 10 M ⊙
Figure A.17. Galaxy image, optical spectrum, and QSO absorption spectra for system s17.
CUBS J011957.16-201003.55 at d p r o j = 141 d_{\mathrm{proj}}=141 d proj = 141
This is a non-isolated, star-forming galaxy at z g a l = 0.4517 z_{\mathrm{gal}}=0.4517 z gal = 0.4517 M s t a r ≈ 4 × 10 10 M ⊙ M_{\mathrm{star}}\approx4\times10^{10}M_{\odot} M star ≈ 4 × 1 0 10 M ⊙ v c − v g a l ≈ 21 v_{\mathrm{c}}-v_{\mathrm{gal}}\approx21 v c − v gal ≈ 21 − 1 ^{-1} − 1 λ 832 \lambda832 λ 832 λ 787 \lambda787 λ 787 v c 3 − v c 2 ≈ 24 v_{\mathrm{c3}}-v_{\mathrm{c2}}\approx24 v c3 − v c2 ≈ 24 − 1 ^{-1} − 1 λ 1031 \lambda1031 λ 1031 λ 1037 \lambda1037 λ 1037 v c − v g a l ≈ − 56 v_{\mathrm{c}}-v_{\mathrm{gal}}\approx-56 v c − v gal ≈ − 56 − 1 ^{-1} − 1 v c − v g a l ≈ 277 v_{\mathrm{c}}-v_{\mathrm{gal}}\approx277 v c − v gal ≈ 277 − 1 ^{-1} − 1 α \alpha α β \beta β
For the ionization modeling of c2, a two-phase solution is necessary. The high-density phase is constrained at log ( n H / c m − 3 ) ≈ − 2.6 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3})\approx-2.6 log ( n H / cm − 3 ) ≈ − 2.6 log ( n H / c m − 3 ) ≈ − 4.3 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3})\approx-4.3 log ( n H / cm − 3 ) ≈ − 4.3 log ( n H / c m − 3 ) ≈ − 3.7 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3})\approx-3.7 log ( n H / cm − 3 ) ≈ − 3.7
Figure A.18. Galaxy image, optical spectrum, and QSO absorption spectra for system s18.
CUBS J045415.97−611602.32 at d p r o j = 162 d_{\mathrm{proj}} = 162 d proj = 162
This is an isolated quiescent galaxy at z g a l = 0.6011 z_{\mathrm{gal}} = 0.6011 z gal = 0.6011 M s t a r ≈ 3 × 10 10 M ⊙ M_{\mathrm{star}} \approx3 \times10^{10}M_{\odot} M star ≈ 3 × 1 0 10 M ⊙
CUBS J224457.76−493159.13 at d p r o j = 177 d_{\mathrm{proj}} = 177 d proj = 177
This is a non-isolated, quiescent galaxy at z g a l = 0.6203 z_{\mathrm{gal}} = 0.6203 z gal = 0.6203 M s t a r ≈ 5 × 10 10 M ⊙ M_{\mathrm{star}} \approx5 \times10^{10}M_{\odot} M star ≈ 5 × 1 0 10 M ⊙ v c − v g a l ≈ − 309 v_{\mathrm{c}} - v_{\mathrm{gal}} \approx-309 v c − v gal ≈ − 309 − 1 ^{-1} − 1 λ \lambda λ λ \lambda λ v c − v g a l ≈ − 394 v_{\mathrm{c}} - v_{\mathrm{gal}} \approx-394 v c − v gal ≈ − 394 − 1 ^{-1} − 1 λ \lambda λ λ λ \lambda\lambda λλ λ λ \lambda\lambda λλ
For the ionization modeling of c2, a high-density phase is constrained at log ( n H / c m − 3 ) ≈ − 2.5 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-2.5 log ( n H / cm − 3 ) ≈ − 2.5 log ( n H / c m − 3 ) ≈ − 3.8 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-3.8 log ( n H / cm − 3 ) ≈ − 3.8
CUBS J011037.69−164915.63 at d p r o j = 364 d_{\mathrm{proj}} = 364 d proj = 364
This is a non-isolated, quiescent galaxy at z g a l = 0.5404 z_{\mathrm{gal}} = 0.5404 z gal = 0.5404 M s t a r ≈ 2 × 10 11 M ⊙ M_{\mathrm{star}} \approx2 \times10^{11}M_{\odot} M star ≈ 2 × 1 0 11 M ⊙ [31 ] . The associated absorption corresponds to a pLLS with N ( H I ) ≈ 3 × 10 16 N(\mathrm{H\,I}) \approx3 \times10^{16} N ( H I ) ≈ 3 × 1 0 16 − 2 ^{-2} − 2 v c 2 − v c 1 ≈ 20 v_{\mathrm{c2}} - v_{\mathrm{c1}} \approx20 v c2 − v c1 ≈ 20 − 1 ^{-1} − 1 λ \lambda λ λ \lambda λ v c − v g a l ≈ 246 k m s − 1 v_{\mathrm{c}} - v_{\mathrm{gal}} \approx246\ \mathrm{km\ s^{-1}} v c − v gal ≈ 246 km s − 1 β \beta β γ \gamma γ
Figure A.19. Galaxy image, optical spectrum, and QSO absorption spectra for system s19.
For the ionization modeling of c1, the C III/C II, N III/N II, O III/O II, and S III/S II ratios suggest a consistent density of log ( n H / c m − 3 ) ≈ − 2.2 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-2.2 log ( n H / cm − 3 ) ≈ − 2.2 log ( n H / c m − 3 ) ≈ − 4.0 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-4.0 log ( n H / cm − 3 ) ≈ − 4.0
For the satellite component c2, the C III, N IV, and O III columns are robustly measured. The corresponding absorption for N III, C IV, and O IV is weak, so a column density upper limit is evaluated from the posterior distribution for these ions.
A single gas phase is density log ( n H / c m − 3 ) ≈ − 3.4 \log(n_{\mathrm{H}}/\mathrm{cm}^{-3}) \approx-3.4 log ( n H / cm − 3 ) ≈ − 3.4
CUBS J045406.34-611518.54 at dproj=640 kpc
This is a non-isolated, quiescent galaxy at z g a l = 0.5811 z_{\mathrm{gal}} = 0.5811 z gal = 0.5811 M s t a r ≈ 4 × 10 11 M ⊙ M_{\mathrm{star}} \approx4 \times10^{11}M_{\odot} M star ≈ 4 × 1 0 11 M ⊙
Figure A.20. Galaxy image, optical spectrum, and QSO absorption spectra for system s20.
Figure A.21. Galaxy image, optical spectrum, and QSO absorption spectra for system s21.
Figure A.22. Galaxy image, optical spectrum, and QSO absorption spectra for system s22.
Figure A.23. Galaxy image, optical spectrum, and QSO absorption spectra for system s23.
Table A.5. Absorption properties along the sightline at d p r o j = 54 d_{\rm proj}=54 d proj = 54
Table A.6. Absorption properties along the sightline at d p r o j = 55 d_{\rm proj}=55 d proj = 55
Table A.7. Absorption properties along the sightline at d p r o j = 58 d_{\rm proj}=58 d proj = 58
Table A.8. Absorption properties along the sightline at d p r o j = 63 d_{\rm proj}=63 d proj = 63
Table A.9. Absorption properties along the sightline at d p r o j = 65 d_{\rm proj}=65 d proj = 65
Table A.10. Absorption properties along the sightline at d p r o j = 66 d_{\rm proj}=66 d proj = 66
Table A.11. Absorption properties along the sightline at d p r o j = 67 d_{\rm proj}=67 d proj = 67
Table A.12. Absorption properties along the sightline at d p r o j = 80 d_{\rm proj}=80 d proj = 80
Table A.13. Absorption properties along the sightline at d p r o j = 82 d_{\rm proj}=82 d proj = 82
Table A.14. Absorption properties along the sightline at d p r o j = 97 d_{\rm proj}=97 d proj = 97
Table A.15. Absorption properties along the sightline at d p r o j = 99 d_{\mathrm{proj}} = 99 d proj = 99
Table A.15.
Table A.16. Absorption properties along the sightline at d p r o j = 114 d_{\mathrm{proj}} = 114 d proj = 114
Table A.16.
Table A.17. Absorption properties along the sightline at d p r o j = 128 d_{\mathrm{proj}} = 128 d proj = 128
Table A.17.
Table A.18. Absorption properties along the sightline at d p r o j = 138 d_{\mathrm{proj}} = 138 d proj = 138
Table A.18.
Table A.19. Absorption properties along the sightline at d p r o j = 141 d_{\rm proj}=141 d proj = 141
Table A.20. Absorption properties along the sightline at d p r o j = 162 d_{\rm proj}=162 d proj = 162
Table A.21. Absorption properties along the sightline at d p r o j = 177 d_{\rm proj}=177 d proj = 177
Table A.22. Absorption properties along the sightline at d p r o j = 364 d_{\rm proj}=364 d proj = 364
Table A.23. Absorption properties along the sightline at d p r o j = 640 k p c d_{\mathrm{proj}}=640\ \mathrm{kpc} d proj = 640 kpc
