The Diffuse Ionized Gas Halo of the Large Magellanic Cloud

All observations are processed using the WHAM pipeline described in Haffner et al. (2003). The spectra are preprocessed in the pipeline to remove cosmic rays before combining. The Fabry–Perot spectrometer produces circular interference patterns, which are summed in annuli to produce a linear spectrum that is a function of velocity. The pipeline then bins the spectra in v_bin = 2 km s⁻¹ intervals. Once binned, the spectra are normalized to exposure time and then scaled according to the airmass of the observation. The pipeline then uses an intensity correction factor to account for the sensitivity degradation of the WHAM instrumentation that occurs over time. The standardization of the spectra allows us to directly average observations across multiple nights and varying exposure times.

We follow the same methods for atmospheric line subtraction as outlined in SHB19 and previous WHAM studies (Haffner et al. 2003; Hill et al. 2009; Barger et al. 2012, 2013; Ciampa et al. 2021). However, we observed Hα emission at higher positive geocentric velocities (v_geo) than these previous WHAM surveys. To characterize the atmospheric lines at these higher velocities, we observed two Hα-faint positions on the sky located at (ℓ, b) = (890, − 710) and (ℓ, b) = (2725, − 280). We combined 24 observations of these two sightlines taken on 2016 September 3 to create the atmospheric template in Figure 4. The faint atmospheric emission lines are well-fit by the Gaussian components listed in Table 1 and have a strength that scales with airmass.

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In addition to removing the faint atmospheric lines listed in Table 1 with an atmospheric template, we separately removed a bright OH atmospheric line that lies at v_geo = + 272.44 km s⁻¹ (as seen in the top panel of Figure 4) as its strength scales with the flux of sunlight on the Earth's upper atmosphere in addition to the airmass. The strength of the OH line compared to the Hα emission allows us to directly fit and subtract the line from our observations.

Unfortunately, some of the faint Hα emission associated with the LMC can be inadvertently removed when the OH line is subtracted. The affected region spans roughly 20 km s⁻¹ in width (see Figure 9), wider than the unresolved core of the OH line (≈12 km s⁻¹) due to non-Gaussian, broader wings in the instrument profile. Some of the brighter and/or broad astronomical emission can be disentangled from the OH line, but the full scope of missing emission can not be estimated well in this work using Hα alone. Future observations of associated emission lines at other wavelengths (e.g., [S ii] and [N ii]) will help reconstruct Hα component structure, when necessary.

Unlike many previous WHAM studies, we used a different Hα filter centered at longer wavelengths to capture the most positive Magellanic velocities with high transmission. However, this shift also introduces an unfortunate out-of-tune Fabry–Pérot "ghost" from the OH doublet at 6577.183 Å/6577.386 Å. Although these lines are at a much higher velocity (v_LSR ≈ + 662 km s⁻¹) than our target range, the combination of the filter and etalon pressure tune introduces a broad emission feature that spans roughly +120 ≲ v_LSR ≲ + 180 km s⁻¹ (seen as the rise in the left edge of Figure 4). To avoid this feature, any sightlines that have data using only this filter are integrated over velocities greater than v_LSR = + 180 km s⁻¹ for measurements in this work. However, sightlines combined with observations from BHB13 and CBL20 utilize the standard WHAM Hα filter (Haffner et al. 2003) to avoid this contamination. These spectroscopic observations can be fully integrated down to v_LSR = + 150 km s⁻¹, our preferred low-velocity limit for this study.

Finally, we use the fitting algorithm described in CBL20 to identify systematic residuals present at specific geocentric velocities after removing the atmospheric lines, as described above. Sightlines with little to no Galactic and Magellanic emission are then used to correct these residuals from all spectra in a data set.

Foreground dust from the MW lies between us and the LMC. Therefore, we need to correct the Hα intensity for any attenuation associated with the MW. Using the average foreground R_V value of 3.1 as outlined in Cardelli et al. (1989), we follow the method described in Section 3.4.1 of BHB13 and apply:

$$\begin{eqnarray}&&{I}_{{\rm{H}}\alpha ,\mathrm{corr}}={I}_{{\rm{H}}\alpha ,\mathrm{obs}}{e}^{A({\rm{H}}\alpha )/2.5},\end{eqnarray} \tag{ 1 }$$

where

$$\begin{eqnarray}&&A({\rm{H}}\alpha )=5.14\times {10}^{-22}\left\langle {N}_{{\rm{H}}\,{\rm\small{I}}}\right\rangle {\mathrm{cm}}^{2}\ {\mathrm{atoms}}^{-1}\ \mathrm{mag},\end{eqnarray} \tag{ 2 }$$

and $\left\langle {N}_{{\rm{H}}\,{\rm\small{I}}}\right\rangle$ is a measure of the average H i column density over the region of interest.

In this work, the extinction correction is only applied during the mass calculations in Section 5. For foreground MW extinction, we integrate H i spectra from the HI4PI survey (Ben Bekhti et al. 2016) over −100 ≤ v_LSR ≤ + 100 km s⁻¹. We calculated the extinction correction using two methods. The first uses the mean H i for the entire region of interest. The second determines the correction for 025 gridded pixels used in a specific mass calculation method (see Section 5.1 for details). In general, we find the range of extinction corrections from these methods increases the observed I_Hα by an average of 10%–18% for the main body of the LMC and the LA.

To calculate the internal extinction correction for the LMC, we use values from Gordon et al. (2003)'s Table 2 and Table 4. Using Table 4, we calculate a A(λ)/A(V) for λ = 0.656 μm using their results for wavelengths between 2.198 μm ≥ λ ≥ 0.440 μ m. Using a third-order polynomial, we have

$$\begin{eqnarray}&&A(\lambda )/A(V)=0.2536\lambda -0.5623{\lambda }^{2}+0.8468{\lambda }^{3}-0.1724,\end{eqnarray} \tag{ 3 }$$

where λ is in μm. At Hα we then have A(Hα)/A(V) = 0.7492 ± 0.0012. We can combine this scaling with N(H i)/A(V) from Gordon et al. (2003)'s Table 2 to estimate A(Hα). This gives us the equation:

$$\begin{eqnarray}&&A\left({\rm{H}}\alpha \right)=2.31\times {10}^{-22}\lt{N}_{{\rm{H}}\,{\rm{I}}}\gt{\mathrm{cm}}^{2}\ {\mathrm{atoms}}^{-1}\ \mathrm{mag}.\end{eqnarray} \tag{ 4 }$$

We only apply this correction to regions near the center of the LMC that have integrated N_{H I} > 1.0 × 10²⁰ cm⁻² over +150 ≤ v_LSR ≤ + 390 km s⁻¹. In these regions, I_Hα > 0.5 R and the average correction is 16%. Since we do not know the location of the Hα emitting gas along the line-of-sight depth with respect to the dust, these extinction-corrected intensities provide an upper limit on Hα emission. Due to the low H i column density of the LA, we do not apply a correction for it. Our foreground and internal extinction values are summarized in Table 2.

Table 2. Extinction

Region	Foreground Extinction			Internal Extinction
	log $\left\langle {N}_{{\rm{H}}\,{\rm\small{I}}}\right\rangle$	A(Hα)	%corr a	log $\left\langle {N}_{{\rm{H}}\,{\rm\small{I}}}\right\rangle$	A(Hα)	%corr
LMC	20.7	0.27	10%	21.2	0.43	16%
Leading Arm 1.1	21.0	0.46, 0.48	17%, 18%	⋯	⋯	⋯

Note.

^a We list two values that represent the effective extinction correction applied for mass calculations in Sections 5.1.1 and 5.1.2 and for those in Section 5.1.3. These methods use different processes for determining $\left\langle {N}_{{\rm{H}}\,{\rm\small{I}}}\right\rangle$, and the average value from each method is displayed. Only foreground extinction for the LA 1.1 region had any appreciable difference in extinction correction when using the different methods.

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3.4.1. LMC Velocity Frame

Due to the large spatial extent of the LMC across the sky, it is useful to adjust velocities to a reference frame with respect to the center of the LMC. In this frame, we can more easily identify extended features that form continuous kinematic structures that are not modified by our local observing geometry. We use the method described in CBL20 to transform the velocity frame of our data set from LSR to the Large Magellanic Cloud Standard of Rest (LMCSR) given by:

$$\begin{eqnarray}&&\displaystyle \frac{{\rm{\Delta }}{v}_{\mathrm{LMCSR}}}{\mathrm{km}\,{{\rm{s}}}^{-1}}=262.55-3.25\left(l-280^\circ \right)+3.66\left(b-(-33^\circ )\right),\end{eqnarray} \tag{ 5 }$$

where l and b are Galactic longitude and latitude in degrees. Using Equation (5), v_LMCSR = v_LSR − Δv_LMCSR.

In Figure 5, we present our non-extinction-corrected total Hα intensity and H i column density maps of the LMC. These maps are integrated over +150 ≤ v_LSR ≤ + 390 km s⁻¹. The Hα data set is sensitive to I_Hα ⋍ 10 mR for the extended emission structures. The H i data set was extracted from the HI4PI survey, which has a theoretical 5σ detection limit of N_HI = 2.3 × 10¹⁸ cm⁻² (Ben Bekhti et al. 2016).

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We compare log(Hα) to log(N_{H I}) in Figure 6. We only plot sightlines where the Hα emission is above 23 mR $\left(1\sigma \right)$ and the H i column density is above 10¹⁹ cm⁻². At $\mathrm{log}({I}_{{\rm{H}}\alpha })\lt -0.3$, there does not appear to be an obvious relationship between the Hα intensity and the H i column density. However, above that cutoff, there does appear to be a correlation. We tested both ranges of points with Kendall's Tau correlation test and found that τ = 0.52 for I_Hα ≥ − 0.3 R and τ = 0.33 for I_Hα ≥ 0.3 R, with the significance level much lower than 0.05. We fit two different slopes to the data using these two data cutoffs:

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$\mathrm{log}\left(\tfrac{{I}_{{\rm{H}}\alpha }}{R}\right)\geqslant -0.3$:

$$\begin{eqnarray}&&\mathrm{log}\left(\displaystyle \frac{{N}_{{\rm{H}}\,{\rm\small{I}}}}{{\mathrm{cm}}^{-2}}\right)=20.58+0.43\times \mathrm{log}\left(\displaystyle \frac{{I}_{{\rm{H}}\alpha }}{R}\right)\end{eqnarray} \tag{ 6 }$$

$\mathrm{log}\left(\tfrac{{I}_{{\rm{H}}\alpha }}{R}\right)\geqslant 0.3$:

$$\begin{eqnarray}&&\mathrm{log}\left(\displaystyle \frac{{N}_{{\rm{H}}\,{\rm\small{I}}}}{{\mathrm{cm}}^{-2}}\right)=20.76+0.28\times \mathrm{log}\left(\displaystyle \frac{{I}_{{\rm{H}}\alpha }}{R}\right).\end{eqnarray} \tag{ 7 }$$

These two fits are shown in Figure 6. For I_Hα ≥ − 0.3 R, 11% of the points fall outside Equation (6) when including an error range of±0.5. For I_Hα ≥ 0.3 R, only 2% of the points fall outside Equation (7) when including an error range of±0.5. The lines were only fit to the central point of each observation and did not include the errors due to the uneven nature of the error bars in log space. Above I_Hα = 0.3 R, the errors in both Hα and H i are on average 1.3% of the total I_Hα intensity, and would have had a minor effect on the overall fit.

Along with the gas within the LMC, there appears to be an ionized counterpart to the neutral LA (boxed region in Figure 2). This emission is more pronounced when viewing the gas in the v_LMCSR frame (Figure 7). The LA appears between −10 ≤ v_LMCSR ≤ + 70 km s⁻¹ in H i and −60 ≤ v_LMCSR ≤ + 30 km s⁻¹ in Hα. The spectra from the regions marked in Figure 8 are presented in Figure 9. The strongest Hα emission lies adjacent to an extended region with N_{H I} ≳ 0.5 × 10¹⁹ cm⁻². Neutral gas is spatially coincident with the ionized hydrogen gas at column densities from ${N}_{{\rm{H}}\,{\rm\small{I}}}=\left(0.1\mbox{--}0.5\right)\times {10}^{19}\,{\mathrm{cm}}^{-2}$ but may not be physically related to each other due to the observed distinct velocities of the two gas components in many locations.

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When viewed in the LMCSR frame of reference, the velocity profile of the H i and Hα weakly trace each other in the center of the galaxy, around H i contours of N_{H I} = 10²⁰ cm⁻² (Figure 10). The velocity profile of the Hα emission leading to the LA appears smoothly related to the velocity profiles of the neutral LMC gas. The Hα gas does not appear to have the extended high-velocity gas that extends out toward the LA. Instead, the lower velocity gas seems dominant in these regions. This difference can be clearly seen in the first moment maps. Unfortunately, the higher velocity is coincident with the OH line at v_geo = + 272.44 km s⁻¹ in this region, which prevents us from easily detecting faint, higher velocity Hα emission from the LA, which coincides with the OH line.

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In Figure 11, we compare the mean velocity of the Hα emission to mean(v_{H I})–mean(v_Hα ) of all points with both an I_Hα ≥0.2 R and N_{H I}≥10¹⁸ cm⁻². We calculate the mean velocity taking the first moment of the spectra between +170 km s⁻¹ and +390 km s⁻¹. We see 28% of the Hα emission has a positive mean velocity, while 72% has a negative mean velocity. Where the H i has a positive mean velocity, the Hα mean velocity has a wider spread. The velocity comparison appears to be weighted toward a more negative mean difference in velocity. This may be biased by the (CBL20) observations, which observed the lower velocity LMC winds. Figure 12 highlights six regions where the absolute difference in velocity between the Hα and H i components is greater than 70 km s⁻¹. These regions appear on the outer edge of the LMC. The difference in mean velocity may be due to multicomponent emission in the region, as seen in Figures 12(b) and (d). In other regions, such as Figure 12(e), the Hα and H i components may be related; however, the broader nature of the Hα emission may shift the mean velocity or the OH line may obscure the true mean velocity of the Hα emission.

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Since the removal of the bright OH line can remove Hα emission, we tested how this could bias the Hα mean velocity. Adding an extra component to reduced spectra with a velocity width of 15 km s⁻¹, a center coincident with the OH line, and with intensities of I_Hα = 0.10 and 0.23 R, we found a maximum positive mean velocity shift of 7 km s⁻¹ for 0.10 R, and 8 km s⁻¹ for 0.23 R. Additionally, even with the shift, the mean H i velocities in the LA are still significantly positive. Any emission underneath the OH line would not shift the Hα mean velocity to values similar to H i mean velocities that are significantly positive.

The separation in velocity space of the ionized and neutral gas may indicate that, despite the spatial coincidence from our perspective, the gas components are physically separated. This could indicate separate origins for the ionized and neutral gas or additional mechanisms affecting different regions of gas. A multicomponent analysis in a future work may reveal a more accurate relationship between the neutral and ionized gas components.

In this work, we estimate the mass of the LMC in three different ways. As the electron density and the line-of-sight depths of the gas is unknown, we must make generalized assumptions about the distribution of the ionized gas in relation to the neutral components.

We first describe a mass estimate where H i measurements from HI4PI are used to constrain the electron density (Section 5.1.1). This assumes the ionized gas exists in a skin around the neutral gas.

We then describe two different mass estimates where we assume the neutral and ionized gas are well mixed and share the same line-of-sight depth. These two scenarios make different assumptions about the geometry of the galaxy. Our second method assumes a cylindrical geometry, which treats the gas as a uniform, cylindrical slab we view from the top down.

The third method assumes the gas within the LMC exists in an ellipsoid centered at the LMC's kinematic center.

In each of these methods, the line-of-sight depth of the LMC's gas components is not known. We use the maximum line-of-sight depth of the stellar component measured in the SMASH survey (7 kpc (Choi et al. 2018)) as a proxy (Sections 5.1.1, 5.1.2 and 5.1.3) to estimate the upper limit of the assumed line-of-sight depths for either the neutral or the hydrogen gas. We also use a 3 kpc depth, which assumes the gas component has a smaller line-of-sight depth than the stellar component. For the LA, we constrain the line-of-sight depth to 1.5 kpc. The assumption of 1.5 kpc comes from the observed spatial extent of the H i features in the LA area on the plane of the sky. It is reasonable to expect that the features have comparable extents along the line of sight as they do in the plane of the sky.

For each method described below, we focus on two regions, the LMC and the LA (Figure 2). The LA region is defined as LA 1.1 in Venzmer et al. (2012). We use a modified version of their bounds, (ℓ, b) = (295, −23.83), (281.78, −25.51), (283.29, −32.41), (295, −29.51), to accommodate the boundaries of our observations. The area covered by the LMC and the LA regions is 171.3 and 78.5 degrees² respectively. All mass calculations were determined by integrating emission over +150 ≤ v_LSR ≤ + 390 km s⁻¹. Any directions with Hα intensities below 25 mR are excluded when calculating the mass using individual sightlines. To match the HI4PI observations to the WHAM beams, we take the HI4PI column density observations contained in each beam and average them together. We adopt a distance to the LMC of D = 50 kpc (Walker 2012; Pietrzyński et al. 2013), although the ellipsoid scenario in Section 5.1.3 uses the inclination, i = 2586 (Choi et al. 2018), to vary D. Values for each region can be found in Table 3. The resulting atomic gas ratios, M_ionized/M_{ionized+neutral}, are listed in Table 4.

Table 3. Neutral and Ionized Properties

Region		Neutral Properties
		Ω	$\mathrm{log}\left\langle {N}_{{\rm{H}}\,{\rm\small{I}}}\right\rangle$ a	${M}_{{{\rm{H}}}^{0}}$ b	${M}_{{{\rm{H}}}^{0}}$ c	${L}_{{{\rm{H}}}^{0}}$	$\mathrm{log}\left\langle {n}_{0}\right\rangle$
		deg2	(cm−2)	(106 M☉)	(106 M☉)	(kpc)	(cm−3)
LMC		171.3–170.0	20.7	850	760	3.0–7.0 d	−1.6–−1.2
Leading Arm 1.1		78.5–79.8	20.0	71	68	1.5	−1.6
Region		Ionized Properties
		Ionized Skin				Cylinder			Ellipsoid
	$\mathrm{log}\left\langle {EM}\right\rangle$	$\mathrm{log}\left\langle {n}_{e}\right\rangle$ e	${L}_{{{\rm{H}}}^{+}}$	${M}_{{{\rm{H}}}^{+}}$	$\mathrm{log}\left\langle {n}_{e}\right\rangle$	${L}_{{{\rm{H}}}^{+}}$	${M}_{{{\rm{H}}}^{+}}$	$\mathrm{log}\left\langle {n}_{e}\right\rangle$ d	${L}_{{{\rm{H}}}^{+}}$	${M}_{{{\rm{H}}}^{+}}$
	(cm−6 pc)	(cm−3)	(kpc)	(106 M☉)	(cm−3)	(kpc)	(106 M☉)	(cm−3)	(kpc)	(106 M☉)
LMC	1.4–1.4	−1.6	5.1	1600	−1.1–−1.2	3.0–7.0	1200–1800	−0.9–−1.1	3.0–7.0 f	600–1000
Leading Arm 1.1 g	−0.2–−0.3	−1.6	0.8	54	−1.7	1.5	62	−1.7	1.5	12–46

Notes.

^a Average H i column density for the region calculated from HI4PI. ^b Mass range calculated using the ionized skin (Section 5.1.1) and cylinder (Section 5.1.2) scenarios. ^c Mass range calculated using the ellipsoid scenario (Section 5.1.3) only. ^d n_e is calculated over the respective line of sight for each scenario. ^e Only uses an assumed line of sight for $\mathrm{log}\left\langle {n}_{0}\right\rangle$ of 3 kpc for the LMC. ^f Maximum line-of-sight depth at the center of the LMC for the ellipsoid scenario. 3 kpc and 7 kpc maximum depths are from Choi et al. (2018). The minimum depth used is 1.5 kpc. ^g Mass range calculated with an emission cutoff of I_Hα = 25 mR and 10 mR.

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5.1.1.  n_e = n₀

5.1.2. Cylindrical Geometry

5.1.3. Ellipsoidal Geometry

The ellipsoid scenario is similar to our cylindrical scenario; however, instead of averaging the region we apply the mass calculation and the extinction correction to individual 025 pixels, allowing a direct comparison between H i and Hα emission in each location. In this scenario, we model the shape of the LMC as a simple ellipsoid, which is then used to estimate the line-of-sight depth of the gas. We centered our ellipsoid around the H i kinematic center at (ℓ, b) = (2798, − 335) from Kim et al. (1998). The projected ellipse can be seen in Figure 2, with a semimajor axis of 7 kpc and a semiminor axis of 6 kpc, and a position angle θ = 149.23 degrees (Choi et al. 2018). We then assume that the maximum line-of-sight depth of the ellipse is similar to the stellar line-of-sight depths or is half the maximum depth (7 kpc or 3 kpc). We then use vary the line-of-sight depth as we move away from the center.

In both cases, if we simply use the equation for an ellipsoid the outer regions of the LMC would fall to sub-kpc line-of-sight depths. We chose to constrain the minimum depth for the model to be no less than 1.5 kpc. In addition to the varying line of sight, we vary the distance, D from Equation (10). The distance to the center of the LMC is assumed to be D = 50 at the center of the galaxy. In this scenario we allow the inclination of the galaxy to change D according to the following equation:

$$\begin{eqnarray}&&D=50\,\mathrm{kpc}+d\times \tan \left(\theta \right)\end{eqnarray} \tag{ 11 }$$

where θ = 25° and d is the positive or negative distance of the pixel in kpc from the center of the LMC along the semimajor axis, with the positive direction toward the NE side of the galaxy and the negative toward the SW. The inclusion of inclination in the model had a ∼1% effect on the mass, due to the inclination having the largest effect on the regions furthest from the center of the galaxy where there is less gas and the symmetry of the inclination minimizing its overall effect.

Extinction corrections are estimated for each pixel in a 025 grid. We also limit our mass calculations to sightlines with Hα emission above I_Hα > 25 mR, except in the case of the LA where we also include observations down to 10 mR to account for the significant amount of continuous faint emission in the region. Using this model, we find ionized hydrogen that is traced by the Hα emission in the LMC to be M_ionized ≈ 6 × 10⁸ M_☉ using a depth of 3 kpc and M_ionized ≈ 10 × 10⁸ M_☉ using a depth of 7 kpc.

In the LA, the depth is constrained to a flat slab with a depth of 1.5 kpc and the distance is kept at 50 kpc similar to the cylindrical calculations. However, we use the individual sightline calculations instead of the averaged intensity and the same intensity cutoffs we use in the ellipsoidal scenario. The resulting ionized gas mass for the LA is $\left(12\mbox{--}46\right)\times {10}^{6}\,{M}_{\odot }$ for 25 mR and 10 mR intensity cutoffs.

Each of these scenarios treats the gas as if it is homogeneously distributed. They ignore the possibility of a clumpy WIM or a diffuse halo with a thick, dense disk of ionized gas. This likely leads to an overestimate of the total ionized gas mass. Thus, our estimates should be treated as upper limits to the total gas mass and ionization fraction. Future work will investigate how varying the gas distribution affects the total mass calculated.

LA 1.1 (see Figure 2) is the closest region of a larger section of the LA system, LA 1 (Nidever et al. 2008; Venzmer et al. 2012). Each sightline marked in Figure 8 and Figure 9 contains H i and Hα emission spanning similar velocities, including several regions where both the H i and Hα appear to have multiple components. The first moment map in Hα in this region appears to be weighted toward lower velocities compared to the H i gas in the region (Figure 10 and Figure 11). However, the higher velocities can overlap with the bright OH sky line where faint Hα emission is difficult to trace. As a result, it is difficult to determine if the Hα emission is genuinely multicomponent or a single, broad feature.

The Cosmic Origins Spectrograph (COS)/Ultraviolet and Visual Echelle Spectrograph (UVES) survey contains one absorption target present in the LA 1.1 region. PKS0637 lies within the region defined as LA 1.1 by Venzmer et al. (2012). Their modeling finds an ionization fraction of 59% along the line of sight coincident with $\mathrm{log}\left({N}_{{\rm{H}}\,{\rm\small{I}}}/\,{\mathrm{cm}}^{-2}\right)=19.29$. Although our determination of the ionization fraction is less constrained for the LA (16%–56%), it is consistent with their estimate. Along b = − 30° H i column density falls of quickly from N_{H I} ≥8 × 10²⁰ cm⁻² down to N_{H I} ≤10²⁰ cm⁻². In contrast, Hα emission appears to connect continuously from the LMC to the LA, both spatially and in velocity.

In Nidever et al. (2008), the radial velocities of the LA complexes were measured and found to be similar to those of the LMC. They initially identified gaps between the separate components of the LA 1; however, deeper observations by the H i Parkes All-Sky Survey (HIPASS) show the region to be continuous in neutral gas. In the Hα emission, we have identified what appears to be a continuous filament of ionized gas that extends out toward the more distant H i components of LA 1. The filament of gas lies at latitudes b > − 29° where $\mathrm{log}\left({N}_{{\rm{H}}\,{\rm\small{I}}}/{\mathrm{cm}}^{-2}\right)\lt 20$ (Figure 8).

Contrary to Nidever et al. (2008), a study by Putman et al. (1998) argues for the SMC as an origin to LA 1. This claim is based on an H i feature originating from the SMC and smoothly extending to the base of LA 1.1 in velocity space (see their Figures 1 and 2). In addition, Staveley-Smith et al. (2003) suggest that the gas from the LMC is merely leaking into the LA and not a dominant contributor.

While our results for the kinematics and spatial extent of the ionized gas in this region match gas velocities from the LMC, the spatial extent of our observations is limited. However, in our observed region, the velocity and spatial extent of the gas appear to connect smoothly with the LMC, suggesting the LMC is the origin of the ionized component of LA 1.1 Additionally, some sightlines (Figure 10) with velocities that match the Bridge region extend above the LMC. This may indicate some material within the region is associated with the SMC or Bridge. Further observations of the full extent of the LA to find Hα emission present in the gaps in H i column density will test the hypothesis that LA 1 is a continuous structure that originated from LMC.

The results of this LMC study extend our WHAM survey of the Magellanic System. SHB19 found a gas ionization fraction of 42%–47% associated with the central region of the SMC and an ionized gas mass range of ${M}_{\mathrm{ionized}}\approx \left(6\mbox{--}7.5\right)\times {10}^{8}\,{M}_{\odot }$ and a neutral gas mass range of ${M}_{\mathrm{ionized}}\approx \left(8.4\mbox{--}8.7\right)\times {10}^{8}$. Our results listed in Table 4 show a significantly higher ionization fraction for the LMC using the ionized skin and cylinder method. Using the ellipsoidal method, we find a similar ionization fraction found for the SMC when assuming a maximum thickness of 3 kpc.

Like the SMC, the LMC ionized gas follows neutral gas rotation within the center of the galaxy where $\mathrm{log}\left({N}_{{\rm{H}}\,{\rm\small{I}}}/{\mathrm{cm}}^{-2}\right)\gt 20.7$, though it appears to be a weak correlation. Similar to the SMC, there appear to be regions of gas that follow the velocity trends seen in the center of the galaxy and extend out into the halo of the galaxy. The LMC has the LA feature, and the SMC has an ionized filament that extends out of the galaxy, which smoothly connects to the central rotating ionized gas in velocity space (SHB19).

Similar to the two galaxies, the Bridge also appears significantly ionized. The ionized gas mass for the Bridge ranges from ${M}_{\mathrm{ionized}}=\left(0.7\mbox{--}1.7\right)\times {10}^{8}\,{M}_{\odot }$ with a neutral gas mass of M_neutral = 3.3 × 10⁸ M_☉ (BHB13). Like the ionized gas in the SMC and LMC, the Hα emission associated with the Bridge extends beyond the boundaries of the neutral gas.

Combined, these three studies have discovered a total of ${M}_{\mathrm{ionized}}=\left(1.3\mbox{--}2.6\right)\times {10}^{9}\,{M}_{\odot }$ in diffuse ionized gas associated with the Magellanic System. Compared to the total neutral gas mass of ${M}_{\mathrm{neutral}}=\left(1.9\mbox{--}2.1\right)\times {10}^{9}\,{M}_{\odot }$ from these studies, the gas in these galaxies and their extended environment is significantly ionized. Hopkins et al. (2012) modeled star-forming galaxies of varying mass ranges and investigated the warm $\left(2000\,{\rm{K}}\lt {\text{}}T\lt 4\times {10}^{5}\,{\rm{K}}\right)$ gas to cold (T < 2000 K) gas ratios of the star-forming disks. Their models found a warm-to-cold gas ratio of 0.5–0.7 for MW and SMC-like galaxies. Another study by Dobbs et al. (2011) investigated the ISM in star-forming spiral galaxies. They split their gas fractions into cold gas $\left(T\lt 150\,{\rm{K}}\right)$, intermediate gas $\left(150\lt T\lt 5000\,{\rm{K}}\right)$, and warm gas $\left(T\gt 5000\,{\rm{K}}\right)$ and find the total fraction of gas in each state is roughly equal. While neither of these studies is directly analogous to our observations, we find comparable warm-to-cold gas ratios.

Many models of the MSys fall short of replicating both the MS and LA in extent and morphology, thus investigations into the extended material are important for informing future models. A recent paper by Lucchini et al. (2020) investigates the impact of adding an envelope of warm–hot coronal gas surrounding the two galaxies, called the Magellanic Corona. The Magellanic Corona is defined as a halo gas at a transition temperature of ∼10⁵ K surrounding the LMC with a mass of 3 × 10⁹ M_☉ extending over the virial radius 100 kpc.

The warm ionized gas studied here is at 10⁴ K (Hoyle & Ellis 1963; Haffner et al. 2003) and is not directly comparable to the warm–hot corona defined in the Lucchini et al. (2020) models. At 10⁴ K, the relation between the measured Hα intensity and the emission measure scales according to Equation (8). For temperatures above 2.6 × 10⁴ K, the Hα intensity is highly sensitive to temperature. From Draine (2011),

$$\begin{eqnarray}\begin{array}{rcl}\mathrm{EM} & = & 2.77\,{\mathrm{cm}}^{-6}\,\mathrm{pc}\left(\displaystyle \frac{{I}_{{\rm{H}}\alpha }}{R}\right){\left({\epsilon }_{{\rm{b}}}\right)}^{-1}\\ & & \times \,{T}_{4}^{\left(0.942+0.031\,\mathrm{ln}{T}_{4}\right)},\end{array}\end{eqnarray} \tag{ 12 }$$

where T₄ = T_e /10⁴ K and _b is the beam dilution factor. Assuming _b = 1 and T_e = 10⁵ K, 25 mR would result in EM = 0.71 cm⁻⁶ pc. While the emission measure of the Magellanic corona has not been measured, if it is similar to the MW with 0.005 ≤ EM ≤ 0.0005 cm⁻⁶ pc (Henley et al. 2010), it is far below the threshold for WHAM to isolate with Hα observations. While the observations in this paper do not trace the corona, we do find an extended, complex, multiphase halo surrounds the LMC.

Lehner & Howk (2007) suggest the warm–hot corona gas surrounding the MCs acts as a shield around stripped gas from the Clouds, preventing the MW hot corona from fully ionizing the structures and allowing them to persist longer. Our observations of the near LMC portion of LA 1.1 shows the region to be 15%–47% ionized (Table 3). This may suggest an envelope of ionized gas surrounding LA 1.1. The presence of neutral gas in this region does not directly indicate a Magellanic Corona, but additional studies of the ionizing radiation in this region may indicate if the Corona is contributing.

The total combined ionized gas mass of the central galaxies from SHB19 and this study is ${M}_{\mathrm{LMC}+\mathrm{SMC},\mathrm{ionized}}\,=\left(1.2\mbox{--}2.4\right)\times {10}^{9}\,{M}_{\odot }$. This value is comparable to the mass needed in the Magellanic Corona for the LA to survive (Lucchini et al. 2020). While the models only considered the Magellanic Corona contributing to the missing mass budget of the MC system and the ionized gas needed to shield the LA, combining both the mass from the WIM as well as the Magellanic Corona may further guide simulations tracing the history and evolution of the Clouds.