Exploring Hydrodynamic Instabilities along the Infalling High-Velocity Cloud Complex A
Kathleen A. Barger, David L. Nidever, Cannan Huey-You, Nicolas Lehner, Katherine Rueff, Paris Freeman, Amber Birdwell, Bart P. Wakker, Joss Bland-Hawthorn, Robert Benjamin, Drew A. Ciampa
DDraft version January 29, 2021
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Exploring Hydrodynamic Instabilities along the Infalling High-Velocity Cloud Complex A
Kathleen A. Barger, David L. Nidever,
2, 3
Cannan Huey-You,
1, 4
Nicolas Lehner, Katherine Rueff, Paris Freeman,
1, 6
Amber Birdwell,
1, 7
Bart P. Wakker, Joss Bland-Hawthorn, Robert Benjamin, andDrew A. Ciampa Department of Physics & Astronomy, Texas Christian University, Fort Worth, TX 76129, USA Department of Physics, Montana State University, P.O. Box 173840, Bozeman, MT 59717-3840 National Optical Astronomy Observatory, 950 North Cherry Ave, Tucson, AZ 85719 Accommodated Learning Academy, Grapevine, TX 76051, USA Department of Physics, University of Notre Dame, Notre Dame, IN 46556, USA Founders Classical Academy of Lewisville, Lewisville, TX 75067, USA Aledo High School, Aledo, TX 76008, USA Supported by NASA/NSF, affiliated with Department of Astronomy, University of Wisconsin-Madison, Madison, WI 53706, USA Sydney Institute for Astronomy, School of Physics A28, University of Sydney, NSW 2006 University of Whitewater, Whitewater, WI 53190, USA
ABSTRACTComplex A is a high-velocity cloud that is traversing through the Galactic halo toward the MilkyWay’s disk. We combine both new and archival Green Bank Telescope observations to construct aspectroscopically resolved H i . (cid:46) log( N H i , σ / cm − ) (cid:46) . − line and ∆ θ = 9 . (cid:48) (cid:46) ∆ d θ (cid:46)
30 pc spatial resolution. Wefind that that Complex A is has a Galactic standard of rest frame velocity gradient of ∆v
GSR / ∆L =25 km s − / kpc along its length, that it is decelerating at a rate of (cid:104) a (cid:105) GSR = 55 km / yr , and that itwill reach the Galactic plane in ∆ t (cid:46)
70 Myrs if it can survive the journey. We have identify numeroussignatures of gas disruption. The elongated and multi-core structure of Complex A indicates thateither thermodynamic instabilities or shock-cascade processes have fragmented this stream. We findRayleigh-Taylor fingers on the low-latitude edge of this HVC; many have been pushed backward byram-pressure stripping. On the high-latitude side of the complex, Kelvin-Helmholtz instabilities havegenerated two large wings that extend tangentially off Complex A. The tips of these wings curveslightly forward in the direction of motion and have an elevated H i column density, indicating thatthese wings are forming Rayleigh-Taylor globules at their tips and that this gas is becoming entangledwith unseen vortices in the surrounding coronal gas. These observations provide new insights on thesurvivability of low-metallicity gas streams that are accreting onto L (cid:63) galaxies. Keywords:
Galaxy: evolution - Galaxy: halo - ISM: individual (Complex A) - Physical Data andProcesses: hydrodynamics - Physical Data and Processes: Instabilities INTRODUCTIONThe star formation in galaxies is dependent on theirability to accrete gas onto their disks. Although boththe Milky Way and Andromeda are surrounded by gas(e.g., Wakker & van Woerden 1997; Wakker et al. 2003;Braun & Thilker 2004; Lehner & Howk 2011; Lehneret al. 2015), their star-formation rates appear to be in adecline (see Bland-Hawthorn & Gerhard 2016 for a re-view). These galaxies may even be transitioning into the“Green Valley” (Mutch et al. 2011; Davidge et al. 2012;Bland-Hawthorn & Gerhard 2016), which is the region between blue star-forming and red quiescent galaxies ona color-magnitude diagram.The halo clouds that surround the Milky Way aretypically put into two different categories that arebased on their local standard of rest (LSR) velocities.Intermediate-velocity clouds (IVCs) are a slower popu-lation (30 (cid:46) | v LSR | (cid:46)
90 km s − ) that tend to lie nearthe Galactic disk. The high-velocity cloud (HVC) popu-lation ( | v LSR | (cid:38)
90 km s − ) has multiple origins, includ-ing galactic-feedback processes, halo-gas condensations,nearby low-mass galaxies, and intergalactic medium fil-aments; therefore, many HVCs provide replenishment a r X i v : . [ a s t r o - ph . GA ] J a n Barger et al. the star-formation reservoir of our galaxy (Putman et al.2012 and Richter 2017 for review).As HVCs clouds travel through galaxy halos, theyheated and ionized by photons that are escaping fromthe galaxies (e.g., Milky Way: Bland-Hawthorn & Mal-oney 1999, 2001; Fox et al. 2005; Magellanic Clouds:Barger et al. 2013). Additionally, the hot coronal gasthat surrounds them acts as a headwind that compressestheir leading material and strips its outer layers througha process known as ram-pressure stripping (e.g., Put-man et al. 2011; For et al. 2014). When the surround-ing gas rubs against the HVC’s surface, it promotesKelvin-Helmholtz instabilities—a type of shear-drivendisturbance—which can cause small cloudlets to frac-ture off the complex’s main body (see Stone & Gar-diner 2007; Bland-Hawthorn et al. 2007; Heitsch & Put-man 2009). Rayleigh-Taylor instabilities, which arebuoyancy-driven disturbances, further disrupt the com-plex because it is resting on top of less dense halo gaswhile situated in a galaxy’s gravitational field. Com-bined, these processes can cause the skin of the cloudto become warmer, ionized, and more diffuse than itscore. Internal temperature and density variations be-tween these two gas phases can generate thermal in-stabilities, which can fragment the cloud (see Murray& Lin 2004). Fragmentation can also occur if strippedleading gas, due to ram-pressure stripping, collides withdown stream material (see Bland-Hawthorn et al. 2007;Tepper-Garc´ıa et al. 2015). As the surface area of theHVC increases, it becomes more exposed to its environ-ment, which will cause it to evaporate more rapidly intothe surrounding coronal gas (e.g., Konz et al. 2002).Complex A is plummeting towards the Galactic diskand could supply our galaxy with up to M total (cid:38) × M (cid:12) (neutral: Kunth et al. 1994; van Woerden &Wakker 2004; ionized: Barger et al. 2012) of new mate-rial ( Z = 0 . Z (cid:12) : Kunth et al. 1994; Schwarz et al. 1995;van Woerden et al. 1999; Wakker 2001; Barger et al.2012). Its chemical composition indicates that it eitheroriginated from a low-mass galaxy or the intergalacticmedium. However, as no complementary stellar streamhas been found (Belokurov et al. 2010; Newberg et al.2010), Complex A was not likely stripped from a satellitegalaxy. This complex has an elongated morphology withmultiple dense cores—dubbed A0–AVI and B—along its∆ L ≈ . θ ≈ ◦ ; Barger et al. 2012).Because this infalling cloud spans 3 (cid:46) z (cid:46) Figure 1.
An H i map of Complex A with the regions that wereobserved through different programs highlighted (see Table 1) andwith the eight high H i column density cores A0–AVI and core Blabeled. The two purple shaded regions mark the locations of ournew observations that primarily span the trailing portion of thiscomplex (PI Barger: GBT13B-068). The region highlighted in redspans the leading portion of this gaseous stream (PI Verschuur:ID GBT1010A-003). The yellow (PI Chynoweth: GBT09A-046)and blue (PI Martin: GBT107A-003) regions indicate observationscover the central region of this HVC. We additionally circle theemission in our surveyed region that is associated with the M81galaxy at ( l, b ) = (142 . ◦ , . ◦ In this study, we investigate how HVC Complex Ais affected by its environment with new and archivalGreen Bank Telescope (GBT) H i i mor-phology and kinematic structure along the length of thecomplex in Section 6 and discuss morphological featuresthat are indicative of hydrodynamic instabilities occur-ring within different regions of Complex A. Finally, wesummarize our main conclusions in Section 8. OBSERVATIONSOur H i − ≤ v LSR ≤−
90 km s − velocity range and are spatially resolved at∆ θ = 9 . (cid:48) (cid:46) ∆ d θ (cid:46)
30 pc at the distance of Com-plex A (6 . (cid:46) d (cid:12) (cid:46) . Throughout this study, we use the kinematic definition ofthe LSR, where the solar motion moves at 20 km s − toward(R . A ., DEC . ) J2000 = (18 h m . s , ◦ (cid:48) . (cid:48)(cid:48) omplex A Table 1.
GBT DataProp ID PI Name l, b
Center Angular Size ∆ θ ∆v LSR log( N H i , σ / cm − ) a (degrees) (degrees) (arcmins) (km s − )GBT09A-046 b Chynoweth M81 143 . ◦ , . ◦ ◦ × ◦ . (cid:48)
75 5.2 17 . . ◦ , . ◦ ◦ × ◦ . (cid:48) . . ◦ , . ◦ ◦ × ◦ . (cid:48) . . ◦ , . ◦ ◦ × ◦ . (cid:48) . c Martin SPC 135 . ◦ , . ◦ ◦ × ◦ . (cid:48) . . ◦ , . ◦ ◦ × ◦ . (cid:48) . . ◦ , . ◦ ◦ × ◦ . (cid:48) . . ◦ , . ◦ ◦ × ◦ . (cid:48) . . ◦ , . ◦ ◦ × ◦ . (cid:48) . . ◦ , . ◦ ◦ × ◦ . (cid:48) . a The 1-sigma sensitivity for a FWHM = 20 km s − line. A sensitivity map of all the mapped observations explored in thisstudy is presented in Figure 2. b This program scanned in ICRS coordinates while all others used Galactic coordinates. c This region also includes data from GBT08B-083 (PI: Goncalves), GBT08B-038 (PI: Goncalves) and GBT10A-078 (PI:Lockman). σ Sensitivity (10 atoms cm −1 ) G a l a c t i c La t i t ude [ D eg r ee s ] Figure 2.
Sensitivity map (1 σ ) of the Complex A GBT obser-vations in units of 10 atoms cm − over 20 km s − . upper velocity limit of this survey (v LSR ≤ −
90 km s − )is to avoid contributions from the Milky Way’s disk andto reduce the contribution from the neighboring HVCComplex C (see Figure 1). For the core A0 region, whichlies nearest to the Galactic disk, we truncated this ve- locity limit to v LSR ≥ −
130 km s − in an effort to avoidMilky Way contamination. It is important to note thatwe did not search for H i emission associated with Com-plex A below a Galactic latitude of b < . ◦ . . ≤ ν L-Band ≤ .
73 GHz) obser-vations on the H i ν = 1420 . ν = 12 . channel = 0 . − channel width.The archival GBT H i ∼
215 deg region on thesky (PI Chynoweth: GBT09A-046), (2) core A0 over a ∼
60 deg region (PI Verschuur: GBT10A-003), and (3)cores AI, AIII, AIV, and AV over a ∼
260 deg region(PI Martin: GBT07A-104, GBT08A-083, and GBT10A-078). Table 1 summarizes the angular extent, angu-lar and spectral resolution, and the sensitivity of thesedatasets. The observations from the Verschuur andPlanck programs were taken with a 4 second exposuretime and the ones from the Chynoweth program weretaken with 5 second exposures. Together, these archival Barger et al.
Figure 3.
Example H i omplex A Figure 4.
Vertically stacked residuals image for each of the ∼ . × sightlines explored in this study as a function of ve-locity, where the residuals are defined as the difference betweenthe H i observations and the Gaussian decomposition fit. Weexclude a small region of our survey shown in Figure 6 that iscontaminated by emission from the M81 galaxy system containedbetween 141 ◦ (cid:46) l (cid:46) ◦ and 40 . ◦ (cid:46) b (cid:46) ◦ . In the bottomright hand corner, the column density of the residuals increaseswhere core A0 overlaps with the Milky Way; to avoid confusionwith our Galactic disk, we only report on the properties of the H i emission below v LSR ≤ −
130 km s − in this region. datasets stretch over a 470 square degrees region acrossthe sky along ∼ DATA REDUCTIONFor this project, we used a reduction and calibrationprocedure that is very similar to the one used by Nideveret al. (2010) with only a few small modifications. Wecalibrated antenna temperature ( T A ) of our frequency-switched H i us-ing the relationship: T A , calib ( ν ) = T refsys ( ν ) × F sig ( ν ) − F ref ( ν ) F ref ( ν ) (1)Here, the T refsys ( ν ) is the system temperature that we de-termined from the reference spectrum, F ref ( ν ) is the ref-erence flux of a calibration target, and F sig ( ν ) is theflux of the on-target signal. We assumed that bright-ness temperature ( T B ) is roughly equal to the antennatemperature (i.e., T B ≈ T A ). Our calibration objectsincluded standard S7 and S8 reference targets located We performed this calibration with the GETFS program inthe GBTIDL (http://gbtidl.nrao.edu/) software. at ( l, b ) = (207 . ◦ , − . ◦
00) and (207 . ◦ , − . ◦ . ◦ , . ◦
04) during each obser-vational run, enabling us to measure the H i emissionalong this sightline very accurately and to use it as asubstitute flux calibration target whenever both S7 andS8 were below the horizon.Once calibrated, we binned the spectra to a ∆v bin =0 . − velocity spacing to decrease small scale fluc-tuations in the signal and further smoothed the spectrawith a Gaussian kernel and then removed the continuumlevel. We fit the baseline for each integration and polar-ization separately after masking out the Galactic emis-sion between − (cid:46) v LSR (cid:46) +100 km s − and emissionlines above 2-sigma. For the XX polarization, we fit a5th-order polynomial to the continuum. We used thesame procedure for the YY polarization with the addi-tion of a sinusoidal component in the fit to remove astanding wave that has a period of ν ≈ . T B (v) = 0 mK km s − by subtract-ing the median emission–free spectral height of all spec-tra in an observing session from each polarization. Fi-nally, we averaged the two polarizations together to pro-duce reduced spectra. The resultant spectra has a typi-cal root-mean-square (RMS) noise that is T B ≈
75 mKper 0 . − channel. This corresponds to a spectralnoise sensitivity of log( N H i , σ / cm − ) = 17 . − (see Fig-ures 2 and 4), using N H i , σ cm − = 1 . × (cid:18) T B , σ K (cid:19) (cid:114) FWHMkm s − ∆v channel km s − (2)to convert between T B and N H i noise sensitivity un-der the assumption that the spectral noise is Gaussianin nature (Wolfe 2014). Therefore, our 3-sigma detec-tion limit is log( N H i , σ / cm − ) = 18 . − width. We determined the N H i ofour lines from the T B under the assumption that emit-ting gas is optically thin to self absorption:N H i cm − = 1 . × (cid:90) (cid:18) T B (v)K (cid:19) (cid:18) dvkm s − (cid:19) (3)For the archival datasets, we use the already reducedand calibrated observations that were shared with us bythose program leaders. The reduction and calibrationprocedures outlined above for our new GBT observa-tions from the GBT13B-068 (PI Barger) program werethe same procedures used to reduce the GBT10A-003(PI Verschuur) dataset, which they outline in their studythat explored the physical conditions of Complex A’s Barger et al.
Figure 5.
A 3 dimensional movie of Complex A’s H i gas distribution that rotates through position-position-velocity maps. The H i column densities and LSR line center are solutions to our Gaussian decomposition fits (see Section 4). All points in the 3D map have atransparency of 50%. The movie begins with an H i position-position map that is rotated by 90 ◦ about the Galactic Longitude axis intoa Galactic longitude position-velocity map. The 3D map in this movie is then rotated by another 360 ◦ about the Galactic Latitude andLongitude plane, passing through Galactic longitude position-velocity and Galactic latitude position-velocity maps twice along the way.Finally, the 3D map is rotated back into its original position by rotating 90 ◦ about the Galactic Longitude axis. The real-time duration ofthe movie is 1 minute and 21 seconds. core A0 (Verschuur 2013). The calibration and contin-uum level removal techniques in for the GBT07A-104(PI Martin) observations are described in Boothroydet al. (2011) and Martin et al. (2015).The observation and reduction procedures used on theGBT09A-046 (PI Chynoweth) dataset differs substan-tially from the other datasets. These observations weretaken in position switching mode instead of frequency-switching mode that was used for all other datasets inour survey. For each observing session, Chynoweth et al.(2011) used a reference spectrum positioned at the edgeof their observing grid as a reference signal as their fluxcalibrator. Through this calibration scheme, most of theMilky Way zero-velocity is removed, but the Complex Afeatures remain essentially intact. The spectra was thenbinned to ∆v bin = 5 . − . We combined and resampled all of the new andarchival datasets—except the GBT09A-046 (PI Chynoweth)dataset—on a large, uniform grid in Galactic coordi-nates with ∆ θ = 3 . (cid:48) . − velocity bins. Because the velocity sampling and an-gular resolution of the GBT09A-046 (PI Chynoweth)dataset differ the most from the other datasets, we onlyused those observations when other data were not avail-able. In Figure 2 we shows the 1 σ sensitivity map ofour all observations. GAUSSIAN DECOMPOSITIONSWe determined the component structure of the H i omplex A Figure 6.
Gas distribution and motions of Complex A over a − ≤ v LSR ≤ −
90 km s − velocity range for the entire complex exceptin the core A0 region, which is over a − ≤ v LSR ≤ −
130 km s − velocity range to avoid emission associated with the Galactic disk.Top Left: H i column density; Top Right: FWHM of the emission line; Bottom Left: Center LSR velocity of the emission line; BottomRight: Center GSR velocity of the emission line. All values in these three maps depict solutions from Gaussian decompositions along eachsightline, where the shown color represents the properties of the component with the highest H i column density. determined the number Gaussians to model the emissionby minimizing the reduced chi-squared ( (cid:101) χ ) with theMPFIT IDL routine (Markwardt 2009). We selectedthe fit that used the smallest number of Gaussians toachieve a reduced-chi squared that was within 0.25 ofthe best fit to avoid arbitrarily adding more and moreGaussian profiles that are not of physical significance tobetter match the H i emission.Our initial guesses for the Gaussian fit parameters de-termined by iteratively searching for peaks in each spec-trum. We identified peaks as locations with the highest N H i (v) emission above the 1 standard deviation noiselevel of a smoothed spectrum. We then masked out allregions with v peak ±
25 km s − and repeated this searchfor emission-lines in the unmasked region. We then fitthe spectrum using for lines at velocity positions with∆v = 10 km s − for the initial guess of the line width.We checked the quality of each fit by searching the resid-uals for any remaining unfitted emission lines, where the The IDL MPFIT routines are available athttp://purl.com/net/mpfit. residuals are the spectral signatures remaining after thefit is subtracted from the H i emission spectrum:residuals(v) = N H i , obs (v) − N H i , fit (v) (4)We only report results for fitted line profiles that havea signal-to-noise ratio (S / N) greater than 2 and thatare above the 3-sigma detection limit of our survey atlog( N H i , σ / cm − ) ≥ .
7. We defined the area of thesignal in the S / N to as the area of the fitted Gaussian-line profile and the area of the noise to be equivalentto the area of a rectangle that has a height equal tothe standard deviation of the continuum and a widthequal to the FWHM of the fitted Gaussian-line profile.This step was done to remove unrealistic fits that char-acterized spikes in the noise and to ensure that the faintemission that is associated with the diffuse outer enve-lope of Complex A was kept.In Figure 3, we illustrate representative Gaussian de-compositions of three sightlines along cores AIII, AVI,and B and the corresponding residuals of those fits.For each of these sightlines, there is a brighter com-ponent that is associated with the H i core. Addi-tionally, we often find fainter and wider components, Barger et al.
Column Density (10 atoms/cm ) − − − − − C o m p l e x A La t i t ude [ D eg r ee s ] V LSR < −
90 km s − Milky Way
Direction of motionA0 AI AII AIII AIV BAVIW1W2
M81
RT RT RT RT
RPS
RPS RT RPSRT
RPSRPS
KH KH KHRT?RT?KH Turbulent wakeof W2RPS
Figure 7.
The total integrated H i column density distributionof Complex A in the CA coordinate system, which places the mainbody of this HVC stream at b CA = 0 ◦ and the center of core A0 at( l CA , b CA ) = (0 ◦ , ◦ ) or ( l, b ) = (133 . ◦ , +25 . ◦ N H i was determined using the Gaussian decompositions procedure de-scribed in Section 4. All have marked all major cores (A0–AVIand B) are marked and high-latitude wing 1 (W1) and wing 2(W2) that lies off the core AVI and AIV regions, respectively. Welabel prominent Rayleigh-Taylor (RT) globules, Kelvin-Helmholtzstructures, and ram-pressure stripping (RPS) features. The redbox ( l, b ) = (142 . ◦ , . ◦
9) encompasses emission from the M81galaxy.
Average Column Density along b CA (10 atoms cm − ) − − − − − − V L S R ( k m s − ) Milky Way Milky WayA0 AI AII AIII AIV AVI BW1W2
Figure 8.
Position-velocity map showing the column density av-eraged in b CA of Complex A and the Milky Way H i l CA = 0 ◦ in the CA coordinate System. The N H i andv LSR center positions of the emission for each component alongall sightlines was determined using the Gaussian decompositionprocedure described in Section 4. which trace the diffuse gas in the outer envelope of this
Figure 9.
The median line widths and hydrogen temperaturesfor each of the major H i core regions (see Figures 11–18) andthe two high-latitude wing regions (see Figures 19–20). The tem-peratures were estimated from median width of the lines underthe assumption that they are only affected by thermal-line broad-ening affects. The uncertainties in the line width represent theaverage deviation from the median. The blue dashed line marksour best fit for the FWHM as a function of Complex A Longitude,which is given by the following linear relationship: FWHM =( − . ± .
07) km s − degree − l CA + (20 . ± .
6) km s − HVC. We include a movie that rotates from througha 3 dimensional position-position-velocity map of Com-plex A that shows all of the LSR line center and N H i solutions to our Gaussian decomposition fits in Fig-ure 5. We additionally provide a vertically stacked spec-tral image of the residuals of all ∼ . × sight-lines explored in this study in Figure 4, which are typi-cally log (cid:0) residuals(v) / [ cm − km s − ] (cid:1) < . COMPLEX A COORDINATE SYSTEMBecause Complex A has a long, filamentary struc-ture it is useful to have a coordinate system wherethe equator lies along the great circle of this extendedH i structure. We define a “Complex A” coordinate sys-tem with a pole at ( l, b )=(202 ◦ , − ◦ ) and the originof the longitude axis (the ascending node) defined suchthat the center of the A0 core at ( l, b ) = (133 . ◦ , +25 . ◦ l CA , b CA ) = (0 ◦ , ◦ ). As in the Mag-ellanic (Wakker 2001) and Magellanic Stream coordi-nate systems (Nidever et al. 2008), l CA decreases alongComplex A towards higher Galactic latitudes. Figure 7shows the column density of Complex A and Figure 8shows the position–velocity diagram in this new coordi-nate system. omplex A Figure 10.
Histogram distributions of the FWHM for each of the major H i core regions (see Figures 11–18) and the two high-latitudewing regions (see Figures 19–20). These line widths were through determined by decomposing the H i spectra into Gaussian components(see Section 4). We additionally mark the gas temperature that these widths would correspond to for a pure thermal broadening scenario,where the 8,000 (cid:46) T (cid:46) α emission peaks. NEUTRAL GAS MORPHOLOGY ANDKINEMATICS6.1.
Global Properties
Complex A is an elongated stream with multiple denseH i cores along its length (see Figures 6 and 7). Thesecores tend to be more compressed on the lower Galacticlatitude and longitude side (or higher l CA side) and morediffuse and elongated on the opposite side, indicatingthat core A0 represents the leading end of this streamand core AVI represents the trailing end (see more dis-cussion below). This stream is wider at its trailing end(see Figures 6 and 7). Because the leading edge of Com-plex A is much closer ( d (cid:12) ≈ d (cid:12) ≈
10 kpc; see Barger et al. 2012), this meansthat the wider angular extent of cores AVI and B corre-sponds to a much larger physical width at ∆ θ core A0 ≈ .
10 kpc / degree vs ∆ θ core AVI ≈ .
17 kpc / degree. How-ever, the relatively inline A0–AVI cores suggests thatthey are part of the main body of Complex A and thatcore B represents material that fractured off this gasstream.There is a relatively coherent Galactic standard ofrest (GSR) velocity gradient along the length of Com-plex A, where its leading gas traveling slower relativeto the Milky Way than its trailing gas (see lower right-hand panel in Figure 6); this indicates that that Com-plex A is deccelorating . The GSR velocity gradient alongits ∆ θ ≈ ◦ (or ∆ L = 5 . GSR / ∆ θ = 4 . − / degree (or∆v GSR / ∆L = 25 km s − / kpc). Assuming an averagevelocity of (cid:104) v GSR (cid:105) ≈ −
70 km s − , it has taken Com-plex A ∆ t = 80 Myrs to travel the length of its body,corresponding to decceleration of (cid:104) a (cid:105) GSR = 55 km / yr .0 Barger et al.
Figure 11.
Gas distribution of the core A0 region over the − ≤ v LSR ≤ −
130 km s − velocity range with the Rayleigh-Taylor (RT)globules, Kelvin-Helmholtz structures, and ram-pressure stripping (RPS) features labeled in the Top-Left panel. The colors in the toprow of panels represent the H i column density and the colors on the bottom panels represent the FWHM line width. All of the valueswere determined through Gaussian decompositions (see Section 4). Left Columns: Galactic Latitude and Longitude position-position map.Middle Columns: Galactic longitude position-velocity map. Right Columns: Galactic latitude position-velocity map. The panels in theLeft Column only display the N H i and FWHM for the component with the highest H i column density along each sightline, though allcomponents are included in the Middle and Right Column position-velocity maps. Only the emission contained within the spatial regionmarked black boundary in the Left panel is the included in the Middle and Right position-velocity maps. The faint gray-scale vertical barsthe in the background of the position-velocity map represent the kinematic extent of the emission and span from v extent = v center ± ∆v width . At a constant acceleration, this complex will cross theGalactic plane at b = 0 ◦ in ∆ t (cid:46)
70 Myrs; this is anupper limit as this time should decrease due to an in-creasing gravitational pull as the Complex A approachesthe Milky Way’s.We also find that there is a graduate increase in thewidth of the H i line toward the trailing end of Com-plex A. In Figure 9, we have plotted the median linewidths for each of the major core regions and the twohigh-latitude wings as a function of Complex A longi-tude. In the core A0 region, the median FWHM linewidth is roughly FWHM ≈
19 km s − at l CA ≈ ◦ ,but grows to FWHM ≈
25 km s − at l CA ≈ − ◦ forcores A6 and B. We have included the histogram distri-butions of the line widths for all fitted components ineach of these regions in Figure 10. In general, the his-togram distributions for each core region are relativelywell behaved with an easy to identify peak in the num-ber of components at a particular line width, except forthe core A0. For this leading core, the line widths peak between 10 (cid:46) FWHM (cid:46)
23 km s − and include a muchlarger distribution of narrow lines than any other coreregion. These narrow lines suggest that this core is cool-ing rapidly, presumably because this low metallicity coreis mixing with the higher metallicity gas near the MilkyWay’s disk.Assuming that the H i emission lines are only broad-ened by thermal broadening, the increasing median linewidths along the length of Complex A would correspondto a rise in the hydrogen gas temperature by roughly4,400 K from T H i , median = 8,700 K along the leadingedge of the complex to 13,100 K along its trailing edge(see Figure 9). Overall, this is relatively inline withthe typical gas temperature that Barger et al. (2012)found for the warm ionized phase of this complex at T H α = 12 ,
600 K in the direction of the H i cores, wheretheir WHAM H α observations were resolved at ∆ θ = 1 ◦ and have an angular area that is larger than the GBTH i observations by a factor of A θ, WHAM ≈ A θ, GBT .However, the elevated line widths on the trailing end of omplex A Figure 12.
Same as Figure 11, but for the core AI region over the − ≤ v LSR ≤ −
90 km s − velocity range. Complex A could also signify that this gas is experienc-ing an increase in non-thermal motions. If that is thecase, then the emission lines associated with the trailingend of Complex A would be additionally non-thermallybroadened by FWHM non − thermal = 14 . − , assum-ing a thermal broadening of FWHM thermal ≈
20 km s − .This significant contribution to the non-thermal broad-ening of the line width would indicate that the trailinggas is being more disrupted.The higher H i column density cores in this streamare connected by lower column density gas. Many of theH i cores are compressed in the direction of their mo-tion, including cores A0, AI, AII, and AIII. These coresalso tend to be moving towards the Galaxy faster (i.e.,larger negative LSR velocities) than the lower columndensity gas that surrounds and connects them. Thisglobal trend is especially apparent in movie found inFigure 5, which rotates Complex A through 3 dimen-sional position-position-velocity space, and in Figure 8.We additional provide two sets of zoomed in position-position and position-velocity maps that are scaled bythe H i column density and FWHM line width of eachcore region in Figures 11–20. This global morphologyis characteristic of ram-pressure stripping in which thesurrounding coronal gas and incident Galactic photonsact as a headwind that compresses the leading gas andstrips the outer layers of this stream to form a laggingdiffuse tail that travels in the anti-direction of motion. The fragmented morphology of Complex A could bea result of thermal cooling instabilities or a slow “shockcascade.” Cooling instabilities often arise as a result ofdensity inhomogeneities in which the high density gascools more efficiently. As the complex descends towardthe disk, it will sweeps up coronal gas, compressing theleading gas (Kereˇs & Hernquist 2009). Further, gas thatthe complex sweeps up gas near the Milky Way’s diskwill have a higher metallicity ( Z CA = 0 . Z (cid:12) : Kunthet al. 1994; Schwarz et al. 1995; van Woerden et al. 1999;Wakker 2001; Barger et al. 2012 and will promote cool-ing. Fragmentation is expected to occur once the streamsweeps up roughly its own mass in ambient material(Murray & Lin 2004), indicating that Complex A hasalready accreted a substantial material during its jour-ney. Unfortunately, the sparseness of metallicity mea-surements along the length of the complex means thatthe level of metal mixing and accretion cannot currentlybe constrained. In a “shock cascade” scenario, leadingmaterial that is stripped via ram-pressure will be slowedby non-conservative forces and can then collide withdown stream gas (Bland-Hawthorn et al. 2007; Tepper-Garc´ıa et al. 2015). This shock cascade can disrupt andand fragment the down stream gas. The rapidly varyingline widths with position and column density, between10 (cid:46) FWHM (cid:46)
40 km s − (see Figures 11–20) are astrongly indicator that the low and high column densityH i gas is either not in thermal-dynamical equilibrium2 Barger et al.
Figure 13.
Same as Figure 12, but for the core AII region. or that the low column density gas is experiencing moresevere non-thermal motions.Interestingly, the gas in the core AII region has a muchlower H i column density than the gas in the adjacentcores that connect to it. Further, this relatively wispycore region is moving much slower toward the Galacticdisk than cores AI and AIII at ∆v LSR ≈
80 km s − off-set from core AI and ∆v LSR ≈
30 km s − offset fromcore AIII (see Figure 8), indicating that it is muchmore influenced by coronal-gas interactions. However,although core AII is morphologically much more dis-rupted, its higher column density gas still has a narrowline profile (10 (cid:46) FWHM (cid:46)
20 km s − ; Figure 13). Thissuggests that the gas in the core AII region is still able toremain relatively cool and compact. Using mapped H α observations, Barger et al. (2012) found that roughlyhalf of Complex A is ionized. This warmer and lowerdensity phase acts as skin that shields the H i coresfrom direct interactions with the surrounding coronalgas. The gas in core AII could also be “drafting” theleading gas in core AI, such that it is not experiencing adirect headwind, though we do not know the locations ofthese cores in 6-dimensional position and velocity spaceand therefore cannot tell how well aligned core AII isbehind core AI.Numerous H i structures protrude or are fractured offComplex A’s main body, which is an indication that itsgas is subject to hydrodynamic instabilities. As these offset structures still have a relatively cohesive structurein H i , they were likely recently stripped off the complex.This displaced gas is now more exposed to the incidentionizing radiation and the surrounding coronal gas asthis material now has a larger surface area and is nolonger “drafting” behind the cloud. This increased ex-posure will cause them to be heated and ionized quicker,which will lead to them rapidly evaporating (Konz et al.2002). We identify these offset structures and discusshow Rayleigh-Taylor and Kelvin-Helmholtz instabilitiesare working with ram-pressure stripping producing thesestructures in the following subsections.6.2. Rayleigh-Taylor Instability Structures
Complex A is surrounded by hot coronal gas, whichmeans that it is essentially resting on top of a lowerdensity medium while being influenced by the MilkyWay’s gravitational field. This is an unstable config-uration that can drive buoyancy related disturbancesknown as Rayleigh-Taylor instabilities. If these instabil-ities are strong enough, they can generate globules andspikes (“finger” like structures) that drip through thewarmer coronal medium that lies below Complex A to-ward the center of the Milky Way’s gravitational field.These globules will therefore form on the lower Galacticlatitude edge of this complex (see Figure 7). Further, themorphology will include a compressed edge that forms omplex A Figure 14.
Same as Figure 12, but for the core AIII region. when globules push through the lower density coronalgas below it.There is an H i arch that hangs off core A0 at ( l, b ) ≈ (132 ◦ , ◦ ) by a thin filament (see Figure 11). Thehigh-latitude portion of this arch is more compressed,indicating that it was the material that initially pushedthrough the coronal gas when the globular began its de-parture from core A0. This gas arches in Complex A’sdirection of motion, which is unusual as ram-pressurestripping should have pushed this material in the otherdirection. Instead, as core A0 is only z ≈ . i structure that looks like a skewed “loop” at ( l, b ) ≈ (141 ◦ , . ◦
5) (see Figure 12). This material connects tocore AI at (140 ◦ , ◦ ) and extends in the anti-directionof Complex A’s motion and then curves back up to-ward core AI. This loop appears to represent a Rayleigh-Taylor spike that was elongated and pushed backwarddue to ram-pressure stripping by the surrounding coro-nal gas. A Rayleigh-Taylor spike also lies on the lowerlatitude side of core AIII at (1 , b ) ≈ (149 ◦ , ◦ ) (seeFigure 14). However, this relatively shorter and wider structure projects downward in Galactic latitude anddoes not curve backward, indicating that this gas onlyrecently “dripped” off core AIII.There is a mini stream that branches off core AIV’slower latitude edge and points in the anti-direction ofComplex A’s motion (see Figure 15). This mini streamrepresents a complex Rayleigh-Taylor instability struc-ture that is strongly influenced by ram-pressure strip-ping. The gas that is positioned directly under core AIVhas only recently “dripped” off core A4. Below the be-ginning of this mini stream, there is a mini H i knotat (1 , b ) ≈ (154 ◦ , ◦ ) that appears to be the start ofa new Rayleigh-Taylor spike. At (151 ◦ , ◦ ), there isa small H i knot that branches off in the direction ofComplex A’s motion, which would occur if this globularis flowing into a low pressure pocket behind core AIII.Along this stream, there are two H i knots at (156 ◦ , ◦ )and (159 ◦ , ◦ ) that indicate that there are smallerRayleigh-Taylor fingers forming off other fingers.The gas associated with core B (Figure 18) doesnot align with the main body of Complex A (i.e.,cores A0–AVI; see Figure 6). The gas that connectscore B to core AVI has a relatively lower column density(log (cid:0) N H i / cm (cid:1) (cid:46) .
5) and is more diffuse comparedto the A0–AVI cores. The entire core B region likelyrepresents a very large globule that is being swept awayvia ram-pressure stripping. However, this core regionmight be more diffuse than the other Rayleigh-Taylor4
Barger et al.
Figure 15.
Same as Figure 12, but for the core AIV region. fingers if it is being further disrupted by a turbulentwake that trails behind Complex A as it travels throughthe Galactic halo.6.3.
Kelvin-Helmholtz Instability Structures
In addition to Rayleigh-Taylor Instabilities, Kelvin-Helmholtz Instabilities are also influencing Complex A.As the outer layers of the complex “rubs” against thesurrounding coronal gas, small tangential perturbationsfrom shear-flow disturbances can form on its surface.If they become amplified, then some of affected gas willraise tangentially off the complex in the ± b CA directions.Elevated material can then be more easily swept away bythe surrounding coronal gas through ram-pressure strip-ping as high pressure zones form on the leading edge ofthese structures and low pressure zones on the trailingedge. This elevated gas additionally can be influenced byRayleigh-Taylor Instabilities in the direction of the hostgalaxy’s gravitational potential well if it is able to main-tain a gas density that is greater than the halo densityand if it is not overpowered by ram-pressure stripping.Kelvin-Helmholtz instabilities can affect all portionsof the complex that are directly sliding against the sur-rounding coronal medium, but their signatures are moredifficult to identify on the lower Galactic latitude half(or higher b CA half) of Complex A. This is because theyare occurring in tandem with Rayleigh-Taylor instabil-ities and ram-pressure stripping, which have a stronger morphological impact on this HVC as evident by thenumerous globules that hang from it (see Figure 7).We therefore only identify Kelvin-Helmholtz instabilitystructures on the higher latitude side of Complex A.However, we stress that these Kelvin-Helmholtz insta-bilities could be exacerbating the Rayleigh-Taylor struc-tures that form on the lower latitude edge of this com-plex.Three small Kelvin-Helmholtz structures branch off ofcores A0 and A1, which are marked in Figures 7. Allof these structures point roughly perpendicularly off thesurface of Complex A with a slight tilt in the direction ofComplex A’s motion. This is interesting as ram-pressureaffects should cause these structures to tilt in the anti-direction of motion, but interactions with the densergas near the Milky Way’s disk may have affected theirorientation. This unusual orientation is also shared bythe low-latitude globular that hangs off of core A0. Ascore A0 is the leading core, its leading edge is beingheated eroded away by direct interactions with densermaterial near the Milky Way’s disk. These interactionsassisted in the formation of the Rayleigh-Taylor glob-ular at ( l, b ) ≈ (132 ◦ , ◦ ) and the Kelvin-Helmholtzstructure at (133 ◦ , ◦ ) (see Figure 11). All three of theKelvin-Helmholtz structures are attached to cores A0and A1 are connected by a thin filament, indicating thatthey will soon detach and evaporate into the surround-ing coronal medium (see Figures 11 and 12). Addition- omplex A Figure 16.
Same as Figure 12, but for the gas distribution of the core AV region. ally, the higher H i column density sub-cores that haveformed at the tips of these structures might indicate thatthese they are developing or will develop Rayeigh-Taylorfingers.Two high-latitude “wings” protrude from Complex A(see Figure 6), one between cores AIII and AIV at( l, b ) ≈ (147 ◦ , ◦ ) (see Figure 20) and another offcore AVI at (160 ◦ , ◦ ) (see Figure 19). Because thestems of these wings extend perpendicularly off of Com-plex A, this indicates that they were formed by Kelvin-Helmholtz instabilities. These structures subsequentlybecame elongated due to interactions with the surround-ing coronal as this HVC fell through the Galactic halo.Interestingly, sub-H i cores have formed in the tips ofthese wings which may indicate that they are starting toform Rayleigh-Taylor fingers. The odd forward leaningmorphology of these wings could be a result of buoy-ancy instabilities that are causing this higher densitygas to fall faster toward the disk. However, in the caseof wing 1, unseen eddies or a low pressure zone in theturbulent wake that lies behind wing 2 could also becausing this wing to curl forward (see Figure 7).In a hydrodynamical simulation of gas streams, Mur-ray & Lin (2004) found that wings can form as a re-sult of evolving thermal and Kelvin-Helmholtz instabili-ties. They found that as the wings grow, they can curvein the direction of motion of the main cloud due to acombination of Rayleigh-Taylor instabilities and entan- glement with vortices that formed in the surroundingcoronal gas, which erode away the middle of the wingon its leading side. The numerous cloud fragments thatlie behind these wings indicates that there is substantialturbulent mixing behind them, presumably caused by awake that follows these wings. DISCUSSIONThe HVCs that are infalling onto the Milky Way willgenerally need to travel for a tens to hundreds of mil-lion years to reach the Galactic disk as they typically lie | z | (cid:46)
10 kpc above or below the disk (van Woerden et al.1999; Wakker 2001; Wakker et al. 2007, 2008; Thomet al. 2006, 2008; Smoker et al. 2011; Richter et al. 2015)and are moving with speeds of 50 (cid:46) | v z | (cid:46)
200 km s − relative to the disk. While they are traversing the Galac-tic halo, they are gradually eroding away into the sur-rounding coronal medium. Heitsch & Putman (2009)and Bland-Hawthorn et al. (2007) predict that HVCswith M H i < . M (cid:12) will become fully ionized throughKelvin-Helmholtz instabilities within τ KH (cid:46)
100 Myrand therefore they will not typically reach the MilkyWay’s disk. Kwak et al. (2011) project that up to 70%of the hydrogen in HVCs with masses of M H i (cid:38) M (cid:12) can remain neutral for a few hundred million years,which means that the large complexes are likely to sur-vive their journey. However, higher mass HVCs thathave a stream morphology, or that have a fractured sur-6 Barger et al.
Figure 17.
Same as Figure 12, but for the gas distribution of the core AVI. face are similarly vulnerable to rapid evaporation due totheir increased surface area. Additionally, HVCs can be-come even more vulnerable to their surroundings if theybecome fragmented as a result of thermal instabilities(see Murray & Lin 2004) or “shock cascade” processes(see Bland-Hawthorn et al. 2007; Tepper-Garc´ıa et al.2015).While hydrodynamical instabilities assist in the de-struction of HVCs, heat conduction (Vieser & Hensler2007; Armillotta et al. 2017), self-gravity, and Magneticfields (Chandrasekhar 1961; Grønnow et al. 2018) areall processes that can suppress them (see Pl¨ockinger& Hensler 2012 and Grønnow et al. 2018). As HVCsmove through hot halo gas, they are being heated viathermal conduction, advection, and ionizing radiation,which means that these instabilities are at least partiallysuppressed due conduction. Although self-gravity wouldhelp stabilize HVCs, it is unlikely that these complexesare embedded within dark matter halos as their corre-sponding H i Virial distances would place them millionsof parsecs away (Oort 1966; Freeman & Bland-Hawthorn2002).The net effect that magnetic fields have on shapingHVCs is uncertain as hydrodynamic instabilities havebeen found to be mildly (Banda-Barrag´an et al. 2016)and strongly (McCourt et al. 2015; Goldsmith & Pit-tard 2016) suppressed and even enhanced (Grønnowet al. 2017) in magnetohydrodynamic simulations (see Grønnow et al. 2018). It may be the case, however, thatmagnetic fields affect each kind of hydrodynamical in-stability differently. For instance, Banda-Barrag´an et al.(2016) and Grønnow et al. (2018) found that magneticfields inhibit Kelvin-Helmholtz instabilities, which helpsto protect clouds against ablation by reducing their con-tact with halo material. In the case of thermal instabili-ties, Ji et al. (2018) found that magnetic fields appear topromote thermal instabilities, which aids in their frag-mentation. Similarly, Gregori et al. (1999), Grønnowet al. (2017), and Grønnow et al. (2018) found that mag-netic fields enhanced Rayleigh-Taylor instabilities in the z direction. Regardless, hydrodynamical effects domi-nate over magnetohydrodynamics in shaping clouds dur-ing most of their journey through the Galactic halo withthe exception of when they are near the Galactic planebecause the compressed leading edge of these clouds am-plifies their magnetic field strength (see Grønnow et al.2017).While there is uncertainty as to whether or not allHVCs have a magnetic field, the Smith Cloud (Hill et al.2013), Magellanic Bridge (Kaczmarek et al. 2017), andLeading Arm (McClure-Griffiths et al. 2010) all have adetected magnetic field. However, as these three HVCsrepresent material that has been displaced from a galaxy(Smith Cloud from the MW: Fox et al. 2016; Magel-lanic Bridge from the Magellanic Clouds: see Putmanet al. 2012 for a review; Leading Arm from the SMC: omplex A Figure 18.
Same as Figure 12, but for the core B region.
Fox et al. 2018), these fields may have been inheritedfrom their parent galaxy (Milky Way: Haverkorn 2015;Small Magellanic Cloud: Lobo Gomes et al. 2015; LargeMagellanic Cloud: Mao et al. 2012). It is unknown ifthe HVCs that originate from an intergalactic mediumfilament or though halo gas condensations will have amagnetic field.No magnetic field has been directly measured for Com-plex A. Nonetheless, Verschuur (2013) placed indirectconstraints on the strength of the Complex A’s toroidalmagnetic field by assuming that its broad H i lines area causes by a combination of thermal broadening andmagnetic turbulence that results from Alfv´en waves. Inthat study, they surmised that the lack of H α emissiondetected in the Wisconsin H α Mapper (WHAM) North-ern Sky Survey (NSS) of the Milky Way (Haffner et al.2003) was an indication that this complex is colder than T H i < × K and therefore the H i emission shouldonly have narrow emission-line profiles of FWHM <
25 km s − . Assuming a distance of d (cid:12) ≈
200 pc toComplex A, they derived a field strength of B ≈ µ Gwith their model. However, there are two major issueswith their assumptions: (1) The WHAM NSS does notspan the kinematic extent of Complex A (WHAM NSS: − (cid:46) v LSR (cid:46) +100 km s − ), so no H α emission fromthis complex would be present in this survey. Bargeret al. (2012) mapped the H α emission in Complex Ausing the WHAM telescope and they detected H α emis- sion from the entire complex, which varied in strengthfrom 30 (cid:46) I H α (cid:46)
100 mR. Therefore, broad H i line pro-files with 25 (cid:46) FWHM (cid:46)
35 km s − is not surprising asH α emission peaks between 8,000 (cid:46) T H α, peak (cid:46) d (cid:12) = 200 pc.Barger et al. (2012) also found that the inferred level ofionization based on the H α emission could be producedby photoionization from the Milky Way and the extra-galactic background if core A0 lies roughly 6 . (cid:46) d (cid:12) (cid:46) . B (cid:46) µ G.While the predicting power of hydrodymanical simu-lations continue to improve, they are currently unableto fully resolve the detailed physics that are influenc-ing HVCs. Smooth particle hydrodynamic simulationsstruggle to produce high resolution models that incor-porate ram-pressure stripping, Kelvin-Helmholtz andRayleigh-Taylor instabilities, turbulent motions, andmagnetic fields (see Mocz et al. 2015) as these codescan suppress entropy generation, underestimate vortic-ity generation, and impede efficient gas stripping (Si-jacki et al. 2012). Adaptive mesh refinement codes havedifficulty modeling diffusion as two mediums rub pasteach other at supersonic bulk velocities (Mocz et al.8
Barger et al.
Figure 19.
Same as Figure 12, but for wing 1 (W1)—a the high-latitude cloud fragment that lies off the core AVI region at ( l, b ) ≈ (154 . ◦ , . ◦ θ ≈ . (cid:48)
1) and high sensitivity (log (cid:0) N H i / cm − (cid:1) ≈ . i z ≈ − Z = 0 . +0 . − . Z (cid:12) : Fox et al.2016) indicates that it likely originated from a Galacticfountain. Although the present mass of this complexof M total (cid:38) × M (cid:12) (neutral: Lockman et al. 2008;ionized: Hill et al. 2009) is much larger than anticipatedfrom the energetic processes occurring within the Galac-tic disk (Fox et al. 2016), hydrodynamical simulationssuggest that when its high metallicity gas mixes withthe surrounding coronal gas it can provide an avenue forthe halo to cool and condense onto the complex (e.g., Marinacci et al. 2010; Marasco et al. 2013; Fraternaliet al. 2015. The cooled coronal gas can accrete ontothe HVC, enabling it to grow as it travels through thehalo (Armillotta et al. 2016). Because this HVC hasa measured average magnetic field strength along theline-of-sight (LOS) of (cid:126)B LOS (cid:38) µ G (Hill et al. 2013), itcould be at least partially resistant to Kelvin-Helmholtzinstabilities. Nonetheless, Rayleigh–Taylor instabilitieshave been observed in the H i of this complex (Bettiet al. 2019).However, low metallicity halo clouds that are ineffi-cient at cooling will instead more easily erode into thehalo (e.g., Joung et al. 2012). HVC Complex A is onesuch low metallicity cloud ( Z = 0 . Z (cid:12) : Kunth et al.1994; Schwarz et al. 1995; van Woerden et al. 1999;Wakker 2001; Barger et al. 2012). This HVC furtherhas no detected magnetic field, so it is unknown if theyare suppressing or enhancing hydrodynamic instabili-ties along its length. In this study, we have resolvedthe H i morphology of this complex in unprecedenteddetail, which enable us to identify morphological struc-tures that are associated with ram-pressure strippingand thermal, Rayleigh-Taylor, and Kelvin-Helmholtz in-stabilities. This study provides the first opportunity totrace all of these hydrodynamic instability signatures ina low metallicity HVC that may have originated froman intergalactic-medium filament. omplex A Figure 20.
Same as Figure 12, but for wing 2 (W2)—a high-latitude cloud fragment that lies off the core AIV region at ( l, b ) ≈ (147 . ◦ , . ◦ (cid:0) N H i / cm (cid:1) ≈
20) cloudlet at (142 ◦ , ◦ ) that isassociated with M81 galaxy, which is enclosed within a red rectangle in the left-hand panels. This M81 emission has been removed fromthe center and right-hand panels. Resolution matters. Previous observations of this gasstream using the Leiden/Argentine/Bonn (LAB) surveywith a ∆ θ = 0 . ◦ ≤ ∆ θ ≤ (cid:48) , butthey were much less sensitive at 0 . (cid:46) T B , σ (cid:46) i θ = 16 . (cid:48) T B , σ = 43 mK per ∆v bin = 1 .
29 km s − or 3-sigmalog( N H i , σ / cm − ) = 18 . − (HI4PI Collaboration et al. 2016),the beam size of the GBT observation we are presentingin this study span an angular diameter that is a factorof 3 . × smaller. At the angular resolution and sensitiv-ity of the HI4PI survey, Rayleigh-Taylor instability fin-gers are difficult to identify (see Figure 3 of Westmeier2018). Without high resolution observations of HVCs,like the ones presented this study, we will not be able toidentify morphological and kinematical signatures of hy-drodynamic instabilities, which are needed understandthe survivability these complexes as they traverse theGalactic halo and anchor simulations. SUMMARYIn this study, we explore the kinematics and mor-phology of the neutral hydrogen gas of Complex A. Wepresent a kinematically resolved H i − ≤ v LSR ≤ −
90 km s − velocity rangethat spans a 600-square degree area across the sky. Thissurvey has a sensitivity of 17 . (cid:46) log( N H i , σ / cm − ) (cid:46) . − width. Wefinish with the main conclusions of our study:1. Bulk Motion:
There is a Galactic standardof rest frame velocity gradient of ∆v
GSR / ∆L =25 km s − / kpc along the ∆ L ≈ . (cid:104) a (cid:105) GSR = 55 km / yr , which will place thiscomplex at the Galactic plane in ∆ t (cid:46)
70 Myrs.2.
Ram-Pressure Stripping:
Numerous H i cloudlets along the Complex A exhibit morpho-logical signatures that are shaped by ram-pressurestripping. The cores A0, AI, AII, and AIII tendto be compressed in the direction of motion withdiffuse following gas (see Figures 6–14). Much ofthe gas that extends off the low-latitude edge ofthe complex is tilted in the anti-direction of mo-tion. This includes an “H i loop” that extends off0 Barger et al. of core A0, wispy gas that hangs from core AII,a relatively multi-core filament and a small fila-ment that branches off of core AIV, and the entirecore B region (see Figure 7).3.
Rayleigh-Taylor Instabilities:
We have iden-tified numerous Rayleigh-Taylor fingers that hangfrom the lower latitude edge of the Complex Astream. This includes a finger that hangs offcore A0 and curves upward in Complex A’s direc-tion of motion (see Figure 11), suggesting that thislow-latitude gas is interacting with higher densitygas near the Galactic disk. Fingers also extend offcores AI, AII, AIII, and AIV (see Figures 12, 14,and 15). The entire elongated and diffuse core Bregion has the morphology of a large globular thatbranches off core AVI and has been pushed back-ward due to ram-pressure stripping (see Figures 6and 18). Additionally, the high density H i sub-cores at the tip of the two high-latitude wings sug-gests that they could be forming globules.4. Kelvin-Helmholtz instabilities:
Because bothRayleigh-Taylor instabilities and Kelvin-Helmholtzinstabilities are simultaneously affecting the gas onthe lower latitude edge of Complex A, the Kelvin- Helmholtz signatures are difficult to isolate onthis edge of the complex. On the high-latitudeedge, there are two wings that extend tangen-tially from Complex A that were formed throughKelvin-Helmholtz instabilities. After their initialformation, ram-pressure stripping elongated thisgas and a combination of Rayleigh-Taylor insta-bilities and/or vortices in the surrounding coronalgas that are eroding caused them to curl slightlyin the direction of motion.We thank the Research Apprentices Program at theDepartment of Physics and Astronomy at Texas Chris-tian University, which enabled high school students toparticipate in this research. Barger received supportedthrough NSF grant AST 1203059. We thank Jay Lock-man and Kevin Blagrave for providing us the deepPlanck GBT H i data cubes before they were publiclyreleased. We also thank Katie Chynoweth for sharingher deep M81 GBT H i data cube with us. Facility:
Green Bank Telescope
Software:
GBTIDL and MPFIT .REFERENCES