Magnetized Gas in the Smith High Velocity Cloud
Alex S. Hill, S. A. Mao, Robert A. Benjamin, Felix J. Lockman, Naomi M. McClure-Griffiths
aa r X i v : . [ a s t r o - ph . GA ] S e p Accepted for publication in the Astrophysical Journal
Preprint typeset using L A TEX style emulateapj v. 5/2/11
MAGNETIZED GAS IN THE SMITH HIGH VELOCITY CLOUD
Alex S. Hill , S. A. Mao , Robert A. Benjamin , Felix J. Lockman , & Naomi M. McClure-Griffiths (Dated: September 13, 2013) Accepted for publication in the Astrophysical Journal
ABSTRACTWe report the first detection of magnetic fields associated with the Smith High Velocity Cloud.We use a catalog of Faraday rotation measures towards extragalactic radio sources behind the SmithCloud, new H I observations from the Green Bank Telescope, and a spectroscopic map of H α from theWisconsin H-Alpha Mapper Northern Sky Survey. There are enhancements in rotation measure of ≈
100 rad m − which are generally well correlated with decelerated H α emission. We estimate a lowerlimit on the line-of-sight component of the field of ≈ µ G along a decelerated filament; this is a lowerlimit due to our assumptions about the geometry. No RM excess is evident in sightlines dominatedby H I or H α at the velocity of the Smith Cloud. The smooth H α morphology of the emission at theSmith Cloud velocity suggests photoionization by the Galactic ionizing radiation field as the dominantionization mechanism, while the filamentary morphology and high ( ≈ α intensity of thelower-velocity magnetized ionized gas suggests an ionization process associated with shocks due tointeraction with the Galactic interstellar medium. The presence of the magnetic field may contributeto the survival of high velocity clouds like the Smith Cloud as they move from the Galactic halo tothe disk. We expect these data to provide a test for magnetohydrodynamic simulations of infallinggas. Subject headings:
ISM: kinematics and dynamics — ISM: structure — Magnetic fields INTRODUCTION
Infall of gas drives the evolution of galaxies(e.g. Wakker & van Woerden 1997; Blitz et al. 1999;Putman et al. 2012) and may be necessary to explain theobserved mass-metallicity relationship (Tremonti et al.2004). High velocity clouds (HVCs) likely tracethis inflow. HVCs are gas observed at veloci-ties inconsistent with Galactic rotation, defined var-iously in terms of their local standard of rest(LSR: | v LSR | >
90 km s − ) or deviation ( | v dev | >
50 km s − ) velocity (Wakker & van Woerden 1991,1997; Moss et al. 2013). Though primarily studied in H I (Wakker & van Woerden 1997), HVCs have a substantialor dominant ionized component, some at ∼ K (e.g.Tufte et al. 1998, 2002; Bland-Hawthorn et al. 2007;Hill et al. 2009; Barger et al. 2012) and some more highlyionized (e.g. Sembach et al. 2000, 2003; Fox et al. 2006;Shull et al. 2009; Lehner & Howk 2011).The Smith Cloud is perhaps the best HVC for trac-ing the interaction between the Galactic halo and in-terstellar medium (ISM). It has a cometary morphol-ogy, with a head near ( l, b ) = (39 ◦ , − ◦ ) and a tail ex-tending well past (45 ◦ , − ◦ ). From three independentmeasurements, the distance to its head is 12 . ± . CSIRO Astronomy & Space Science, Marsfield, NSW, Aus-tralia; [email protected], naomi.mcclure-griffi[email protected] Department of Astronomy, University of Wisconsin-Madison, Madison, WI, USA; [email protected] Jansky Fellow, National Radio Astronomy Observatory, So-corro, NM, USA Department of Physics, University of Wisconsin-Whitewater,Whitewater, WI, USA; [email protected] National Radio Astronomy Observatory, Green Bank, WV,USA; [email protected] tic plane and moving upwards towards the disk. L08also derived all three components of the space velocityof the cloud. The H I mass of gas mapped by L08is 10 M ⊙ , though there is considerable emission out-side the area they observed as well. While the originof the cloud is unknown, it has been the subject of muchspeculation including an association with the Sgr dwarf(Bland-Hawthorn et al. 1998) and a jet from the molecu-lar ring (Sofue et al. 2004). Nichols & Bland-Hawthorn(2009) suggested that an associated dark matter halocould confine the cloud. Its low nitrogen abundance sug-gests an extragalactic origin (Hill et al. 2009, but see dis-cussion in Section 6.1 below), while its prograde orbitinclined at only 30 ◦ suggests knowledge of the Galac-tic gravitational potential (Smith 1963, L08). Its calcu-lated orbit suggests that it survived a passage throughthe Galactic disk at a Galactocentric radius of ≈
13 kpc(L08).Hydrodynamical simulations suggest that HVCs likethe Smith Cloud should lose most of their neutral hy-drogen content after traveling .
10 kpc through theGalactic halo and that clouds need to be denser thanthe ambient medium by a factor of ∼ −
20 to sur-vive infall (Kereˇs & Hernquist 2009; Heitsch & Putman2009; Joung et al. 2012). If HVCs are disrupted onthese scales, much of the accreting gas likely feedsthe warm ionized medium (Benjamin & Danly 1997;Bland-Hawthorn 2009; Kwak et al. 2011; Henley et al.2012). However, although the Smith Cloud likely hasan ionized envelope as massive as the neutral cloud(Hill et al. 2009) and the H I is disrupted, & M ⊙ remains concentrated in a coherent H I structure.The few simulations of HVCs to date which have ex-plored the effects of magnetic fields in two (Konz et al.2002) or three (Santillan et al. 2004) dimensions (2D or3D) support analytic expectations that magnetic fieldsof at least a few µ G can stabilize a cloud against disrup-tion. In 2D simulations, Santillan et al. (1999) showedthat a magnetized disk can prevent an unmagmetizedHVC from penetrating into the disk due to increasedmagnetic tension from compressed field lines, althoughthe interaction drives waves on both sides of the disk.The orientation of the field both within the cloud andin the ambient ISM substantially affects its evolution(Kwak et al. 2009), so observations of the magnetic fieldgeometry within HVCs are a necessary constraint on ac-cretion models.Measuring magnetic fields in interstellar gas is diffi-cult. Attempts to measure Zeeman splitting of the H I
21 cm line, the best direct measure of magnetic fields inneutral hydrogen, have not been successful for the SmithCloud due to potential structure in H I within the beam.Faraday rotation measures the integrated line of sightmagnetic field in ionized gas weighted by the electrondensity. The measured rotation in the linear polariza-tion angle of a background radio source at wavelength λ is ∆ θ = RM λ , where the rotation measure isRM = 0 . Z observersource n e ( s ) B ( s ) · d s rad m − . (1)Here n e ( s ) is the electron density as a function of positionalong the line of sight (in cm ), B ( s ) is the magnetic field(in µ G), and s is the line-of-sight vector (in pc) from thesource to the observer; thus, a positive RM indicates afield pointing towards the observer.A recent all-northern-sky catalog of RMs towards ex-tragalactic points sources (Taylor et al. 2009) has al-lowed many studies of the interstellar magnetic field.McClure-Griffiths et al. (2010) used these data to mea-sure a line-of-sight magnetic field of B || ≥ µ G in anHVC in the Leading Arm of the Magellanic System.Here, we apply the technique of McClure-Griffiths et al.(2010) to the Smith Cloud to measure the magnetic fieldin a cloud interacting with the Galactic disk. In Sec-tion 2, we describe the H I , H α and Faraday rotationdata we use. We describe the spatial and kinematicstructure of the Smith Cloud in Section 3 and argue thatessentially all of the emission studied here is physicallyassociated with the cloud in Section 4. We then esti-mate electron densities and magnetic field strengths inSection 5. We discuss our results in Section 6 and sum-marize the paper in Section 7. DATA
For Faraday rotation, we use the catalog ofTaylor et al. (2009). They derived RMs from the NRAOVLA Sky Survey (NVSS; Condon et al. 1998) for ≈ − − for | RM | .
650 rad m − . At larger RMs, awrapping ambiguity can be of concern because the NVSSpolarization data sample only two frequencies. However,such large RMs are unlikely at the Galactic latitude ofthe Smith Cloud: at − ◦ < b < − ◦ , Taylor et al.(2009) list mean RMs of ≈ −
20 rad m − with a stan-dard deviation of <
90 rad m − , so | RM | = 650 rad m − is 7 σ from the mean of the distribution. Mao et al. (2010,2012), using measurements which better sample λ space, have confirmed that high latitude RMs from Taylor et al.(2009) do not suffer from the wrapping ambiguity.The H I data used here comes from a new survey ofthe Smith Cloud and its environs made with the 100meter Robert C. Byrd Green Bank Telescope (GBT) ofthe NRAO. The spectra cover 700 km s − around zerovelocity LSR. The velocity resolution is 0 .
65 km s − , andthe typical rms noise level is 90 mK in a 0 .
65 km s − channel. The angular resolution is 9 . α spectroscopic maps from the Wisconsin H-Alpha Mapper (WHAM) Northern Sky Survey (WHAM-NSS; Haffner et al. 2003). These data cover the entireregion we consider here with a 3 σ sensitivity of ≈ .
15 Rand a spectral resolution of 12 km s − . The 1 ◦ WHAMbeam is spaced on a grid with ≈ ◦ between pointing cen-ters, undersampling the image. The map presented herecovers v LSR . +70 km s − . Hill et al. (2009) presenteda WHAM H α map for a subset of this region coveringthe range −
70 km s − . v LSR . +130 km s − . We haveconsidered these data in our analysis but do not presentthem here. An estimated extinction correction in thesightlines and latitudes we consider would add ≈
40% tothe true H α intensities at b ≈ − ◦ (Hill et al. 2009) andless at higher latitudes. Because this extinction correc-tion is highly uncertain and affects derived densities andmagnetic field strengths by .
18% (see equations 4 and5 below), we do not perform an extinction correction. DESCRIPTION OF SMITH CLOUD ENVIRONS
The region around the Smith Cloud is morphologi-cally and kinematically complex. Here, we describe andidentify features evident in the data (Figures 1 −
4) andsketched in Figure 5. In Section 4, we interpret the mor-phology as the result of an interaction between the SmithCloud and the ambient ISM.
Neutral hydrogen emission
The Smith Cloud shows a gradient in v LSR from+100 km s − at its head to +70 km s − at its tail, whichcan be understood as the change of the Sun’s projectedmotion over the large angle. In the Galactic standardof rest (GSR) frame (the frame in which the Galacticcenter is at rest) the hydrogen has a uniform velocity, v GSR ∼ +230 km s − . In Figure 1, we show the H I emission at v GSR = +247 km s − . In addition to themain head-tail cloud extending from ( l, b ) ≈ (38 ◦ , − ◦ )to (47 ◦ , − ◦ ), there are several knots of emission spa-tially offset from the cloud but at the same v GSR . Onenear (55 ◦ , − ◦ ) is close to the major axis of the cloud,but there are also clumps of emission offset from the ma-jor axis of the cloud, near (38 ◦ , − ◦ ), (41 ◦ , − ◦ ), and(49 ◦ , − ◦ ). These three clumps are identified as clumpA, B, and C, respectively, in Figure 5.Lower-velocity H I emission at v LSR ≈ +40 km s − is also likely associated with the Smith Cloud. TheH I “decelerated ridge” along the northwest side of the G a l a c t i c L a t i t u d e
500 pc 0.000.050.200.501.001.502.00 (K)
Figure 1.
Taylor et al. (2009) RMs overlaid on GBT H I data. The grayscale shows H I emission in the v GSR = +247 km s − channel;the background is white in regions for which we have no GBT data. RMs are shown as red (RM >
0) and blue (RM <
0) circles. Themagnitude of the RM is proportional to the diameter of the circle, with the largest circles corresponding to | RM | ≥
200 rad m − . Thescale bar assumes a distance of 12 . Smith Cloud was identified as gas decelerated from theSmith Cloud by L08. Figure 2 shows additional rel-atively straight ridges in intermediate velocity H I at v LSR ≈ +40 km s − , identified as the “ l = 39 ◦ ridge”and “ b = − ◦ ridge” in Figure 5. These ridges extendaway from near the midpoint of each side at ≈ ◦ anglesrelative to the major axis of the Smith Cloud. The b = − ◦ ridge is somewhat difficult to disen-tangle from foreground emission due to its lower lati-tude. Thus, we focus on the l = 39 ◦ ridge, shown in alatitude-velocity diagram in Figure 2 c . The emission ex-tends from v LSR = +100 km s − at (39 ◦ , − ◦ ), the edgeof the main portion of the Smith Cloud, until it mergeswith the v LSR ≈ +40 km s − emission near b = − ◦ .The v LSR ≈ +40 km s − emission in Figure 2 is anti-correlated with the v GSR = +247 km s − emission near(41 ◦ , − . ◦ ), indicating interaction between gas at thetwo velocities. Ionized hydrogen emission
Hill et al. (2009) presented a map of H α emission at Because of the contamination of v LSR ≈ v LSR ≈ +100 km s − . The morphology of the H α emis-sion is smoother than that of the H I , with a typicalintensity of 0 . − . v LSR ≈
100 km s − H I emission, both along themain axis of the cloud and in clumps A, B, and C. Forcomparison, the typical integrated H I intensity in clumpB is ≈
40% of the integrated intensity along the main axisof the Smith Cloud. The relative lack of structure in H α compared to H I is probably due in part to the differingangular resolutions of 1 ◦ and 9 ′ , respectively.The H α emission at v LSR ≈ +40 km s − (Fig. 3) ismore structured than the +100 km s − emission. To theeast and south of the l = 39 ◦ and b = − ◦ H I ridgesis a roughly semicircular enhancement in H α (Fig. 3)called the “H α ridge” in Figure 5; this ridge is mostclear in Figure 4 (described in Section 5). The H α ridgeis separated from the H I in the l = 39 ◦ ridge by ≈ ◦ ,and the l = 39 ◦ ridge emission curves eastward outsidethe H α ridge near ( l, b ) = (38 ◦ , − ◦ ). The brightestemission in the H α ridge is in a filament ≤ I H α ≈ . − . ≈ . α ridge. Faraday rotation
Figure 2.
Panel a shows a comparison of H I at the Smith Cloud velocity and at lower velocities. The grayscale shows H I in the v LSR = +40 . − channel with a square root intensity scale, with black contours of the same data set at 1 . . v GSR = +247 km s − corresponds to v LSR = +115 km s − at ( l, b ) = (38 ◦ , − ◦ ) and v LSR = +98 km s − at (47 ◦ , − ◦ ). Panels b and c show longitude-velocity and latitude-velocity diagrams along the red lines in panel a . The raw RMs from Taylor et al. (2009) are shown inFigure 1. Off the major axis of the cloud, in the upperleft portion of the figure, the RMs are relatively largeand negative (shown as blue), while in the lower rightportion of the figure, the RMs are relatively small butstill negative. The distribution of RMs appears simi-lar in the main portion of the Smith Cloud and outside,suggesting that the v GSR ≈ +230 km s − H I does notcontribute significantly to the RMs. There are positiveRMs (shown as red) near the cloud but offset from thebrightest v GSR = +247 km s − emission, with particu-larly strong concentrations of positive RM near clump Band the H α ridge, along the l = 51 ◦ ridge, and to higherlongitudes, near ( l, b ) = (55 ◦ , − ◦ ). This emission iscoincident with the edge of the fainter H α emission as-sociated with the l = 51 ◦ ridge.The large-scale distribution of RMs is well fit by a lin-ear surface, which we interpret as a foreground contri-bution due to the magneto-ionic medium of the Galaxy(following McClure-Griffiths et al. 2010). To model thecontribution of the foreground to the Faraday rotation,we performed a 2D fit of all the RMs in Figure 1 asa function of l and b . The fit RMs are −
89 rad m − at (60 ◦ , − ◦ ) and −
27 rad m − at (30 ◦ , − ◦ ), consis- tent with the wide-area fits by Taylor et al. (2009). Wesubtracted this fit from the observed RMs; the differ-ence, RM HVC , represents magnetized ionized gas associ-ated with the Smith Cloud. We show RM
HVC in Fig-ure 4. There is no clear evidence of non-zero RM
HVC inthe main portion of the Smith Cloud, where the H I at v GSR ∼ +230 km s − is brightest.Figure 4 shows that the regions of enhanced RM arecoincident with enhancements in v LSR ≈ +40 km s − H α emission. Two regions of enhanced Faraday rota-tion are coincident with the H α ridge, with RM HVC ≈ +120 to +250 rad m − near ( l, b ) = (41 ◦ , − ◦ ) andRM HVC ≈ −
50 to −
100 rad m − near (47 ◦ , − ◦ ). TheRM enhancements are each coincident with H α emissionslightly fainter than the brightest emission in the H α ridge. We note that the enhancement in negative RMsnear (58 ◦ , − ◦ ) is in a region with bright v LSR ≈ α emission which extends off the figure towards the planeand Galactic east, so this RM enhancement is likely as-sociated with a foreground H II region. INTERACTION WITH THE HALO ISM
The morphology and kinematic structure of the fea-tures described in the previous section suggest that G a l a c t i c L a t i t u d e
500 pc 0.000.250.500.751.001.25 (R)
Figure 3.
Emission from ionized and neutral gas in the Smith Cloud region. The grayscale shows WHAM-NSS data integrated over+25 km s − < v LSR < +50 km s − , with the gray circles denoting WHAM beams. The orange and black contours show H I data exactlyas in Fig. 2. the gas responsible for the v GSR ≈ +230 km s − (or v LSR ≈ +100 km s − near the head of the cloud) andthe v LSR ≈ +40 km s − H I and H α emission are physi-cally associated and likely the result of interaction withthe ambient ISM (L08). This is most clearly indicated bythe anti-correlation of the v LSR ≈ +40 km s − H I emis-sion from the l = 39 ◦ ridge (black contours in Fig. 2)and the v GSR ≈ +230 km s − (orange contours). Theconnection between the v LSR = +40 and +100 km s − emission along the l = 39 ◦ ridge is also evident in Fig-ure 2 c . The similar morphology of the H α ridge and theH I l = 39 ◦ ridge, both at v LSR ≈ +40 km s − , stronglysuggests that they are, in turn, physically associated.The dominance of H α in some regions and H I in oth-ers may indicate shock ionization in some regions of thecloud or may reflect an emission measure or column den-sity threshold above which neutral gas is shielded fromphotoionization by the Galactic radiation field.The H I mass in the decelerated ridge is too large toaccount for with swept-up ISM gas that far from theGalactic plane. L08 thus argued that the deceleratedridge consists of material which has been stripped fromthe Smith Cloud and decelerated by the interaction withdisk gas. The same argument applies to the b = − ◦ ridge (which may be the same structure as the deceler-ated ridge) and the l = 39 ◦ ridge; the latter has an H I mass of ∼ × M ⊙ , assuming the gas is at the SmithCloud distance.The l = 51 ◦ ridge is connected in velocity space to themain emission from the cloud, as shown in the longitude-velocity diagram in Figure 2 b , suggesting physical as-sociation. The emission from the ridge at l ≈ ◦ , v LSR = +40 km s − connects to the tail of the H I emis-sion associated with the main Smith Cloud ( l = 44 ◦ , v LSR = +80 km s − ; see the orange contours in Fig-ure 2). However, in contrast to the bow wave-like ap-pearance of the l = 39 ◦ and b = − ◦ ridges with respectto the Smith Cloud, the H I and H α morphology of the l = 51 ◦ ridge does not clearly indicate interaction withthe v GSR ≈ +230 km s − H I . Therefore, it is possiblethat this ridge is a foreground feature.Because the majority of the v LSR ∼ +40 km s − emis-sion in both H I and H α is associated with the v GSR ≈ +230 km s − emission, we adopt the Smith Cloud dis-tance of 12 . ± . ELECTRON DENSITY AND MAGNETIC FIELDSTRENGTH
The correlation of enhancements in RM with deceler-ated H α structures associated with the Smith Cloud sug-gests that RM HVC is due to magnetic fields in the SmithCloud. To estimate the magnetic field strength, we first G a l a c t i c L a t i t u d e
500 pc 0.000.250.500.751.00 (R)
Figure 4.
Fit-subtracted RMs (RM
HVC ) overlaid on H α and H I data. The grayscale shows WHAM-NSS H α data integrated over+25 km s − < v LSR < +50 km s − . We have subtracted the estimated Sagittarius Arm contribution (eq. 2) and smoothed the H α observations from the WHAM beams in Fig. 3. The H I contours are as in Fig. 2. The red and blue circles denote RM HVC (Section 3.3);as in Fig. 1, the largest-diameter symbols have | RM HVC | = 200 rad m − . The green shapes denote regions with positive and negativeRM HVC values which we use to calculate B || in Table 1. Figure 5.
Schematic diagram of the features in the Smith Cloud,drawn to scale in comparison to Figs. 3 and 4. The v GSR ≈ +230 km s − H I emission from the Smith Cloud is outlined inorange. Prominent v LSR ≈ +40 km s − H I features are shown inblack, and the +40 km s − H α ridge in green. Red and blue cir-cular regions denote concentrations of positive and negative RMs(respectively) which are likely associated with the Smith Cloud. estimate the electron density from the H α emission. Toderive the H α emission at v LSR ≈ +40 km s − due tothe Smith Cloud, we estimate the H α contribution of the warm ionized medium in the foreground SagittariusArm assuming (Haffner et al. 1999) I Sgr ( b ) = I exp (cid:18) − DH tan | b | (cid:19) . (2)We performed a weighted linear fit of the H α emissionfor latitudes − . > tan b > − . − 35. We used the mean absolute de-viation about the median as the uncertainty in I Sgr . Here D is the distance to the emission; the emission is likelydistributed along distances from ∼ − 10 kpc because thearm is near tangency at this longitude (Benjamin 2008). D is also double-valued (G´omez 2006). The warm ion-ized medium scale height is H . Because H ≈ H ≈ 300 pc for the neutral gas (e.g.Cox 2005), the foreground contribution of the SagittariusArm to the H I at these latitudes is negligible.We derive the observed emission measure due to the Table 1 Magnetic field estimates l b N h RM HVC i h EM i L H + h B || i (rad m − ) (pc cm − ) (pc) ( µ G)41 ◦ − ◦ 17 +108 ± . ± . 04 220 +8 . ± . ◦ − ◦ − ± . ± . 08 220 − . ± . ◦ − ◦ 16 +90 ± . ± . 04 220 +8 . ± . ◦ − ◦ 30 +73 ± . ± . 03 1000 +4 . ± . Note . — h RM HVC i and h EM i are measured within the greenshapes in Fig. 4. L H + is an assumed upper limit based on the H α morphology. | B || | is the lower limit derived from equation (5).The number of WHAM beams in the region is N . Smith Cloud from the H α intensity in each beam,EM ≡ Z n e ( s ) ds = 2 . (cid:18) T K (cid:19) . (cid:18) I H α − I Sgr ( b )R (cid:19) pc cm − , (3)assuming T = 8000 K in the H α -emitting gas. Becausethe underlying density distribution n e ( s ) is uncertain, weassume the simplest case: n e ( s ) is a step function with avalue of n e or zero along the line of sight, and the pathlength in which n e ( s ) = n e is L H + ≡ f L , where the gasoccupies a fraction f of the volume along the total pathlength L . Then the electron density is n e = p EM /L H + . (4)If the gas density is determined by a series of compres-sions and rarefactions such as might be produced byturbulence, we would expect a lognormal distribution ofdensity (V´azquez-Semadeni 1994). As a test, we replacedthe uniform distribution assumed above with a lognor-mal distribution of electron density with a width appro-priate for the warm ionized medium (Hill et al. 2008).This yields only a ∼ . h B || i µ G = h RM HVC i . h n e i L H + = h RM HVC i . × ( L H + h EM i ) / , (5)under the additional assumption that the field does notvary along the line of sight with h RM HVC i , n e , and L H + in rad m − , cm − , and pc, respectively. Here h RM HVC i is the weighted mean of RM HVC for all sources in a de-fined region. The weights are w i = σ − i / (Σ i σ − i ), where σ i is the quadrature sum of the statistical uncertaintyin RM HVC ,i and 7 rad m − , an estimate of the intrinsicvariation in the extragalactic radio sources (Schnitzeler2010; Stil et al. 2011). There is an additional, systematicuncertainty due to our Milky Way foreground RM sub-traction which is not included in our uncertainty. Theuncertainties reported in Table 1 are the standard devi-ation of the weighted mean, (Σ i w i σ i ) / .We determine the average field for four regions. Ourresults are shown in Table 1. We choose L H + to bethe largest value which is reasonable based upon themorphology, so our estimates of the magnetic field arelower limits. Along the H α ridge (first two rows of Ta- ble 1), the brightest H α emission and the RM enhance-ment are one WHAM beam wide, so we assume that L H + < 220 pc, the projected beam size. The l = 51 ◦ ridge (third row of Table 1) is also . L H + would be smaller and B || larger than we estimate.The downstream (53 ◦ . l . ◦ ) region of enhanced H α and RM HVC is larger; we estimate L H + . DISCUSSION The morphologies of the H α and H I emission at dif-ferent velocities and of the RMs provide insight into thephysical processes at work. Shock ionization or photoionization? The H α emission at v LSR ≈ +100 km s − is rel-atively smooth, with the exception of the tip near(39 ◦ , − ◦ ), with an observed H α intensity of ≈ . − . α intensity fromthe Smith Cloud by a model of the Galactic ionizingradiation field (Putman et al. 2003). In this scenario,the gas from the cloud would be largely ionized in askin up to a constant emission measure as seen fromthe midplane, EM ⊥ , while gas beyond this thresholdwould be neutral. Assuming that the depth of theionized skin, L ⊥ = EM ⊥ n − e , is small enough so that L ⊥ / tan | b | is less than the width of the cloud projectedonto the plane, we have EM ⊥ = EM sin | b | . From the ob-served H α intensities, EM ⊥ ≈ . − . − (vary-ing based upon the assumed extinction correction andtemperature). This photoionization scenario explainsthe H α emission from most other HVCs in which H α has been detected (Tufte et al. 1998, 2002; Haffner 2005;Putman et al. 2003; Barger et al. 2012).The intensity of the v LSR ≈ +40 km s − emission fromthe H α ridge is less smooth, has a higher H α intensity of ≈ . − . I ,so it is more difficult to explain this emission as photoion-ization by the Galactic radiation field. Ionization relatedto shocks may better explain this emission, either di-rectly or through photoionization from shock emission(Bland-Hawthorn et al. 2007). Given the sound speedin 8000 K ionized gas of 10 km s − and the velocity ofthe Smith Cloud with respect to the corotating ISM of130 km s − (L08), shocks are expected, although the am-bient density is uncertain and presumably small.The velocity is sufficient to produce some Si IV , C IV ,N V , and O VI , with the ratios depending upon thefraction of the kinetic energy which converts to mag-netic rather than thermal energy (Allen et al. 2008;Henley et al. 2012). The shock is slow enough so thatthe majority of this emission would be in the shock, notthe post-shock cooling-zone, so the expected shock is ina regime in which the ratios of these lines are useful di-agnostics of the shock conditions. None of the sightlinesprobed with the Far Ultraviolet Spectroscopic Explorer by Sembach et al. (2003) or others are near the SmithCloud. The Si IV , C IV , and N V lines are all accessi-ble with the Cosmic Origins Spectrograph on the HubbleSpace Telescope . Optical forbidden line emission, acces-sible with WHAM, can also constrain shock ionizationscenarios. Hill et al. (2009) did not detect [O III ] λ III ] in a limited range of shockvelocities around 100 − 200 km s − .Hill et al. (2009) measured an optical line ratio[N II ] / H α = 0 . ± . 05 towards the tip of the SmithCloud, which they used to estimate a nitrogen abun-dance N / H = 0 . − . 44 times solar. This estimate wasbased on the assumption that the gas is photoionized.However, their measurement was towards the portion ofthe cloud at the leading edge, the brightest portion ofthe HVC in both H α and H I at v LSR ≈ +100 km s − .Both the +100 km s − H α intensity (0.43 R) and the ge-ometry suggest that shocks may be located within thisbeam, potentially making the photoionization assump-tion invalid. The [S II ] line width in this beam impliesa nonthermal velocity of 11 . ± . − (Hill et al.2009). Putman et al. (2003) measured [N II ] / H α = 0 . α map of Hill et al. (2009), this region has smootheremission more suggestive of photoionization. With anassumed temperature of 8000 − II ] / H α = 0 . . < (N / H) / (N / H) ⊙ < . 8. Wenote that the shocked region at the tip of the cloud maybe quite small; it is not sampled by any RMs used in thepresent paper. Magnetic field geometry The enhanced Faraday rotation is better correlatedwith filaments of v LSR ≈ +40 km s − H α emis-sion than any other emission tracer. To constrainthe cause, we estimate the electron column density, n e L H + . For the observed emission measures and assum-ing a constant electron density, the column density of v LSR ≈ +40 km s − H + is n e L H + ∼ (EM L H + ) / =4 × cm − ( L H + / 200 pc) / in the H α ridge, com-pared to (1 − × cm − for the diffuse, v LSR ≈ +100 km s − H α (Hill et al. 2009). Therefore, if L H + ∼ 200 pc (our best estimate; see Section 5) in the ridgeswhich contribute most to the observed RMs, n e L H + issimilar in the ridges and in the diffuse H α -emitting gas.We would then expect a comparable RM ∼ B || n e L H + across the cloud to that in the filament if B || were con-stant.Instead, the RMs through the bulk of the H I emissionfrom the Smith Cloud (Figures 1 and 5) are consistentwith the foreground. This could be explained by a larger B || in the H α ridge or by a magnetic field which is lessordered in the Smith Cloud than in the H α ridge, pro-ducing more field reversals along the line of sight. Eitherscenario could be explained by a physical compression ofthe gas, which would draw together and order field lines.In addition to the sightlines with | RM HVC | & 100 rad m − which are correlated with filaments, theRMs downstream of the H I and H α emission from theSmith Cloud (in the lower left quadrant of Fig. 4) are This corresponds to a mass integrated over +25 km s − 200 pc) / M ⊙ . positive compared to the foreground. This may suggestan extended, low column density wake which is ionizedand magnetized but at too low a column to detect inemission tracers. However, the association of this down-stream gas with the Smith Cloud is less certain than forthe H α ridge and the l = 39 ◦ and b = − ◦ ridges.Faraday rotation is sensitive only to magnetic fieldsin ionized gas. Although magnetized gas in the Galaxywith a low ionization fraction does not contribute sig-nificantly to Faraday rotation, gas with the ionizationfraction of the warm and cold neutral media is stillaffected by Lorentz forces due to ion-neutral collisions(Kulkarni & Heiles 1987; Ferri`ere 2001). Therefore, if afield is present in the neutral gas but does not have anevident RM signature due to the low n e , it could still bedynamically important. Because most of the RM detec-tions reported here are concentrated in narrow, possibly-shocked filaments, the magnetic field lines are likely com-pressed. Because the mass contained in the filaments islarger than can be explained by swept-up gas from theambient ISM, the filaments are most likely gas strippedand decelerated from the HVC.The positive values of RM HVC on the l ≈ ◦ sideof the H α ridge and the negative values of RM HVC onthe l = 47 ◦ side of the H α ridge (see Figure 4), com-bined with the near-zero RM HVC values along the ridgebetween these data, suggest a toroidal field pointing to-wards the observer at l ≈ ◦ , away from the observerat l ≈ ◦ , and perpendicular to the line of sight inbetween. In 2D magnetohydrodynamic (MHD) simu-lations, Konz et al. (2002) found that an HVC movingthrough a magnetized medium would establish a mag-netic barrier which provides thermal insulation betweenthe cool cloud and the hot ambient medium, reducingthe disruption of the cloud. A magnetic field shouldalso reduce Kelvin-Helmholtz instabilities in the inter-face between the cloud and the ambient medium. In theKonz et al. (2002) model, an HVC can compress fieldlines in a much weaker ambient field to create a field ofa few µ G in the head of a cloud, qualitatively similar tothe observations we report here.The peak magnetic field of & µ G in the SmithCloud is stronger than the typical field in the Galac-tic midplane of a few µ G, particularly the regular (non-turbulent) component of ∼ − µ G (Ferri`ere 2001). Es-timates of the field at | z | ≈ ≈ . . µ Gand typically ≈ µ G (Cox 2005; Mao et al. 2010, 2012;Jansson & Farrar 2012). However, these measurementsare all made locally and vary between magnetic spiralarms (Jansson & Farrar 2012). These weak Milky Wayhalo magnetic field strength measurements in an envi-ronment similar to that surrounding the Smith Cloudsupport our interpretation that the ambient field musthave been compressed in order to produce the strong ob-served field. Implications and future prospects The observed morphology of the neutral and ionizedgas and the Faraday rotation described here promise toprovide an excellent benchmark for 3D MHD simula-tions of HVCs entering the Galactic disk. For example,such simulations could help to determine whether the ob-served RM distribution is due primarily to stronger fieldsin the decelerated gas or to observational selection effectsrelated to either the orientation of the field or the elec-tron density. MHD simulations could also test whethersmall seed fields in the HVC could be amplified to pro-duce the observed fields, since existing models rely uponthe swept-up ambient field. Such models will improveour understanding of the role the magnetic field plays inregulating the disk-halo interaction.The technique for measuring magnetic fields in isolatedinterstellar clouds developed by McClure-Griffiths et al.(2010) and expanded here can be used in any region inwhich H α spectra (or some other estimate of the elec-tron column) and background RMs are available. TheTaylor et al. (2009) catalog of RMs covers the northernsky ( δ > − ◦ ) with ≈ − on average. Thetechnique works best in regions of the sky with rela-tively smooth large-scale features in Faraday rotation,as RMs from background sources do not provide depthinformation to separate out foregrounds. WHAM-NSSprovides H α spectra for δ > − ◦ , and the observa-tions for the WHAM Southern Sky Survey, which willfill in the remainder of the sky, are complete and nowbeing reduced (Haffner et al. 2010, Hill et al in prep).WHAM survey data only provide complete H α veloc-ity coverage to | v LSR | < 80 km s − , but WHAM is alsoused to obtain spectra of HVCs in other velocity win-dows (Tufte et al. 1998, 2002; Haffner 2005; Hill et al.2009; McClure-Griffiths et al. 2010; Barger et al. 2012).The new broadband backends on both the Karl G.Jansky Very Large Array (Perley et al. 2009) and theAustralia Telescope Compact Array (Wilson et al. 2011)make these telescopes efficient machines for measuringFaraday rotation towards point sources, making RMsaccessible with short exposure times over the full skyfor much fainter sources than Taylor et al. (2009) in-clude. Future telescopes with phased array feeds like theAustralian Square Kilometre Array Pathfinder (ASKAP;Johnston et al. 2008) and the Square Kilometre Arraywill provide far more RM measurements over the en-tire southern sky. In particular, the ASKAP Polariza-tion Sky Survey of the Universe’s Magnetism (POSSUM;Gaensler et al. 2010) is designed to provide a large, all-southern sky database of RMs which will be well-suitedto this kind of work. SUMMARY Through an analysis of Faraday rotation and spectro-scopic maps of H I and H α , we have analyzed the struc-ture of the neutral and magnetoionized gas in the SmithCloud at a range of velocities. The Smith Cloud is a lab-oratory in which to study the behavior of a gaseous cloudas it interacts with the high-altitude Galactic ISM. Wefound that neutral hydrogen which has been stripped anddecelerated from the cloud (Section 3.1; Figure 2) has acounterpart a few hundred pc downstream in ionized gas(Section 3.2; Figure 3) which is most likely shocked. Anenhancement in Faraday rotation, relative to the smoothbackground, is coincident with the decelerated ionizedgas (Section 3.3; Figure 4). Using these observations, wemeasured a magnetic field in the cloud with a line-of-sight component & µ G (Section 5). We interpret ourresults as a signature of compressed and likely shockedgas as it interacts with the ambient ISM. This HVC is an excellent target for future observations of shock diag-nostics and MHD simulations to further understand therole magnetic fields play in regulating the accretion ofgas into galactic disks.We acknowledge useful discussions with A. J. Fox, B.M. Gaensler, J. A. Green, V. A. Moss, and E. K. Braden.The National Radio Astronomy Observatory is a facilityof the National Science Foundation operated under a co-operative agreement by Associated Universities, Inc. TheWisconsin H-Alpha Mapper is supported by the NationalScience Foundation. Facilities: GBT, VLA, WHAM REFERENCESAllen, M. G., Groves, B. A., Dopita, M. A., Sutherland, R. S., &Kewley, L. J. 2008, ApJS, 178, 20Barger, K. A., Haffner, L. M., Wakker, B. P., et al. 2012, ApJ,761, 145Benjamin, R. A. 2008, ASPC, 387, 375Benjamin, R. A., & Danly, L. 1997, ApJ, 481, 764Bland-Hawthorn, J. 2009, IAUS, 254, 241Bland-Hawthorn, J., Sutherland, R., Agertz, O., & Moore, B.2007, ApJ, 670, L109Bland-Hawthorn, J., Veilleux, S., Cecil, G. N., et al. 1998,MNRAS, 299, 611Blitz, L., Spergel, D. N., Teuben, P. J., Hartmann, D., & Burton,W. B. 1999, ApJ, 514, 818Boothroyd, A. I., Blagrave, K., Lockman, F. J., et al. 2011, A&A,536, 81Condon, J. J., Cotton, W. D., Greisen, E. W., et al. 1998, AJ,115, 1693Cox, D. 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