# Observations of a Magellanic Corona

Dhanesh Krishnarao,1,2,3∗ Andrew J. Fox,4 Elena D’Onghia,5
Bart P. Wakker,5 Frances H. Cashman,1 J. Christopher Howk,6
Scott Lucchini,7 David M. French,1 Nicolas Lehner6

1Space Telescope Science Institute
3700 San Martin Dr, Baltimore, MD 21218, USA
2William H. Miller III Department of Physics & Astronomy, Johns Hopkins
University
3400 N. Charles Street, Baltimore, MD 21218, USA
3Department of Physics, Colorado College
14 East Cache La Poudre Street, Colorado Springs, CO 80903, USA
4AURA for ESA, Space Telescope Science Institute
3700 San Martin Dr, Baltimore, MD 21218, USA
5Department of Astronomy, University of Wisconsin-Madison
475 N Charter St, Madison, WI 53706, USA
6Department of Physics & Astronomy, University of Notre Dame
225 Nieuwland Science Hall, Notre Dame, IN 46556, USA
7Department of Physics, University of Wisconsin-Madison
475 N Charter St, Madison, WI 53706, USA

∗To whom correspondence should be addressed; E-mail:
<EMAIL_ADDRESS>
The Large and Small Magellanic Clouds (LMC/SMC) are the closest major
satellite galaxies of the Milky Way. They are likely on their first passage on
an infalling orbit towards our Galaxy [1] and trace the ongoing dynamics of
the Local Group [2]. Recent measurements of a high mass for the LMC ($M_{\rm
halo}\approx 10^{11.1-11.4}\,M_{\odot}$) [3, 4, 5, 6] imply the LMC should
host a Magellanic Corona: a collisionally ionized, warm-hot gaseous halo at
the virial temperature ($10^{5.3-5.5}~{}K$) initially extending out to the
virial radius ($100-130$ kpc). Such a Corona would have shaped the formation
of the Magellanic Stream [7], a tidal gas structure extending over
$200^{\circ}$ across the sky [8, 9, 2] that is bringing in metal poor gas to
the Milky Way [10]. No observational evidence for such an extended Corona has
been published previously, with detections of highly ionized gas only reported
in directions directly toward the LMC, where winds from the LMC disk may
dominate [11, 12] Here we show evidence for this Magellanic Corona with a
potential direct detection in highly ionized oxygen (O+5), and indirectly via
triply-ionized carbon and silicon, seen in ultraviolet absorption toward
background quasars. We find that the Magellanic Corona is part of a pervasive
multiphase Magellanic circumgalactic medium (CGM) seen in many ionization
states with a declining projected radial profile out to at least $35$ kpc from
the LMC and a total ionized CGM mass of
$\log_{10}(M_{\mathrm{H~{}{II},CGM}}/M_{\odot})\approx 9.1\pm 0.2$. The
evidence for the Magellanic Corona is a crucial step forward in characterizing
the Magellanic Group and its nested evolution with the Local Group.

We use a sample of 28 Hubble Space Telescope (_HST_)/Cosmic Origins
Spectrograph (COS) spectra of background UV-bright quasars within an angular
separation of the LMC of 45°, corresponding to an impact parameter
$\rho_{\mathrm{LMC}}<35~{}\mathrm{kpc}$, spanning one-third of the LMC’s
initial virial radius. 6 of these sightlines also have archival Far
Ultraviolet Spectroscopic Explorer (_FUSE_) spectra with high enough S/N to
measure O VI absorption. All spectra have a signal-to-noise ratio S/N$>$7 per
resolution element. Our analysis reveals pervasive low- and high-ion
absorption in several phases centered around the LMC both on the sky and in
velocity, with radial velocities distinct from the Milky Way. The C IV
absorption has a high covering fraction of $78\%$ within $25~{}\mathrm{kpc}$,
but between $25~{}\mathrm{kpc}<\rho_{\mathrm{LMC}}<35~{}\mathrm{kpc}$ the
covering fraction is $30$%. Fig. 1 shows two maps of the Magellanic System
superimposed on 21-cm H I emission maps of neutral hydrogen [13], with our
sightline locations color-coded by C IV column density and mean velocity. We
limit our analysis of the high-ion absorption to components with the following
properties: (1) high-ion column densities not explained by photoionization;
(2) velocities $v_{\mathrm{LSR}}>150$ km s-1 to select Magellanic gas and
avoid Milky Way contamination [14, 15, 16]; (3) velocities not associated with
known intermediate- and high-velocity clouds [15, 17].

Figure 1: The Magellanic system in an orthographic projection. The maps are
centered on the LMC and color-coded by column-density (a) and velocity (b).
Magellanic 21-cm H I emission is shown in a blue-scale integrated in
velocities to encompass the Magellanic System [9] with colored symbols showing
_HST_ /COS sightlines color-coded by C IV column densities. Upper limits are
shown using open symbols. The grey background shows Galactic 21-cm H I
emission from HI4PI [13] integrated between $-75<v_{\mathrm{LSR}}<+75$ km s-1.
Panel b) shows the mean velocity of H I with our sightlines color-coded by
mean C IV absorption velocity. Dotted circles mark the LMC impact parameter
and the south galactic pole (SGP) is marked with a white star.

Photoionization accounts for the low-ion absorbers (singly and doubly ionized
species) with temperatures $\log_{10}(T_{e}/K)=4.02^{+0.07}_{-0.04}$,
densities $\log_{10}(n_{e}/\mathrm{cm}^{-3})=-1.4\pm 0.3$, and line-of-sight
cloud sizes $\log_{10}(L/\mathrm{kpc})=-1.5^{+0.6}_{-0.4}$. The high-ion
absorbers have column densities too high to be explained by photoionization,
and instead are well explained by equilibrium or non-equilibrium (time-
dependent) collisional ionization models[18]. The C IV/Si IV column-density
ratio primarily yields a temperature of $T\sim 10^{4.9}~{}\mathrm{K}$, but
solutions as low as $T\sim 10^{4.3}~{}\mathrm{K}$ are possible in non-
equilibrium conditions for certain metallicities. However, the measured ratios
of O VI/C IV and O VI/Si IV yield higher temperatures. Our ionization modeling
and the component kinematics instead suggest that the O VI ions are tracing a
distinct and hotter phase of gas near $\sim 10^{5.5}$ K, where O VI peaks in
fractional abundance in collisional ionization [18].

We show the relation between the C IV, Si IV, and O VI column densities in the
Magellanic CGM and the LMC impact parameter in Fig. 2/A significant declining
profile is observed for C IV and Si IV, indicating that the gas content falls
off with radius, a characteristic signature of a diffuse CGM [19]. Sightlines
at small impact parameters of $\rho_{\mathrm{LMC}}<7~{}\mathrm{kpc}$ tend to
show a deficit of collisionally ionized high-ions compared to sightlines at
$7~{}\mathrm{kpc}<\rho_{\mathrm{LMC}}<12~{}\mathrm{kpc}$. These inner-CGM
absorbers are more susceptible to photoionization and winds, due to their
close proximity to the LMC. When only considering the absorbers at
$\rho_{\mathrm{LMC}}>7~{}\mathrm{kpc}$, the significance of the anti-
correlation between $N$(Si IV) and $\rho_{\mathrm{LMC}}$ becomes stronger. A
similar trend is seen with O VI, but is more uncertain owing to the small
sample size.

Figure 2: Declining radial profile of high-ion absorption. Sum total of O VI
(yellow triangles), C IV (pink circles) and Si IV (orange squares) column
density of Magellanic velocity absorbers that are not photoionized in each
sightline as a function of LMC impact parameter, $\rho_{\mathrm{LMC}}$, with
open symbols marking upper limits. The best-fit line and $68\%$ confidence
intervals for all data (dotted line) and only data at $\rho_{\mathrm{LMC}}>7$
kpc (solid line) are also shown (see Methods). Error bars represent standard
deviations.

For the $\sim 10^{4}$ K low-ion CGM phase, the relation between the modeled
ionized hydrogen column density ($N_{\rm H~{}{II}}$) and the impact parameter
($\rho_{\mathrm{LMC}}$) is shown in panel a) of Fig. 3.The ionized hydrogen
column densities for the $\sim 10^{4.9}$ K CGM phase traced by C IV and Si IV
and derived from both an equilibrium and non-equilibrium collisional
ionization model [18] are shown in panel b). Both the low- and high-ion gas
show similar radial profiles, and all but one sightline where high-ion
absorption is observed also show low-ion absorption. A linear-regression model
is used to find the ionized hydrogen mass of these two phases within
$\rho_{\mathrm{LMC}}<35~{}\mathrm{kpc}$, using a measured metallicity of
$\mathrm{[Z/H]}$ = $-$0.67 for the photoionized gas and an assumed metallicity
of $\mathrm{[Z/H]}$ = $-$1 for the high-ion gas. We find a total ionized
hydrogen mass in the $\sim 10^{4}$ K phase of
$\log_{10}(M_{\mathrm{H~{}{II},photo}}/M_{\odot})=8.7^{+0.2}_{-0.1}$. For the
$\sim 10^{4.9}$ K phase, we find a similar total ionized hydrogen mass of
$\log_{10}(M_{\mathrm{H~{}{II},\,Eq}}/M_{\odot})=8.5\pm 0.1$ from a
equilibrium model, while an isochoric non-equilibrium model results in
$\log_{10}(M_{\mathrm{H~{}{II},\,Non-Eq}}/M_{\odot})=8.6\pm 0.1$. If this
high-ion gas was instead at a lower temperature of $\sim 10^{4.3}$ K as is
possible in the isochoric model, the mass would increase by an order of
magnitude.

Figure 3: Total ionized hydrogen profile from ionization models. Ionized
hydrogen column density as a function of LMC impact parameter,
$\rho_{\mathrm{LMC}}$, for each of three gas phases at $\sim 10^{4}$ K
(photoionized; low ions; a), $\sim 10^{4.9}$ K (interfaces; C IV and Si IV;
b), and $\sim 10^{5.5}$ K (Corona; O VI; c). The photoionized results are from
_Cloudy_ models (blue X’s). Collisional ionization equilibrium (grey circles)
and non-equilibrium isochoric models (brown squares) are used for the $\sim
10^{4.9}$, and $\sim 10^{5.5}$ K gas based on C IV and O VI measurements,
respectively [18]. Linear regression fits and $68\%$ confidence intervals are
shown in the same color lines, with the $\sim 10^{4.9}$ K fit made only to
data at $\rho_{\mathrm{L}MC}>7$ kpc, and the $\sim 10^{4}$ K and $\sim
10^{5.5}$ K fits made to all data. The slope of the $\sim 10^{5.5}$ K is only
considered within the bounds of our observations. Error bars represent
standard deviations.

For the $\sim 10^{5.5}$ K CGM phase traced by O VI, we again assume a
metallicity of $\mathrm{[Z/H]}$ = $-$1 and derive an ionized hydrogen mass
within the range of LMC impact parameters where O VI is observed to be
$\log_{10}(M_{\mathrm{H~{}{II}}}/M_{\odot})\sim 8.3\pm 0.1$ for the
equilibrium and $\log_{10}(M_{\mathrm{H~{}{II}}}/M_{\odot})\sim 8.5\pm 0.1$
for the isochoric non-equilibrium model. Since we do not expect any of our
high-ion gas phases to maintain collisional ionization equilibrium, the non-
equilibrium case is more likely. Combined across phases, we estimate a total
ionized Magellanic CGM gas mass of
$\log_{10}(M_{\mathrm{H~{}{II},CGM}}/M_{\odot})\approx 9.1\pm 0.2$.

We consider three possible explanations for the C IV- and Si IV-bearing gas
around the LMC. It could exist in either a) a diffuse Magellanic Corona at
$\sim 10^{5}$ K, b) turbulent or conductive interfaces [20, 21] between cool,
low-ion gas clouds and a hotter diffuse Magellanic Corona at $\sim 10^{5.5}$
K, or c) turbulent or conductive interfaces between cool, low-ion gas clouds
and a hot ($\sim 10^{6}$ K) gaseous Milky Way halo. Fig. 4shows a cartoon
schematic of these three scenarios. When considering all our observations, we
conclude that the C IV exists in the interfaces between cooler $\sim 10^{4}$ K
clouds and a $\sim 10^{5.5}$ K Magellanic Corona (model b). This model
explains the high-ion radial profile, because there are more cool clouds
closer in to the LMC, each with interfaces tracking the kinematics of low-
ions. It also explains the presence of O VI, which shows velocity offsets and
directly traces the $\sim 10^{5.5}$ K Corona, but some of which may exist in
interfaces. The line widths are also consistent, with Si IV line widths that
are broader than the Si II line widths as expected with interfaces (though a
difference is not confirmed at high significance between C IV and C II line
widths). The alternative models are less favored because they either result in
a thermally unstable Corona with no explanation for O VI (model a) or cannot
explain the radial profile dependent on distance from the LMC, not distance
from the Milky Way (model c).

Figure 4: Summary of three possible C IV CGM models. Our three possible models
for C IV in a) a LMC Corona, b) interfaces with an LMC Corona, c) interfaces
with a Milky Way Corona. Cartoon renditions show $\sim 10^{4}$ K, $\sim
10^{5}$ K, $\sim 10^{5.5}$ K, and $\sim 10^{6}$ K gas in blue, orange, yellow,
and pink hues, respectively. The green dashed-line outlines model b) which
best explains our observations.

Several sightlines in our sample have been studied as part of a survey of the
Magellanic Stream[10], where some high-ion absorption may be attributed to
photoionization associated with a Seyfert flare in the Milky Way [22, 23] or
to a shock cascade [24]. However, the modeled Seyfert flare only affects gas
within a relatively small ionization cone outside of our sample. Gas
associated with the Stream may contribute to the absorption measured in
several other sightlines in our sample, but overall the tidally-stripped
Stream gas is less likely to be a major contributor to the Magellanic CGM
since it cannot easily explain the radial profile seen in both low- and high-
ions. The radial profile is also seen in sightlines away from the Stream,
further supporting that our observations are not biased by the tidally-
stripped gas. This radial profile appears truncated when compared to those
seen in a previous survey of the CGM of $43$ low mass, $z<0.1$ dwarf galaxies
(COS Dwarfs) [25] or the profile seen around the more massive Andromeda
Galaxy, M31 [26], though large uncertainties in virial radius estimates make
such a comparison difficult (see Extended Data Figure 2, panel b).

While some sightlines in our sample pass closer to the SMC than the LMC, the
common history of the two galaxies should result in a single enveloping
Magellanic Corona dominated by the LMC, which is $\sim 10$ times more massive
[8]. While in isolation, SMC-mass galaxies are not massive enough to host
their own warm coronae, they can host cool gas in their halos [25]. However,
the strong interactions with the LMC and the Milky Way would have strongly
disrupted such a cool halo during the SMC’s infall, and so a single Magellanic
Corona dominated by the LMC prevails.

The Magellanic Corona should be detectable via its dispersion measure induced
in radio observations of extragalactic fast radio bursts[27], since it
contains a high column density of free electrons. The presence of such a
pervasive corona around the LMC supports the picture of a hierarchical
evolution for the Local Group, in which the LMC and SMC accreted onto the MW
as part of a larger system of dwarf galaxies, a Magellanic Group[28, 29, 30],
not in isolation. Earlier work has detected the (ultra-faint) galaxies
associated with the Magellanic Group [31, 32]; our evidence for the Magellanic
CGM and Corona suggests we have now detected its gas, an important part of its
baryon budget. This provides a more complete understanding of the overlapping
and co-evolving ecosystems within the Local Group.

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## Methods

This work uses archival _HST_ /COS and _FUSE_ spectra to show evidence for the
Magellanic Corona. Here we describe our data reduction, Voigt profile fitting,
and ionization modeling methods. Throughout this work, all reported values and
uncertainties are medians and $68\%$ confidence intervals, unless otherwise
specified.

### Impact Parameters and Projection Effects

Unlike CGM studies of extragalactic systems, this work focuses on the CGM
surrounding the LMC at a distance of just $D=50$ kpc [1]. This proximity means
that background QSOs at large angular separations from the LMC, $\theta$,
correspond to relatively small physical separations. The impact parameter,
$\rho_{\mathrm{LMC}}$, is found using
$\rho_{\mathrm{LMC}}=D\sin\left(\theta\right)$. At $\theta>45$°, this
assumption no longer results in realistic impact parameter estimates, and we
would require a true 3D model of the location of gas absorbers to calculate a
physical separation between gas absorbers and the LMC. Additionally, for
sightlines at large $\theta$ it is harder to kinematically distinguish
absorption lines from the LMC and the Milky Way. A larger scale understanding
of the Magellanic Corona and multiphase CGM would only be possible with more
reliance on models and simulations to identify the 3D locations of gas
absorbers. To keep this work more focused on the observationally-derived
results, our analysis is thus strictly limited to sightlines within $45$° of
the LMC, corresponding to $35$ kpc in impact parameter.

###  _HST_ /COS Observations

We design our sample to consist of _HST_ /COS FUV observations of background
quasars using both the G130M and G160M gratings, covering the wavelength
ranges $\approx 1150-1450$ Å and $\approx 1405-1775$ Å, respectively. The
combination of these gratings enable us to examine the following absorption
lines: O I $\lambda$ 1302, N I $\lambda\lambda$ 1199, 1200, 1200.7, C II
$\lambda$ 1334, Al II $\lambda$ 1670, Si II $\lambda\lambda$ 1260, 1193, 1190,
1526, 1304, S II $\lambda\lambda$ 1250, 1253, 1259, Fe II $\lambda\lambda$
1608, 1144, Si III $\lambda$ 1206, C IV $\lambda\lambda$ 1548, 1550, and Si IV
$\lambda\lambda$ 1393, 1402. The COS spectra are processed following
previously developed custom reduction and wavelength calibration methods [10,
2] based on the raw products from the _calcos_[3] data reduction pipeline. In
order to remove geocoronal airglow contamination in O I $\lambda$ 1302 and Si
II $\lambda$ 1304, we use a second _calcos_ reduction of the data, using only
observations taken during orbital night-time.

The COS FUV observations have a native pixel size of
$2.5~{}\mathrm{km~{}s^{-1}}$ and a spectral resolution (FWHM) of $\approx
20~{}\mathrm{km~{}s^{-1}}$ and $\approx 15~{}\mathrm{km~{}s^{-1}}$ for G130M
and G160M spectra, respectively. We bin all spectra such that the resulting
spectra are Nyquist sampled with two pixels per resolution element.

###  _FUSE_ Observations

For $15$ sightlines in our sample, archival _FUSE_ spectra are also available
and analyzed to search for O VI $\lambda$ 1031, 1037 absorption. However, only
$6$ sightlines had high enough signal-to-noise to make a measurement. These
wavelengths fall on the _FUSE_ LIF1 channel with a spectral resolution of
$\sim 20~{}\mathrm{km~{}s^{-1}}$ and native pixel size of
$2~{}\mathrm{km~{}s^{-1}}$, which we bin to Nyquist sample with 2 pixels per
resolution element. These _FUSE_ data are reduced and aligned following the
customized methods similar to those used for the _HST_ /COS spectra [4, 5].
The O VI $\lambda$1031 may have contamination from molecular $H_{2}$
absorption at $\lambda$ 1032.356, which corresponds to $\sim 130$ km s-1 in
the O VI frame. However, the expected contribution from this contamination is
very small because of the high Galactic latitude of the sightlines, and in
most cases negligible.

### Absorption-Line Measurements

We use the open-source Python software, _VoigtFit_ [6], to perform Voigt
profile fitting of the absorption in several ions observed with _HST_ /COS
with the G130M and G160M gratings. This process uses a least-squares optimizer
[7] with recent atomic data [8, 9, 10] and convolves the Voigt profile with an
approximate instrumental profile of a Gaussian with FWHM corresponding to the
observed grating resolution. While this Gaussian approximation of the
instrumental profile is not an exact representation, it has been shown to have
a nearly negligible effect on fit results for weak high-velocity components
[23]. For all ion fits, we normalize the spectra using a third-order
polynomial fit to the continuum surrounding absorption lines of interest.
Regions of the absorption spectra that are contaminated by high-redshift
absorption components are then flagged to avoid fitting.

Source Name | RA | Dec | $\rho_{\mathrm{LMC}}$ | $v_{\mathrm{C~{}{IV}}}$ | $b_{\mathrm{C~{}{IV}}}$ | $\log_{10}(N_{\rm C~{}{IV}}/\mathrm{cm}^{-2})$ | $f_{\mathrm{PI,C~{}{IV}}}$ | $v_{\rm Si~{}{IV}}$ | $b_{\rm Si~{}{IV}}$ | $\log_{10}(N_{\rm Si~{}{IV}}/\mathrm{cm}^{-2})$ | $f_{\mathrm{PI,Si~{}{IV}}}$ | $\log_{10}(N_{\rm C~{}{IV}}/N_{\rm Si~{}{IV}})$
---|---|---|---|---|---|---|---|---|---|---|---|---
| [deg] | [deg] | [kpc] | [km s-1] | [km s-1] |  |  | [km s-1] | [km s-1] |  |  |
RX J0503.1-6634 | $75.77$ | $-66.56$ | 3.2 | $196\pm 18$ | $45\pm 27$ | $13.20\pm 0.31$ | $6.05\pm 4.30$ | $196\pm 18$ | $42\pm 20$ | $12.87\pm 0.22$ | $5.93\pm 3.05$ | $0.33\pm 0.38$
RX J0503.1-6634 | $75.77$ | $-66.56$ | 3.2 | $279\pm 6$ | $45\pm 12$ | $13.80\pm 0.10$ | $1.54\pm 0.35$ | $279\pm 6$ | $39\pm 15$ | $13.14\pm 0.17$ | $3.21\pm 1.29$ | $0.66\pm 0.20$
RX J0503.1-6634 | $75.77$ | $-66.56$ | 3.2 | $355\pm 8$ | $23\pm 13$ | $13.30\pm 0.25$ | $4.80\pm 2.75$ | $355\pm 8$ | $40\pm 11$ | $12.97\pm 0.14$ | $4.66\pm 1.47$ | $0.33\pm 0.28$
RX J0503.1-6634 | $75.77$ | $-66.56$ | 3.2 | $404\pm 9$ | $24\pm 11$ | $13.29\pm 0.19$ | $4.08\pm 1.78$ | — | — | $<12.29$ | — | $>1.01$
PKS0552-640 | $88.10$ | $-64.04$ | 5.5 | $397\pm 6$ | $40$ | $13.34\pm 0.26$ | $<0.01$ | $397\pm 6$ | $9$ | $12.45\pm 0.19$ | $0.03\pm 0.01$ | $0.89\pm 0.32$
PKS0552-640 | $88.10$ | $-64.04$ | 5.5 | $336\pm 3$ | $19\pm 5$ | $13.24\pm 0.15$ | $<0.01$ | — | — | $<12.29$ | — | $>0.96$
PKS0552-640 | $88.10$ | $-64.04$ | 5.5 | $262\pm 4$ | $32\pm 7$ | $13.24\pm 0.07$ | $<0.01$ | — | — | $<11.99$ | — | $>1.25$
PKS0552-640 | $88.10$ | $-64.04$ | 5.5 | $453\pm 13$ | $40\pm 14$ | $13.26\pm 0.24$ | $<0.01$ | — | — | $<12.28$ | — | $>0.98$
PKS0637-75 | $98.94$ | $-75.27$ | 6.7 | $230\pm 10$ | $10$ | $12.63\pm 0.56$ | $<0.01$ | $230\pm 10$ | $34$ | $11.94\pm 2.26$ | $<0.01$ | $0.69\pm 2.33$
PKS0637-75 | $98.94$ | $-75.27$ | 6.7 | $393\pm 8$ | $14$ | $12.91\pm 0.26$ | $<0.01$ | — | — | $<12.04$ | — | $>0.87$
PKS0637-75 | $98.94$ | $-75.27$ | 6.7 | $348\pm 5$ | $19\pm 9$ | $13.40\pm 0.20$ | $<0.01$ | $348\pm 5$ | $44\pm 25$ | $12.76\pm 0.26$ | $<0.01$ | $0.64\pm 0.33$
PKS0637-75 | $98.94$ | $-75.27$ | 6.7 | $290\pm 9$ | $37\pm 19$ | $13.52\pm 0.17$ | $<0.01$ | $290\pm 9$ | $43$ | $12.60\pm 0.68$ | $0.03\pm 0.04$ | $0.92\pm 0.70$
IRAS Z06229-6434 | $98.94$ | $-64.61$ | 6.7 | $448\pm 2$ | $27\pm 4$ | $13.53\pm 0.05$ | $<0.01$ | $448\pm 2$ | $16\pm 3$ | $12.79\pm 0.10$ | $<0.01$ | $0.74\pm 0.11$
IRAS Z06229-6434 | $98.94$ | $-64.61$ | 6.7 | $366\pm 10$ | $22$ | $13.22\pm 0.33$ | $<0.01$ | $366\pm 10$ | $23$ | $12.05\pm 0.91$ | $0.02\pm 0.04$ | $1.17\pm 0.96$
IRAS Z06229-6434 | $95.78$ | $-64.61$ | 6.7 | $306\pm 8$ | $35\pm 11$ | $13.33\pm 0.13$ | $<0.01$ | $306\pm 8$ | $30$ | $12.51\pm 0.45$ | $0.04\pm 0.05$ | $0.82\pm 0.47$
IRAS Z06229-6434 | $95.78$ | $-64.61$ | 6.7 | — | — | $<12.43$ | — | $262\pm 10$ | $10$ | $12.02\pm 0.67$ | $0.02\pm 0.03$ | $<0.41$
IRAS Z06229-6434 | $95.78$ | $-64.61$ | 6.7 | $394\pm 7$ | $11$ | $12.86\pm 0.66$ | $<0.01$ | $394\pm 7$ | $36$ | $12.20\pm 0.71$ | $<0.01$ | $0.66\pm 0.97$
UVQSJ045415.95-611626.6 | $73.57$ | $-64.61$ | 7.8 | $319\pm 5$ | $19\pm 11$ | $13.50\pm 0.57$ | $<0.01$ | $319\pm 5$ | $32\pm 6$ | $13.49\pm 0.17$ | $<0.01$ | $0.02\pm 0.59$
UVQSJ045415.95-611626.6 | $73.57$ | $-61.27$ | 7.8 | $394\pm 10$ | $28\pm 14$ | $13.12\pm 0.19$ | $<0.01$ | — | — | $<12.35$ | — | $>0.77$
UVQSJ045415.95-611626.6 | $73.57$ | $-61.27$ | 7.8 | $216\pm 5$ | $15\pm 8$ | $12.84\pm 0.15$ | $<0.01$ | — | — | $<12.39$ | — | $>0.46$
UVQSJ045415.95-611626.6 | $73.57$ | $-61.27$ | 7.8 | $302\pm 13$ | $38\pm 6$ | $13.88\pm 0.26$ | $<0.01$ | $302\pm 13$ | $40\pm 22$ | $12.96\pm 0.55$ | $0.02\pm 0.02$ | $0.92\pm 0.61$
RBS563 | $69.62$ | $-61.80$ | 8.0 | $311\pm 13$ | $24$ | $13.90\pm 0.59$ | $<0.01$ | $311\pm 13$ | $30\pm 16$ | $13.39\pm 0.54$ | $<0.01$ | $0.51\pm 0.80$
RBS563 | $69.62$ | $-61.80$ | 8.0 | $353\pm 26$ | $39\pm 18$ | $14.09\pm 0.38$ | $0.05\pm 0.04$ | $353\pm 26$ | $38\pm 16$ | $13.46\pm 0.42$ | $0.12\pm 0.12$ | $0.63\pm 0.56$
RBS563 | $69.62$ | $-61.80$ | 8.0 | $157\pm 37$ | $40$ | $13.21\pm 0.53$ | $<0.01$ | — | — | $<12.28$ | — | $>0.93$
RBS563 | $46.90$ | $-72.83$ | 8.0 | $253\pm 5$ | $20\pm 9$ | $13.70\pm 0.14$ | $<0.01$ | $253\pm 5$ | $27\pm 6$ | $13.30\pm 0.12$ | $<0.01$ | $0.40\pm 0.18$
ESO031-G08 | $46.90$ | $-72.83$ | 9.7 | $303\pm 12$ | $24\pm 9$ | $13.78\pm 0.27$ | $<0.01$ | $303\pm 12$ | $26\pm 11$ | $13.27\pm 0.44$ | $<0.01$ | $0.51\pm 0.52$
ESO031-G08 | $46.90$ | $-72.83$ | 9.7 | $262\pm 11$ | $19$ | $13.64\pm 0.36$ | $<0.01$ | $262\pm 11$ | $20$ | $12.95\pm 0.74$ | $<0.01$ | $0.69\pm 0.82$
ESO031-G08 | $40.79$ | $-72.28$ | 9.7 | $214\pm 18$ | $30$ | $13.41\pm 0.38$ | $<0.01$ | $214\pm 18$ | $27$ | $13.13\pm 0.70$ | $0.03\pm 0.05$ | $0.28\pm 0.80$
UKS0242-724 | $40.79$ | $-72.28$ | 11.3 | $194\pm 12$ | $24$ | $13.21\pm 1.01$ | $7.52\pm 17.44$ | $194\pm 12$ | $50\pm 8$ | $13.04\pm 0.18$ | $5.31\pm 2.16$ | $0.18\pm 1.02$
UKS0242-724 | $40.79$ | $-72.28$ | 11.3 | $253\pm 11$ | $50$ | $13.67\pm 0.74$ | $0.74\pm 1.27$ | $253\pm 11$ | $31\pm 17$ | $13.10\pm 0.28$ | $1.54\pm 0.99$ | $0.57\pm 0.79$
UKS0242-724 | $66.50$ | $-57.20$ | 11.3 | $312\pm 17$ | $36\pm 12$ | $13.41\pm 0.76$ | $<0.01$ | $312\pm 17$ | $31\pm 19$ | $12.86\pm 0.33$ | $<0.01$ | $0.55\pm 0.83$
1H0419-577 | $66.50$ | $-57.20$ | 12.1 | $311\pm 1$ | $24\pm 2$ | $13.55\pm 0.03$ | $<0.01$ | $311\pm 1$ | $14\pm 3$ | $12.73\pm 0.06$ | $0.03\pm 0.00$ | $0.82\pm 0.06$
1H0419-577 | $65.22$ | $-56.85$ | 12.1 | $356\pm 3$ | $20\pm 0$ | $13.09\pm 0.09$ | $0.02\pm 0.00$ | $356\pm 3$ | $29\pm 7$ | $12.65\pm 0.08$ | $0.05\pm 0.01$ | $0.43\pm 0.12$
HE0419-5657 | $65.22$ | $-56.85$ | 12.6 | $300\pm 4$ | $25\pm 4$ | $13.62\pm 0.10$ | $<0.01$ | $300\pm 4$ | $32\pm 11$ | $12.78\pm 0.12$ | $0.03\pm 0.01$ | $0.84\pm 0.15$
HE0419-5657 | $30.56$ | $-76.33$ | 12.6 | $360\pm 10$ | $35\pm 7$ | $13.50\pm 0.15$ | $<0.01$ | — | — | $<12.36$ | — | $>1.14$
PKS0202-76 | $30.56$ | $-76.33$ | 13.3 | $258\pm 6$ | $20\pm 8$ | $13.34\pm 0.14$ | $<0.01$ | $258\pm 6$ | $29\pm 16$ | $12.82\pm 0.23$ | $<0.01$ | $0.52\pm 0.27$
PKS0202-76 | $30.56$ | $-76.33$ | 13.3 | $294\pm 7$ | $16$ | $12.80\pm 0.48$ | $<0.01$ | — | — | $<12.45$ | — | $>0.35$
RBS567 | $69.91$ | $-53.19$ | 14.9 | — | — | $<12.68$ | — | — | — | $<12.28$ | — | —
HE0435-5304 | $69.21$ | $-52.98$ | 15.1 | $296\pm 16$ | $22$ | $13.31\pm 0.45$ | $<0.01$ | — | — | $<12.29$ | — | $>1.01$
HE0435-5304 | $69.21$ | $-52.98$ | 15.1 | $350\pm 24$ | $33$ | $13.34\pm 0.39$ | $<0.01$ | $350\pm 24$ | $24$ | $12.75\pm 0.74$ | $<0.01$ | $0.59\pm 0.83$
HE0435-5304 | $69.21$ | $-52.98$ | 15.1 | $250\pm 5$ | $25$ | $13.18\pm 0.32$ | $<0.01$ | $250\pm 5$ | $9$ | $12.67\pm 0.65$ | $<0.01$ | $0.51\pm 0.72$
HE0439-5254 | $70.05$ | $-52.80$ | 15.2 | $303\pm 2$ | $34\pm 6$ | $13.40\pm 0.06$ | $<0.01$ | $303\pm 2$ | $15\pm 6$ | $13.10\pm 0.08$ | $<0.01$ | $0.30\pm 0.10$
HE0439-5254 | $70.05$ | $-52.80$ | 15.2 | $243\pm 4$ | $15$ | $12.92\pm 0.16$ | $<0.01$ | — | — | $<12.28$ | — | $>0.64$
PKS0558-504 | $89.95$ | $-50.45$ | 16.9 | — | — | $<12.81$ | — | — | — | $<12.20$ | — | —
PKS0355-483 | $59.34$ | $-48.20$ | 20.3 | $203\pm 7$ | $44\pm 6$ | $13.24\pm 0.10$ | $<0.01$ | — | — | $<12.17$ | — | $>1.07$
FAIRALL9 | $20.94$ | $-58.81$ | 22.6 | $185\pm 3$ | $13\pm 8$ | $12.87\pm 0.29$ | $5.51\pm 3.66$ | $185\pm 3$ | $24\pm 5$ | $12.79\pm 0.06$ | $3.95\pm 0.57$ | $0.08\pm 0.29$
RBS1992 | $350.46$ | $-70.45$ | 23.4 | $244\pm 8$ | $33\pm 11$ | $13.27\pm 0.13$ | $0.02\pm 0.00$ | $236\pm 4$ | $12$ | $12.31\pm 0.14$ | $0.09\pm 0.03$ | $0.96\pm 0.19$
RBS1992 | $350.46$ | $-70.45$ | 23.4 | $173\pm 8$ | $32\pm 11$ | $13.27\pm 0.14$ | $<0.01$ | $186\pm 3$ | $20\pm 5$ | $12.72\pm 0.06$ | $0.05\pm 0.01$ | $0.55\pm 0.15$
HE0331-4112 | $53.28$ | $-41.03$ | 26.6 | — | — | $<12.50$ | — | — | — | $<12.21$ | — | —
RBS144 | $15.11$ | $-51.23$ | 28.6 | $181\pm 10$ | $44\pm 7$ | $13.11\pm 0.14$ | $<0.01$ | — | — | $<12.19$ | — | $>0.92$
HE0246-4101 | $42.03$ | $-40.81$ | 28.8 | — | — | $<12.98$ | — | $179\pm 6$ | $25\pm 10$ | $12.91\pm 0.09$ | $<0.01$ | $<0.08$
HE0153-4520 | $28.81$ | $-45.10$ | 29.0 | — | — | $<12.38$ | — | $197\pm 5$ | $12$ | $12.22\pm 0.14$ | $0.02\pm 0.01$ | $<0.16$
HE0226-4110 | $37.06$ | $-40.95$ | 29.8 | $203\pm 3$ | $12$ | $12.70\pm 0.36$ | $0.02\pm 0.01$ | $203\pm 3$ | $14\pm 5$ | $12.58\pm 0.07$ | $0.03\pm 0.01$ | $0.12\pm 0.36$
HE0226-4110 | $37.06$ | $-40.95$ | 29.8 | $160\pm 6$ | $33\pm 14$ | $13.20\pm 0.15$ | $<0.01$ | $160\pm 6$ | $11$ | $12.09\pm 0.21$ | $0.03\pm 0.02$ | $1.11\pm 0.26$
HE0038-5114 | $37.06$ | $-40.95$ | 30.0 | — | — | $<12.83$ | — | — | — | $<12.65$ | — | —
HE2336-5540 | $354.81$ | $-55.40$ | 30.9 | — | — | $<12.71$ | — | — | — | $<12.37$ | — | —
HE0003-5023 | $1.43$ | $-50.12$ | 30.9 | — | — | $<12.54$ | — | — | — | $<12.07$ | — | —
IRAS F21325-6237 | $324.09$ | $-62.40$ | 32.5 | $170\pm 1$ | $14\pm 2$ | $13.32\pm 0.04$ | $<0.01$ | $170\pm 1$ | $9\pm 6$ | $12.17\pm 0.13$ | $0.02\pm 0.01$ | $1.15\pm 0.14$
HE2305-5315 | $347.16$ | $-52.98$ | 33.9 | — | — | $<12.50$ | — | — | — | $<12.97$ | — | —

Extended Data Table 1: Voigt profile model parameters for the Magellanic CGM.
Best fit model parameters for C IV and Si IV for components associated with
Magellanic CGM with $v_{\mathrm{LSR}}>150$ km s-1. For each sightline, the LMC
impact parameter $\rho$, centroid velocity $v$, Doppler parameter (linewidth;
$b$), column density log${}_{10}(N/\mathrm{cm}^{-2})$, and photoionized
fraction $f_{\mathrm{PI}}$
($f_{\mathrm{PI}}=N_{\mathrm{Cloudy}}/N_{\mathrm{Obs}}$) for both ions are
given. Components with $f_{\mathrm{PI}}\leq 0.1$ are considered as not
photoionized and shown in Fig. 3.Linewidths from data with low S/N spectra or
that are fixed in the fitting process do not show errors. Uncertainties in
this table correspond to $1\sigma$ standard deviations. Additional fit results
for all ions and sightlines can be viewed on
%\ref{fig:3}.%****␣main.tex␣Line␣1425␣****https://github.com/Deech08/HST_MagellanicCorona.

We first fit the absorption in all low and intermediate ions (O I, N I, C II,
C II*, S II, Si II, Si III, Al II, Fe II) simultaneously, allowing component
line centers to be tied across ions when they show general agreement. The C
II* line always contaminates the measurement of absorption in C II at $+250$
km s-1. When there are blended C II components at this velocity, we fix the C
II* column density to a constant value of $10^{13.8}$ cm-2, based on average
measurements from previous work [11], but in these cases the measured C II
columns near $+250$ km s-1 are not used in our analysis. The Si III $\lambda$
1206 transition is frequently saturated, requiring the linewidths to be tied
to match the fit Si II linewidths. A minimum allowed linewidth of $9$ km s-1
is applied based on the instrumental resolution, and maximum linewidths are
only added as a constraint for highly blended components if they are needed to
converge to a best fit.

Source Name | RA | Dec | $\rho_{\mathrm{LMC}}$ | $v_{\mathrm{O~{}{VI}}}$ | $b_{\mathrm{O~{}{VI}}}$ | $\log_{10}(N_{\rm O~{}{VI}}/\mathrm{cm}^{-2})$
---|---|---|---|---|---|---
| [deg] | [deg] | [kpc] | [km s-1] | [km s-1] |
IRAS Z06229-6434 | 95.78 | -64.61 | 6.7 | $231\pm 12$ | $40$ | $14.07\pm 0.14$
IRAS Z06229-6434 | 95.78 | -64.61 | 6.7 | $352\pm 9$ | $40$ | $14.30\pm 0.11$
IRAS Z06229-6434 | 95.78 | -64.61 | 6.7 | $466\pm 11$ | $36\pm 16$ | $14.07\pm 0.16$
ESO031-G08 | 46.90 | -72.83 | 9.7 | $231\pm 10$ | $35$ | $14.17\pm 0.11$
ESO031-G08 | 46.90 | -72.83 | 9.7 | $308\pm 5$ | $25$ | $14.33\pm 0.10$
1H0419-577 | 66.50 | -57.20 | 12.1 | $302\pm 7$ | $16\pm 8$ | $14.03\pm 0.13$
1H0419-577 | 66.50 | -57.20 | 12.1 | $358\pm 11$ | $28\pm 18$ | $13.90\pm 0.22$
RBS144 | 15.11 | -51.23 | 28.6 | $199\pm 9$ | $32$ | $13.88\pm 0.12$
HE0226-4110 | 37.06 | -40.95 | 29.8 | $180\pm 8$ | $26$ | $13.72\pm 0.13$
IRAS F21325-6237 | 324.09 | -62.40 | 32.5 | $165\pm 6$ | $31\pm 10$ | $14.04\pm 0.10$

Extended Data Table 2: Voigt profile model parameters for O VI. Best fit model
parameters for O VI for components associated with Magellanic CGM with
$v_{\mathrm{LSR}}>150$ km s-1. Linewidths without errors are shown when they
are fixed in the fitting process. Additional fit results for all ions and
sightlines can be viewed on https://github.com/Deech08/HST_MagellanicCorona.

C IV and Si IV are then fit simultaneously following the same procedure, but
independent of the low-ion results to avoid biasing our analysis, since the
high-ion component structure may be different. O VI absorption from _FUSE_ is
also fit independently when data are available and a reasonable continuum can
be determined. If the component structure of the low and high-ions match, they
are flagged after the fitting process so that their column densities,
linewidths, and line centers can be compared in the subsequent steps.
Additionally, we calculate upper limits of any transitions where absorption is
not seen based on the S/N of the observed spectra [12, 13]. Lastly, fit
components attributed to the Milky Way or known intermediate- or high-velocity
clouds are flagged to avoid contaminating our analysis. We note that some
contamination from fixed pattern noise persists in our reduced spectra, which
may impact our measured column densities and is not accounted for in our
estimated errors.

In total, across 28 sightlines, we initially identify $112$ unique velocity
components that may be attributed to the Magellanic System. We then impose a
velocity threshold and only consider absorbers at $v_{\mathrm{LSR}}>150$ km
s-1 to avoid contamination from absorbers associated with the Milky Way [14].
The 150 km s-1 velocity threshold was determined using a combination of the
observed component velocities and simulations of the Magellanic System [7]; it
represents the value that best separates the Galactic and Magellanic
components and is consistent with previous kinematic studies of Magellanic
absorption [14, 16]. Additionally, this velocity threshold is supported by
dynamical arguments: given the LMC mass, the virial theorem predicts that
Magellanic gas has a velocity dispersion of $50$ km s-1 centered on the LMC
velocity of $280$ km s-1, implying that $95\%$ of Magellanic gas should be
within $180$ km s-1 and $380$ km s-1. As a result, our final sample has $52$
unique Magellanic velocity components that are further analyzed based on their
kinematics and photoionization modeling. The Voigt-profile model parameters
for these $52$ C IV and Si IV components are given in Extended Data Table 1and
the $10$ unique Magellanic O VI absorption components are shown in Extended
Data Table 2.Extended Data Figure 1.shows our measured C IV$\lambda$ 1548 and
O VI$\lambda$ 1031 absorption-line spectra for our sample. Panel a) of
Extended Data Figure 2shows the total measured _HST_ /COS column densities in
a number of low and high ions from the Magellanic absorbers at
$v_{\mathrm{LSR}}>150$ km s-1 as a function of LMC impact parameter. All low-
ions show a declining radial profile, similar to the relation shown in the
high-ions (Fig. 2).

A comparison of our observed radial profile to that seen in the COS Dwarfs
survey [25] and M31 [26] is shown in panel b) of Extended Data Figure 2.We
normalize impact parameter measurements across these surveys based on the
radius enclosing a mean overdensity of 200 times the critical density,
$R_{200}$, which is often used as a measure of the virial radius in CGM
studies. In the radial region of overlap between these surveys and our work,
the declining profile of the LMC is more concentrated, with a possibly
truncated profile. Because the LMC halo is already within the virial radius of
the MW, it is expected to be tidally truncated, hence such a truncated profile
is expected. However, the uncertainties in estimates of $R_{200}$ are
estimated to be $50\%$ in the COS Dwarfs and M31 surveys, with the LMC value
we use at $R_{200}=115\pm 15$ kpc.

Extended Data Figure 1: Sample absorption-line spectra. Normalized _HST_ /COS
spectra of the C IV $\lambda$ 1548 (upper panels) and O VI $\lambda$ 1031
(lower panels, red outlines) absorption lines in the LSR velocity frame,
ordered by their LMC impact parameters (low to high). The normalized flux is
shown in black, with 1$\sigma$ uncertainties shaded in gray around them. The
solid teal line shows the full Voigt-profile fit to the C IV $\lambda$ 1548,
1550 and O VI $\lambda$ 1031, 1037 doublets, with individual Magellanic
components shaded in red-orange hues, corresponding to the same color scheme
used in Fig. 1. Component centers are shown with tick marks. The dashed
vertical line marks the $150$ km s-1 threshold used in this work. Grayed-out
portions of the spectra are contaminated by higher-redshift absorbers and are
not considered in our analysis. C IV $\lambda$ 1550 spectra are shown in cases
where the C IV $\lambda$ 1548 is highly contaminated. Additional fit results
for all ions and sightlines can be viewed on
https://github.com/Deech08/HST_MagellanicCorona.

Our spectra can be accessed publicly on the Barbara A. Mikulski Archive for
Space Telescopes (MAST). A full table of our fit parameters, along with
summary plots of our best fits can be accessed on GitHub at
https://github.com/Deech08/HST_MagellanicCorona.

Extended Data Figure 2: Radial profiles of ionic column densities. Measured
column densities of Magellanic components at $v_{\mathrm{LSR}}>150$ km s-1 as
a function of LMC impact parameter, $\rho_{\mathrm{LMC}}$, with open circles
marking upper limits in panel a). For each ion, the best-fit line and
uncertainties are found using a MCMC analysis with censoring, as in Fig. 3,
for all impact parameter for low-ions and only $\rho_{\mathrm{LMC}}>7$ kpc for
C IV and Si IV. Each panel corresponds to the ion labeled in the lower left.
Panel b) shows a comparison of C IV column densities as a function of impact
parameter normalized by the radius enclosing a mean overdensity of 200 times
the critical density, $R_{200}$, of the host galaxy with the COS Dwarfs survey
[25] and M31 [26]. The LMC data use $R_{200}=115\pm 15$ kpc.

### Ionization Models

We use 1D _Cloudy_ [15] radiative transfer models to simulate the physical
conditions of the absorbing gas. Our _Cloudy_ models require four key inputs
in order to run: (1) an external radiation field, (2) the observed column
density measurements, (3) a specified stopping condition to reach for
convergence, and (4) a gas-phase metallicity. All models assume a plane-
parallel geometry and constant gas density.

Incident radiation fields in _Cloudy_ require a shape and intensity. We adopt
the Milky Way escaping radiation field model to set the shape of the radiation
field, assuming the radiation field from the LMC and SMC have the same
spectral shape[16, 17, 10]. The intensity of the radiation field toward each
sightline is set by a hydrogen ionizing photon flux $\Phi_{\rm H}$ determined
from published ionization models, which includes contributions from the LMC,
SMC and Milky Way [22]. We reconstruct this model in 3-dimensional space to
interpolate an initial value for $\Phi_{\rm H}$ for any specified location. In
our model, we allow $\Phi_{\rm H}$ to be a free parameter, since a precise
distance to the absorbing material is not known. We also include a constant
contribution from an extragalactic UV background [18], and cosmic ray
background [19].

We use _Cloudy_ ’s built-in _optimize_ command to vary our free parameters and
find optimal parameters to explain our observed column densities and upper
limits [20, 15]. The optimize models use up to three possible free parameters:
(1) the hydrogen ionizing photon flux, $\Phi_{\rm H}$, described above, (2)
the total hydrogen number density $n_{\rm H}$, which is the sum of the ionic,
atomic, and molecular hydrogen densities of the plasma that is to be modeled,
and (3) the neutral hydrogen column density ($N_{\mathrm{H~{}{I}}}$) stopping
condition. For sightlines with an H I or O I detection, the observed H I or O
I column density measurement serves as the stopping condition and the model
only utilizes the first two free parameters ($\Phi_{\rm H}$, and $n_{\rm H}$).
For sightlines without an H I or O I, we utilize all three free parameters
($\Phi_{\rm H}$, and $n_{\rm H}$, $N_{\mathrm{H~{}{I}}}$). Once _Cloudy_ ’s
_optimize_ method has found a possible solution of parameters, we run one
final _Cloudy_ model at the specified optimal parameters to produce
predictions of ion column densities and gas temperatures, including
predictions for high-ion (Si IV, C IV, O VI) column densities. To ensure
_Cloudy_ does not settle at local minima in the optimization process, we use a
broad range of initial densities from $\log_{10}(n_{\rm
H}/\mathrm{cm}^{-3})=-3$ to 1, and ionizing fluxes $\Phi_{\rm H}$ in a range
of $3$ dex around the model prediction at $D=50$ kpc, but still find a
resulting narrow range of total hydrogen densities ($n_{\rm H}$), ionized gas
temperatures $(T_{e})$, neutral atomic hydrogen columns $(N_{H~{}{I}})$, and
ionized-to-neutral atomic hydrogen ratios $N$(H II)/$N$(H I), across all
sightlines and velocity components. Additionally, we have also run a coarse
grid at a larger range of free parameters to help confirm that our solutions
are indeed optimal, and not local minima.

While ISM gas-phase metallicities have been measured in the LMC, SMC,
Magellanic Bridge [21], and Magellanic Stream [22, 23, 24, 25], the
metallicity of the Magellanic CGM is highly uncertain. In order to estimate
the gas metallicity, we use a sightline in our sample towards HE0226-4110 that
overlaps with recently published analysis of _FUSE_ spectra to measure neutral
hydrogen column densities [26]. Two absorption components toward this
sightline may belong to the Magellanic Corona at
$v_{\mathrm{LSR}}=+174~{}\mathrm{km~{}s^{-1}}$ and
$+202~{}\mathrm{km~{}s^{-1}}$, providing a measured neutral hydrogen column
density to set as a stopping condition in _Cloudy_. Unfortunately, there is no
detected O I absorption in either the COS or _FUSE_ data, so a metallicity is
calculated using a _Cloudy_ optimize model (described above), allowing the
total hydrogen density, hydrogen ionizing photon flux, and metallicity to
vary. The _Cloudy_ models are optimized based on the measured COS column
densities across all available metal ions and any upper limits when absorption
is not detected. The results for these two components are
$\log_{10}(\Phi_{\rm{H}}/\mathrm{photons}~{}\mathrm{s}^{-1})=5.06$,
$\log_{10}(n_{\rm{H}}/\mathrm{cm}^{-3})=-1.58$, and $\mathrm{[Z/H]}=-0.72$ for
the $v_{\mathrm{LSR}}=+174~{}\mathrm{km~{}s^{-1}}$ component and
$\log_{10}(\Phi_{\rm{H}}/\mathrm{photons}~{}\mathrm{s}^{-1})=4.95$,
$\log_{10}(n_{\rm{H}}/\mathrm{cm}^{-3})=-1.91$, and $\mathrm{[Z/H]}=-0.62$ for
the $v_{\mathrm{LSR}}=+202~{}\mathrm{km~{}s^{-1}}$ component. Based on these
results, we adopt the average $\mathrm{[Z/H]}=-0.67$ as the gas-phase
metallicity for photoionized gas. For hotter gas in interfaces and the corona,
we assume a gas-phase metallicity of $\mathrm{[Z/H]}=-1$, since we expect this
more primordial gas to be at lower metallicity.

Our optimal set of _Cloudy_ models provide predictions for the expected column
densities of the high-ions Si IV, C IV, and O VI for a single-phase
photoionized gas. However, the observed high-ion columns are significantly
greater (by orders of magnitude) than the photoionization predictions. Across
all sightlines and absorption components that may be associated with the
Magellanic System, we find that $72\%$ of Si IV and $84\%$ of C IV absorption
components are under $10\%$ photoionized. We use this $10\%$
($1~{}\mathrm{dex}$) threshold to define our sample of Magellanic absorbers
that are not photoionized (see shaded components in Fig. 2). These C IV and Si
IV absorbers likely arise in interfaces in the range
$T=10^{4.3-4.9}~{}\mathrm{K}$.

Extended Data Figure 3: Collisional ionization model relations. Predicted
column-density ratios as a function of temperature based on the radiative
cooling gas models in collisional ionization equilibrium (CIE; black; panel
a), time-dependent isochoric cooling (brown; panel b) or time-dependent
isobaric cooling (green; panel b) [18]. Dotted pink lines show observed values
of column density ratios in the Magellanic components. Solid, dash, dash-dot,
and dotted lines correspond to metallicities of $[Z/H]=0$, $-1$, $-2$, and
$-3$, respectively.

The observed triply-ionized Magellanic absorption is well described using
either equilibrium or time-dependent non-equilibrium collisional ionization
models [18]. In both cases, we can infer an electron temperature based on the
ratio of C IV and Si IV column densities, because the close similarity of the
C IV and Si IV line profiles indicates the two ions are co-spatial. The
modeled relation of this column-density ratio with temperature for the
equilibrium model and for isobaric and isochoric time-dependent models is
shown in Extended Data Figure 3for a range in metallicities. The inferred
temperature is then used to determine a C IV ionization fraction, from which
the total ionized hydrogen (H II) column density can be calculated, resulting
in the measurements shown in Fig. 3(middle panel). In total, the temperature
distributions of the photoionized and collisionally ionized gas are shown in
Extended Data Figure 4(lower-left panel). In the sightlines where we have
measured O VI absorption, we find that the O VI absorbing gas requires a
higher temperature than the C IV and Si IV absorbing gas, indicating that the
O VI arises in a separate, hotter phase. While at high metallicity, lower-
temperature solutions for our observed column density ratios are possible,
this is not the case at the lower metallicities (below 0.1 solar) expected for
Magellanic Coronal gas.

Extended Data Figure 4: Evidence for a multi-phase Magellanic CGM. Panel a)
shows histograms and Gaussian kernel density estimates (KDEs) of the residuals
between predicted ion column densities from _Cloudy_ photoionization models
and measured ion column densities, normalized by their standard deviation
uncertainties. The histograms are shown for low-ions (C II, Si II, and Al II;
dashed lines in blue, yellow, and green, respectively) and high-ions (C IV, Si
IV; solid lines in pink, orange). Panel b) shows the inferred gas temperatures
for the photoionized gas (blue; $\log_{10}(T_{e}/K)=4.02^{+0.07}_{-0.04}$) and
the collisionally ionized gas under an equilibrium (grey;
$\log_{10}(T_{e}/K)=4.92^{+0.05}_{-0.02}$) and non-equilibrium isochoric
(brown; $\log_{10}(T_{e}/K)=4.87^{+0.09}_{-0.06}$; high temperature solution
only) model [18]. Panel c) and d) show measured linewidths for C II and C IV
absorption (upper; c) and Si II and Si IV absorption (lower; d).

We also consider more recent collisional ionization models that include
photoionization from an extragalactic background [27]. However, these models
do not include the non-isotropic radiation fields necessary for modeling
clouds near the Milky Way and LMC, and only offer approximate predictions
using a general background radiation field. Instead, we only consider the two
cases of entirely photoionized or entirely collisionally ionized in this work,
but note that a full picture will require considering collisional ionization
and photoioniation from the Milky Way and Magellanic Clouds together.

### Statistical Significance of Results

Here we describe the statistical tests we used to support our claims of
significance. Throughout this work, we adopt a significance threshold p-value
of $p=0.05$.

Extended Data Figure 5: Kinematic differences of paired high and low ions.
Histograms and Gaussian KDEs of pairs of high $-$ low-ion linewidths (panel a)
and velocity centroids (panel b), separated by C ions (pink; solid line) and
Si ions (orange; dotted line). The velocity centroid panel also includes
differences between O VI and either Si III (grey; dashed line) or C IV (green;
dot-dashed line). The paired linewidths undergo a one-sided Wilcoxon signed-
rank test, while the velocity centroids undergo a two-sided test.

#### Velocity Structure

In our Voigt-profile fitting process, individual components are initially
paired across low- and high-ions based on their approximate centroid
velocities. This pairing process is inherently biased as it assumes components
across ions are physically tied and results in the lowest possible differences
in velocity centroids for our analysis. However, for the low- and high-ions,
the velocity structure was qualitatively well matched to one another, with
absorption components at similar velocities for both cases. This
correspondence is less clear for the O VI absorption line centroids, so
matching O VI components in the same manner is much more uncertain. Combined
with the relatively low S/N ($\approx$10) and moderate velocity resolution (20
km s-1) of our spectra, we are unable to fully resolve all absorption
components. We therefore find that comparisons of the kinematic properties of
low- and high-ions are generally inconclusive. However, the kinematics are
still consistent with our primary conclusion that C IV and Si IV arise in the
interfaces between cool clouds and a Magellanic Corona, because in an
interface model the velocity structure of the low-ions and high-ions should be
linked. When considering O VI, we calculate the velocity offset from the
closest absorption component in other ions (Si III or C IV), and find the
widths of the velocity offset distributions have standard deviations of
$\sigma_{\mathrm{O~{}{VI}-Si~{}{III}}}=22^{+7}_{-4}$ km s-1 and
$\sigma_{\mathrm{O~{}{VI}-C~{}{IV}}}=22^{+10}_{-6}$ km s-1, respectively. This
is $7^{+12}_{-8}$ km s-1 greater in width of the distribution of velocity
differences between the low-ions and C IV matched in the same manner,
supporting the result that O VI exists in a different phase.

#### Linewidths

We show the paired (matched based on their velocities during the Voigt profile
fitting process) differences of component line widths in panel a) of Extended
Data Figure 5.Differences in paired linewidths do not show statistical
significance. However, when consider our populations of line width
measurements, we do find a statistically significant difference between the
line width distributions of singly ionized C and Si, in comparison with triply
ionized C and Si (see their distributions in panels c) and d) of Extended Data
Figure 4.The Anderson-Darling statistical test of the null hypothesis that the
singly and triply ionized linewidths are drawn from the same underlying
population can generally be rejected at the p-value threshold of $0.05$ for
both C and Si. We perform the test on $10,000$ bootstrap samples to account
for measurement errors of linewidths. The C IV and C II linewidths return a
p-value (with $68\%$ confidence intervals) of
$p_{\mathrm{C}}=0.008^{+0.08}_{-0.007}$, with $78\%$ of bootstrap samples
below our p-value threshold of 0.05. Similarly, the Si IV and Si II linewidths
return p-values of $p_{\mathrm{Si}}=0.001^{+0.01}_{-0.0}$, with $93\%$ of
bootstrap samples below our significance threshold.

Extended Data Figure 6: Statistical significance of radial profile. Results of
a correlation analysis. Kendall’s Tau correlation test with censoring [28] is
performed using $10,000$ bootstrap samples, with the resulting distributions
of test statistics (panel a) and $p$-values (panel b) shown for the null
hypothesis of no correlation. The dotted line shows the Gaussian kernel
density estimate for the distribution using all data points, while the solid
line and shaded histograms show the result for only data at
$\rho_{\mathrm{LMC}}>7$ kpc.

#### Declining Radial Profile

We test the statistical significance of the anti-correlation between the C IV
and Si IV with LMC impact parameter using Kendall’s Tau rank correlation
coefficient with censoring, which provides a robust measure of the monotonic
relationship between two variables [29, 28]. The Magellanic Corona shows a
distribution of coefficients that are negative for both C IV and Si IV, with
mean values of $\tau=-0.4\pm 0.1$ and $\tau=-0.3\pm 0.1$, respectively, as
shown in Extended Data Figure 6.The $p$-values for C IV allow the null
hypothesis of no correlation to be rejected at the $0.05$ level for $97\%$ of
bootstrap samples, whereas the $p$-values for Si IV can only be rejected for
$73\%$ when considering all of our data. When only considering the absorbers
at $\rho_{\mathrm{LMC}}>7~{}\mathrm{kpc}$, the significance of the Si IV anti-
correlation becomes stronger with $p<0.05$ for $89\%$ of $10,000$ bootstrap
samples and a mean value of $\tau=-0.4\pm 0.1$, but the change for C IV is
negligible. The best-fit lines for the anti-correlation are found using an
MCMC analysis with censoring to account for upper limits and measurement
errors [30]. For the O VI measurements, Kendall’s Tau rank correlation
coefficient is less reliable since we only have 6 data points, and is not
conclusive.

### Magellanic Corona vs. Tidally Stripped Stream with Interfaces

Previous simulations have been able to explain much of the ionized gas
associated with the Magellanic Stream by tidal stripping, without the presence
of a Corona [31]. If this were the case, and the Stream were the dominant
source of ionized gas, we would expect to see a stronger correlation of C IV
column density as a function of distance from the Magellanic Stream (absolute
Magellanic Stream latitude) than as a function of LMC impact parameter. We use
the partial Spearman rank-order correlation test to assess the strength of the
correlation between our measured ion column densities and either the LMC
impact parameter or the absolute Magellanic Stream latitude, while removing
the effects of the other. We note that for this test, we are only considering
the colissionally ionized C IV and Si IV columns, while considering all the
observed columns for low-ions. The correlation coefficients and $p$-values of
the test with null hypothesis of no correlation are given in Extended Data
Table 3.For most ions, the correlation is significantly stronger with LMC
impact parameter, after removing the effects of the absolute Magellanic Stream
latitude. However, the partial correlation test for Fe II is inconclusive and
the test for O I suggests a stronger correlation with absolute Magellanic
Stream latitude. These tests are consistent with a Magellanic Corona and CGM
origin to the gas absorbers we have measured, with the exception of O I, which
may be more biased towards tracing cooler, tidally stripped gas in the
Magellanic Stream.

Ion | # of Detections | $r_{N\,\rho_{\mathrm{LMC}};\left|B_{\mathrm{MS}}\right|}$ |  $p$-value | $r_{N\,\left|B_{\mathrm{MS}}\right|;\rho_{\mathrm{LMC}}}$ |  $p$-value
---|---|---|---|---|---
C IV | 17 |  $-$0.793 |  $<$0.001 | 0.323 | 0.222
Si IV | 17 |  $-$0.524 | 0.037 | 0.224 | 0.362
Si III | 22 |  $-$0.664 | 0.001 |  $-$0.158 | 0.494
Si II | 21 |  $-$0.583 | 0.007 |  $-$0.328 | 0.159
C II | 20 |  $-$0.525 | 0.021 |  $-$0.076 | 0.758
Al II | 18 |  $-$0.534 | 0.027 |  $-$0.468 | 0.058
Fe II | 12 |  $-$0.143 | 0.676 |  $-$0.561 | 0.072
O I | 13 |  $-$0.117 | 0.717 |  $-$0.693 | 0.012

Extended Data Table 3: Partial Spearman rank-order correlation tests. Tests of
the relation between ion column density and LMC impact parameter, after
removing the effects of absolute Magellanic Stream latitude
($r_{N\,\rho_{\mathrm{LMC}};\left|B_{\mathrm{MS}}\right|}$) and between ion
column density and absolute Magellanic Stream latitude, after removing the
effects of LMC impact parameter
($r_{N\,\left|B_{\mathrm{MS}}\right|;\rho_{\mathrm{LMC}}}$). $p$-values for
each partial test are shown in boldface if they are greater than the 0.05
threshold, not allowing the null hypothesis of no correlation to be rejected.

In Extended Data Figure 7,we show our measurements of collisionally ionized C
IV columns on a map of the Magellanic System in Magellanic Coordinates,
alongside measurements of all C IV absorption from a previous survey of the
Magellanic Stream[10]. When considering all C IV, the surface density profile
is much more extended along the direction of the Magellanic Stream, but with
our adopted velocity threshold and removal of photoionized gas, the radial
profile centered on the LMC is apparent, especially when considering
sightlines that overlap on our sample and the previous sample.

Extended Data Figure 7: Magellanic Corona sample vs. Magellanic Stream sample.
Map of measured C IV column density shown in Magellanic Coordinates of our
collisionally ionized Magellanic sample (panel a; photoionized phase removed)
and previous measurements of all C IV absorption in a survey of the Magellanic
Stream[10] (panel b). Square and triangle symbols in the lower panel mark
sightlines that were flagged as belonging to the Magellanic Bridge and Stream,
respectively, while other sightlines are marked with a plus symbol. Open
symbols show upper limits. Symbols with a dotted white outline in the upper
panel denote sightlines that have overlap with the previous survey[10], while
those without a white outline are unique to this work.

In this previous work, much of the observed C IV absorption was interpreted to
arise from interfaces around the tidally stripped, cooler gas from the LMC
with a hot $\sim 10^{6}$ K Milky Way Corona. The basic premise of this
conclusion is still valid in our sample, but the strong radial profile
centered on the LMC suggests that the hotter gas interacting to form the
interfaces should also be centered on the LMC, not the Milky Way. Therefore, a
Magellanic Corona at $\sim 10^{5.5}$ K can explain our observed radial profile
and the observed C IV absorption.

### Mass Estimates

Our estimates of the mass for each phase of the Magellanic CGM are derived
from the relation between ionized hydrogen column density and LMC impact
parameter. For each phase ($\sim 10^{4}$ K, $\sim 10^{4.9}$ K, $\sim 10^{5.5}$
K), a best-fit linear regression model is fit to the ionized hydrogen column
as a function of $\rho_{\mathrm{LMC}}$. Then the ionized hydrogen mass in each
phase is calculated using

$M_{\rm{H~{}{II}}}=\int_{0~{}\mathrm{kpc}}^{35~{}\mathrm{kpc}}~{}N_{\mathrm{H~{}{II}}}\left(\rho_{\mathrm{LMC}}\right)m_{p}~{}2\pi\rho_{\mathrm{LMC}}~{}f_{\mathrm{cov}}~{}d\rho_{\mathrm{LMC}}$
(1)

where $m_{p}$ is the proton mass, and $f_{\mathrm{cov}}$ is the covering
fraction.

For the $\sim 10^{4}$ K gas, the ionized hydrogen column density in each
direction is derived directly from the _Cloudy_ models, with a covering
fraction $f_{\mathrm{cov}}=0.82$, since low-ions are detected at Magellanic
velocities in $23/28$ directions in our sample. However, we note that the
covering fraction of low-ions tends to decrease as a function of LMC impact
parameter, but use a constant covering fraction as an approximation.

For the $10^{4.9}$ K gas, the total ionized hydrogen column density in each
sightline is derived based on the C IV column density and best-fit temperature
from the collisional ionization models [18] using

$N_{\rm{H~{}{II}}}=\frac{N_{\rm{C~{}{IV}}}}{f_{\rm{C~{}{IV}}}\left(\rm{Z/H}\right)},$
(2)

where $f_{\rm{C~{}{IV}}}\equiv\rm{C^{3+}/C}$ is the fraction of triply ionized
Carbon at the best-fit temperature, and
$\left(\rm{Z/H}\right)=0.21\left(\rm{Z/H}\right)_{\odot}$ is the metallicity.
For C IV, the covering fraction is set to $f_{\mathrm{cov}}=0.78$ for
$\rho_{\mathrm{LMC}}<30~{}\mathrm{kpc}$ and $f_{\mathrm{cov}}=0.3$ for
$\rho_{\mathrm{LMC}}>=30~{}\mathrm{kpc}$ based on the observed detection rate
of C IV absorption in our sample. The relation between the derived column
densities and LMC impact parameter allow for our mass calculations for the
$\sim 10^{4}$ K and $\sim 10^{4.9}$ K gas to converge, changing by at most
$0.1$ dex if instead we integrate out to $500$ kpc.

For the $\sim 10^{5.5}$ K gas, the mass is again found based on the O VI
absorption columns in the collisional models, using

$N_{\rm{H~{}{II}}}=\frac{N_{\rm{O~{}{VI}}}}{f_{\rm{O~{}{VI}}}\left(\rm{Z/H}\right)},$
(3)

and using the same covering fraction correction used for the $\sim 10^{4.9}$ K
gas. Here, we use the maximal $f_{\rm{O~{}{VI}}}$ value for each of the
collisional models, which peak near $10^{5.5}$ K at $f_{\rm{O~{}{VI}}}\sim
0.2$. The best-fit line for this phase does not converge, so integrating the
radial profile depends on highly on the radial range considered. Instead, we
only integrate between the bounds of our observations
($6.7~{}\mathrm{kpc}<\rho_{\mathrm{LMC}}<32.5~{}\mathrm{kpc}$) and present an
approximate Corona mass for this region only.

Data Availability: _HST_ /COS spectra used in this work are publicly available
on MAST at https://archive.stsci.edu/. These archival observations can be
found under the following HST Program IDs: 11692, 15163, 12263, 11686, 11520,
12604, 12936, 11541, and 14655.

Code Availability: Voigt-profile fit results, summary fit spectra, and custom
code used can be found in our GitHub repository at
https://github.com/Deech08/HST_MagellanicCorona. Additionally, the following
software was used in this work. _astropy_[32, 33], _calcos_[3], _cartopy_[34],
_lmfit_[7], _scipy_[35], _voigtfit_[6], _Cloudy_[15], _pingouin_[36]

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Acknowledgments: This work was funded through HST Archival Program 16363,
provided by NASA through a grant from the Space Telescope Science Institute,
which is operated by the Association of Universities for Research in
Astronomy, Inc., under NASA contract NAS5-26555. This work uses observations
made with the NASA/ESA Hubble Space Telescope, obtained from the Data Archive
at the Space Telescope Science Institute, which is operated by the Association
of Universities for Research in Astronomy, Inc., under NASA contract
NAS5-26555. D.K. is supported by an NSF Astronomy and Astrophysics
Postdoctoral Fellowship under award AST-2102490.

Author Contributions: D.K. led the investigation, formal analysis,
methodology, visualization, and writing. A.J.F. led the project development
and management and contributed heavily to the project inception, funding, and
writing. E.D.O. led the conceptualization and is Principal Investigator of the
_Hubble Space Telescope_ grant that funded the research. B.W. led the data
curation and contributed heavily to the proposal writing. A.J.F., E.D.O, and
B.W. contributed to validation, methodology, and reviewing and editing. D.M.F
contributed towards data curation. F.H.C contributed towards methodology and
visualization. S.L, F.H.C, D.M.F, J.C.H, and N.L contributed towards
validation and reviewing and editing.

Competing Interests: The authors declare no competing interests.

Additional Information: Correspondence and requests for materials should be
addressed to D.K. (email: dkrishnarao@coloradocollege.edu).

Reprints and permissions information is available at www.nature.com/reprints.