Highly Ordered and Pinched Magnetic Fields in the Class 0 Proto-Binary System L1448 IRS 2
Woojin Kwon, Ian Stephens, John Tobin, Leslie Looney, Zhi-Yun Li, Floris van der Tak, Richard Crutcher
DDraft version May 29, 2019
Typeset using L A TEX preprint style in AASTeX61
HIGHLY ORDERED AND PINCHED MAGNETIC FIELDS IN THE CLASS 0 PROTO-BINARYSYSTEM L1448 IRS 2
Woojin Kwon,
1, 2
Ian W. Stephens, John J. Tobin,
4, 5, 6
Leslie W. Looney, Zhi-Yun Li, Floris F. S. van der Tak,
9, 10 and Richard M. Crutcher Korea Astronomy and Space Science Institute (KASI), 776 Daedeokdae-ro, Yuseong-gu, Daejeon 34055, Republic ofKorea University of Science and Technology, Korea (UST), 217 Gajeong-ro, Yuseong-gu, Daejeon 34113, Republic of Korea Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA, USA National Radio Astronomy Observatory, 520 Edgemont Rd., Charlottesville, VA 22903, USA Homer L. Dodge Department of Physics and Astronomy, University of Oklahoma, 440 W. Brooks Street, Norman,Oklahoma 73019, USA Leiden Observatory, Leiden University, P.O. Box 9513, 2300-RA Leiden, The Netherlands Department of Astronomy, University of Illinois, 1002 West Green Street, Urbana, IL 61801, USA Astronomy Department, University of Virginia, Charlottesville, VA 22904, USA SRON Netherlands Institute for Space Research, Landleven 12, 9747 AD Groningen, The Netherlands Kapteyn Astronomical Institute, University of Groningen, PO Box 800, 9700 AV Groningen, The Netherlands (Received ...; Revised ...; Accepted ...)
ABSTRACTWe have carried out polarimetric observations with the Atacama Large Millimeter/submillimeterArray (ALMA) toward the Class 0 protostellar system L1448 IRS 2, which is a proto-binary embed-ded within a flattened, rotating structure, and for which a hint of a central disk has been suggested,but whose magnetic fields are aligned with the bipolar outflow on the cloud core scale. Our highsensitivity and high resolution ( ∼
100 au) observations show a clear hourglass magnetic field mor-phology centered on the protostellar system, but the central pattern is consistent with a toroidal fieldindicative of a circumstellar disk, although other interpretations are also possible, including field linesdragged by an equatorial accretion flow into a configuration parallel to the midplane. If a relativelylarge disk does exist, it would suggest that the magnetic braking catastrophe is averted in this system,not through a large misalignment between the magnetic and rotation axes, but rather through someother mechanisms, such as non-ideal magneto-hydrodynamic effects and/or turbulence. We have alsofound a relationship of decreasing polarization fractions with intensities and the various slopes of thisrelationship can be understood as multiple polarization mechanisms and/or depolarization from achanging field morphology. In addition, we found a prominent clumpy depolarization strip crossingthe center perpendicular to the bipolar outflow. Moreover, a rough estimate of the magnetic fieldstrength indicates that the field is strong enough to hinder formation of a rotationally supporteddisk, which is inconsistent with the feature of a central toroidal field.
Corresponding author: Woojin [email protected] a r X i v : . [ a s t r o - ph . S R ] M a y Kwon et al.
Keywords:
ISM: magnetic fields — submillimeter: ISM — stars: protostars — stars:formation
ASTEX Magnetic fields in L1448 IRS 2 INTRODUCTIONMagnetic fields are thought to play a significant role in star formation on all scales from the cloud( ∼ ∼
100 au). For example, it has been found that magnetic field directions arewell ordered and typically perpendicular to parsec-scale filamentary structures (e.g., Palmeirim et al.2013), which indicates that magnetic fields are important to form such intermediate scale structures.Also, the magnetic energy is comparable to the kinetic energy down to a few thousand au scales (e.g.,Li et al. 2014; Pattle et al. 2017).In addition, magnetic fields can affect circumstellar disk formation at the early protostellar stages.Hull et al. (2013) found that magnetic field directions of 16 young stellar objects (YSOs) are ratherrandom with respect to their bipolar outflows on a few hundred au scales, although a couple ofexamples with hourglass morphology magnetic fields aligned to its bipolar outflow had been knownat the time (Girart et al. 2006; Stephens et al. 2013). Later, it has been suggested that the magneticfield directions of YSOs can be understood with the existence of an extended disk structure at theyoungest YSOs, the so-called Class 0 YSOs (e.g., Segura-Cox et al. 2015). For example, L1527 has amagnetic field morphology perpendicular to the bipolar outflow and has an extended Keplerian disk(radius ∼
54 au) (Tobin et al. 2012; Hull et al. 2013; Ohashi et al. 2014), whereas L1157 that has amagnetic field aligned with the bipolar outflow has no disk structure larger than 15 au (Stephens et al.2013; Tobin et al. 2013). Such features have been explained by magnetic braking, which can be soefficient in YSOs having a magnetic field aligned with the bipolar outflow that a rotation-supporteddisk structure is largely suppressed at the early stages (e.g., Mellon & Li 2008; Hennebelle & Ciardi2009; Joos et al. 2012; Maury et al. 2018), often called the magnetic braking catastrophe of diskformation.However, not all Class 0 YSOs fit the interpretation connecting the magnetic field morphology andthe disk structure. For example, although a rotating disk-like structure has been detected aroundL1448 IRS 2 and its central binary companion (e.g., Tobin et al. 2015, 2016), the magnetic fielddirection detected on 500–1000 au scales is mostly aligned with the bipolar outflow (Hull et al. 2014),in which magnetic braking is expected to be efficient. In addition, Davidson et al. (2014) reportedthat the most preferred magnetic field for the Class 0 YSO L1527, which has a large Kepleriandisk (Tobin et al. 2012; Ohashi et al. 2014), is a weak field aligned with the bipolar outflow, not aperpendicular field, when considering magnetic fields of large ∼ ∼ Kwon et al. observations of the Atacama Large Millimeter/submillimeter Array (ALMA) toward L1448 IRS 2focusing on how the magnetic fields change on 100 au scales. TARGET AND OBSERVATIONSSeveral Class 0 YSOs with flattened envelope structures have been identified by Tobin et al. (2010)in the Perseus molecular cloud at a distance of about 300 pc (Zucker et al. 2018; Ortiz-Le´on et al.2018). Of these envelopes, L1448 IRS 2 has the clearest flattened structure and has been observedin polarization indicating a magnetic field aligned with the bipolar outflow (Hull et al. 2014). Pre-vious studies have imaged a large, extended bipolar outflow originating from the target at variouswavelengths: e.g., near IR observations using the
Spitzer Space Telescope (Tobin et al. 2007). Inaddition, Kwon et al. (2009) reported that grains have significantly grown based on the dust opac-ity spectral index estimated from 1 and 3 mm observations of the Combined Array for Research inMillimeter-wave Astronomy (CARMA). Regarding polarimetric observations, Caltech SubmillimeterObservatory SHARP observations detected 350 µ m continuum polarization perpendicular (inferredmagnetic field parallel) to the bipolar outflow with a 10 (cid:48)(cid:48) angular resolution (Chapman et al. 2013).Hull et al. (2014) also detected polarization in the same direction, particularly in the blueshifted lobeon the northwest side from the center with an angular resolution of ∼ (cid:48)(cid:48) .Polarimetric observations toward L1448 IRS 2 in ALMA Band 6 were made on 2016 November 12and 14 (2016.1.00604.S, PI: Woojin Kwon). Individual tracks were run over 3 hours to achieve a goodparallactic angle coverage. HH 211 was observed simultaneously with L1448 IRS 2 and shared thesame phase calibrator; HH 211 will be reported in a separate paper. The November 12 and 14 trackshave 42 and 40 antennas in the array, respectively. J0238+1636 was used as polarization calibratorand flux calibrator. Its flux was set to 1.085 Jy at 233 GHz with a spectral index of about -0.45.J0237+2848 and J0336+3218 were bandpass and phase calibrators, respectively. All the executionblocks were calibrated separately and combined when making images. Images were made using aBrigg’s weighting with a robust parameter of 0.5. We found this weighting was a good compromisebetween resolution and sensitivity. The final image has a synthesized beam of 0 . (cid:48)(cid:48) × . (cid:48)(cid:48)
37 (PA =9 . ◦ ). The noise levels of Stokes I (total intensity), Q, and U maps are ∼ . ∼ . ∼ . − , respectively. The polarization intensity map achieved by the Stokes Q and U mapswith a debias using the Stokes Q and U map noise level has a RMS noise of 0.008 mJy beam − . RESULTS3.1.
Magnetic field morphology
Magnetic fields are inferred in perpendicular to polarizations of dust thermal emission in millime-ter/submillimeter wavelengths, because non-spherical dust grains are aligned with their minor axis(spin axis) parallel to magnetic fields (e.g., Lazarian & Hoang 2007). In Figure 1, we rotate thepolarization directions by 90 degree and infer the magnetic fields. The magnetic field morphologyshows a beautifully clear hourglass morphology perpendicular to the elongated structure. Consistentwith previous CARMA observations (Hull et al. 2014), the region northwest of the center shows apoloidal field, which is approximately aligned with the bipolar outflow direction. However, in thecentral region the field direction rapidly changes to an orientation perpendicular to the bipolar out- Assuming magnetic field grain alignment, the magnetic field directions are inferred by 90 ◦ rotation of polarizationdirections. ASTEX Magnetic fields in L1448 IRS 2 − ,which corresponds to an optical depth τ ∼ . T = 30 K, so the polarization maynot be caused by extinction. Therefore, the magnetic fields may be inferred by 90 degree rotationof polarization directions, which results in a pattern in the central region that is broadly consistentwith a toroidal magnetic field projected in the sky plane. On the other hand, self-scattering cannotbe ruled out by the data set taken only in Band 6. However, note that even in this case, it mayprovide indirect evidence for the presence of a circumstellar disk(-like) structure at the center, sinceself-scattering has been found so far only toward disks with large grains. Therefore, the main resultof this paper does not change. The ambiguity will be tackled by further polarimetric observationsin different wavelengths. Indeed, we have carried out polarimetric observations in ALMA Band 3 aswell, and the preliminary results support the interpretation of magnetic fields. These results will bediscussed in a following paper (Kwon et al. in prep).In addition, there is the possibility that the central polarozation pattern is produced by magneticfield lines that are dragged by an equatorial accretion flow into a configuration that is parallel tothe midplane of the system. However, we believe this possibility is less likely than the rotationally-induced toroidal field interpretation because there is already evidence for significant velocity gradientalong the equatorial plane (see, e.g. Fig. 11 of Tobin et al. 2018). Nevertheless, higher resolution lineobservations are needed to firmly establish whether the intriguing field orientation near the centeris produced by accretion or rotation. Note the seven vectors around the source A within or on theinner most contour in the right zoomed-in plot of Figure 1. A detailed modeling may be required toconstrain the structure size of a toroidal magnetic field. However, based on the number of vectors Kwon et al.
RA (J2000) +30°45'10"12"14"16"18" D e c ( J ) L1448 IRS 2
300 AU . mm i n t e n s i t y ( m J y b e a m ) RA (J2000) +30°45'11"12"13"14"15"16" D e c ( J ) L1448 IRS 2
A B
300 AU . mm i n t e n s i t y ( m J y b e a m ) Figure 1.
Magnetic field morphology around L1448 IRS 2. The green vectors have been rotated by 90 ◦ from the polarization directions to indicate the inferred magnetic field direction. Vectors of 2 σ level or betterdetections are marked every 0 . (cid:48)(cid:48) , which is comparable to the Nyquist sampling. The gray scale and blackand white contours present the total intensity distribution with levels of 2, 3, 5, 9, 17, 33, 65, and 129 times0.1 mJy beam − . The blue and red contours are CO 2-1 intensity distributions integrated in velocity rangesof − . − , respectively, at levels of 3, 4, 5, 7, and 9 times 11 K km s − . Thecentral region zoomed-in is presented in the right. The synthesized beams of the CO and the continuumdata are marked in the bottom right corner in blue and white, respectively. The binary system positions are2A(03:25:22.407 +30:45:13.21) and 2B(03:25:22.363 +30:45:13.12). showing the shifted direction, which is 1.5 of the beam area, the central disk could be up to about50 au in radius: considering a beam smoothing, the real structure would be about 0.5 of the beamarea, so r = (0 . A beam /π ) . d , where A beam is the beam area and d is the target distance. Furthermolecular line observations at a high angular resolution will provide the info on how large the diskis, if it existed, and whether it is rotationally supported.Note that the self-scattering polarization pattern is expected to be parallel to the minor axis ofan inclined disk (e.g., Yang et al. 2016; Kataoka et al. 2017; Stephens et al. 2017b), which meansthat after 90 ◦ rotation, the corresponding B vectors look like a toroidal feature. If the toroidal fieldinterpretation is correct, it would indicate that rotation has become fast and energetic enough towind up the field lines, which likely signals the formation of a rapidly rotating disk. If the dust self-scattering interpretation is correct, it would indicate that grains in the flattened structure on the scaleof several tens au have grown to roughly 100 µ m sizes or more, which again would favor the existenceof a rotationally supported disk that is conducive to grain growth through a higher density and longertime compared to a dynamically collapsing inner envelope. Recently, some other ALMA polarimetricobservations have also presented polarization patterns of self-scattering or toroidal magnetic fieldsin the central regions of Class 0 and I YSOs with a disk (Lee et al. 2018; Sadavoy et al. 2018) anddisk candidates (Cox et al. 2018). Additional polarimetric observations at different wavelengths willallow us to distinguish magnetic field alignment from self-scattering (e.g., Alves et al. 2018, , Kwonet al. in prep).The contours in the right panel of Figure 1 also show the binary companion (L1448 IRS 2B), whichis separated from the primary by about 0 . (cid:48)(cid:48) (corresponding to ∼
180 au at the target distance)toward the west. This companion is less bright in the 1 mm continuum and has been detected at
ASTEX Magnetic fields in L1448 IRS 2 Figure 2.
Stokes Q, Stokes U, and polarization intensity maps are on left, in the middle, and on right,respectively. The contours present distributions of total intensity (Stokes I) at the same levels of Fig. 1. Thecolor edges are for the Stokes Q and U maps (the left covering minus and plus values) and the polarizationintensity map (the right). . (cid:48)(cid:48) × . (cid:48)(cid:48)
25 (PA: 21 ◦ ), which is slightlybetter than the polarimetric continuum data. Since these observations lack short baselines, onlythe cavity walls were detected, as the extended features between the walls (as detected in Tobinet al. 2015) are filtered out. Interestingly, the less bright companion L1448 IRS 2B seems to bemore coincident with the bipolar outflow. While the redshifted lobe is primarily centered on thecombination of both L1448 IRS 2A and 2B, the main blueshifted lobe seems to be centered on L1448IRS 2B only. However, there is a weak blueshifted feature from the L1448 IRS 2A and along the leftcontinuum branch. It is possible that the blueshifted component from L1448 IRS 2A might be in thesame velocity regime of the ambient cloud and thus could be filtered out by the interferometer.3.2. Polarization intensity and fraction
In Figure 2, the outflow cavity walls have relatively high polarization intensities, while there is astrip across the center, almost perpendicular (P.A. ∼ ◦ ) to the bipolar outflow (P.A. ∼ ◦ atlarge scales; Stephens et al. 2017a), with weak polarization signal lower than a few tens µ Jy beam − .This depolarized strip is shown in more detail in Figure 3 and discussed below. The polarizationintensity is not symmetrically distributed. The central region of the total intensity peak has thehighest polarization intensity. Also, the region south of the depolarized strip to the east of the centeris high in polarization intensity. In addition, polarization intensity is very clumpy compared to thetotal intensity distributions, which is indicative that polarization is significantly affected by the localenvironment.As shown in Figure 3, the depolarization regions of L1448 IRS 2 clearly appear at the centralregion and along the strip perpendicular to the bipolar outflow, whose polarization fractions areonly a few percent or less. The central region with the binary system presumably has the mostcomplicated magnetic fields. Also, as addressed, the magnetic fields are changing from aligned toperpendicular with respect to the bipolar outflow going from large to small scales. These complicatedpolarization patterns that are smaller than the beam reduce the measured polarization fraction: i.e., Kwon et al.
RA (J2000) +30°45'10"12"14"16"18" D e c ( J ) L1448 IRS 2
300 AU P o l a r i z a t i o n F r a c t i o n Figure 3.
Polarization fraction distribution indicated in color scales. The other plot components arethe same as Fig. 1. The schematic diagram on the right illustrates hourglass morphology magnetic fieldsdetected in different alignment: small scales in red and large scales in blue. beam smearing. A complicated magnetic field could also be caused by turbulence. However, in thecase of the turbulence-induced, depolarized regions are rather randomly distributed (Lee et al. 2017).On the other hand, the depolarized strip is similar to the case of an inclined cloud with a hourglass-shaped magnetic field (Kataoka et al. 2012). Indeed, Tobin et al. (2007) reported that the bipolaroutflow is inclined by about 57 ◦ (where 90 ◦ indicates a bipolar outflow on the sky plane). Theinferred inclination makes it possible for the radially pinched field lines along a given line of sight toproduce polarizations that cancel one another, yielding a less polarized equatorial region (see Fig.6d ofKataoka et al. 2012). Furthermore, rotation of a cloud introduces a mis-alignment to the depolarizedstrip (see Fig.10d of Kataoka et al. 2012), which is broadly consistent with the depolarized featureshown in Figure 3.Depolarization can also occur due to high optical depth (see Fig.3 of Yang et al. 2017), but this isnot likely the case here since the optical depth is expected to be low along this depolarization strip: τ ∼ .
08 even at the highest contour level assuming T = 30 K.In addition, along the depolarized strip, there are several depolarized clumps, whose sizes are notresolved at our angular resolution, as seen in the white in Figure 3. The clumps are separatedby ∼ . (cid:48)(cid:48) , which corresponds to about 240 au at the distance of Perseus. These clumps mayindicate relatively more turbulent areas with chaotic magnetic fields and/or areas with magneticfields pointing along the line of sight direction. They may even be de-magnetized “islands” producedby reconnection of sharply pinched magnetic field lines (see Fig. 7 of Suriano et al. 2017), althoughdetailed exploration of magnetic reconnection is beyond the scope of this paper.The highest polarization fractions, reaching levels of up to 40%, are located near the border of thenorthwest and southeast part of the envelope. Such high polarization fractions can occur only by ASTEX Magnetic fields in L1448 IRS 2 P frac ) with Stokes I intensities ( I ) overall. Thistrend has been reported by many previous polarimetric observations of arc-second (e.g., Girart et al.2006; Kwon et al. 2006; Liu et al. 2013; Hull et al. 2014; Galametz et al. 2018) and tens arc-secondor coarser angular resolutions (e.g., Dotson 1996; Collaboration et al. 2016; Soam et al. 2018). Goingtoward a central denser region, dust grains get larger causing less alignment in a magnetic field,optical depth increases, and/or magnetic fields likely become complicated, which all result in a lowerpolarization fraction. Note that dust grains getting larger, such as above 10 µ m in disk conditions,cause less alignment with the magnetic field because the Larmor precession rate becomes slower thanthe gas randomization rate (Hoang & Lazarian 2016; Tazaki et al. 2017).Regarding power-law indices of the relationship, when polarized emission is dominated from thesurface of a structure, the polarization fraction is expected to be inversely proportional to the intensityin the optically thin case ( P frac ∝ I − ): e.g., polarization caused by dust grains aligned by magneticfields due to the interstellar radiative torque (RAT, Lazarian & Hoang 2007) mainly around themolecular cloud surface. Interestingly, Figure 4 shows multiple slopes that encompass the limitsof the distributions in a qualitative manner. In the regime fainter than about 1 mJy beam − , theslope is roughly − .
4. This region corresponds to the area from the second lowest contour to aboutthe fourth contour in the 1 mm continuum map of Figure 3. The slope shallower than − ∼ ∼
10 mJy beam − ,the slope is close to −
1, which is indicative that the regions have no further significant polarization.Approaching to 10 mJy beam − , the slope becomes a little bit steeper than − s = − .
3. This canbe interpreted as polarization directions changing, resulting in depolarization. Indeed, the region iswhere the field directions switch from the poloidal to the toroidal pattern. For the central region thatis brighter than 10 mJy beam − the slope is − .
5. This shallower slope than − DISCUSSION4.1.
Magnetic field strength Kwon et al.
Figure 4.
Polarization fraction versus intensity. Each blue circle represents a pixel value, and the blacklines indicate individual power-law slopes, not fitting results. Data points of intensities greater than 3 σ andpolarization intensities larger than 2 σ have been selected. We estimate a very rough magnetic field strength using the Davis-Chandrasekhar-Fermi (DCF)method (Davis 1951; Chandrasekhar & Fermi 1953). We are aware that the field orientations can beaffected by outflows and gravitational collapse near a protostellar system, which will likely degradethe accuracy of the DCF method, but quantifying such effects would require more detailed dynamicalmodeling that is beyond the scope of this paper. In the DCF technique, the magnetic field strengthis estimated based on the dispersion of magnetic fields with respect to the background field directioncompared to its turbulence in a given density medium. In other words, a medium at a given densityand a turbulence indicated by a non-thermal linewidth would have a stronger magnetic field strengthwhen it has a smaller field position angle dispersion: the plane-of-sky strength of a magnetic field B P OS = Q √ πρ δV /δφ ≈ . (cid:112) n ( H )∆ V /δφ [ µ G], where Q , ρ , δV , δφ , and n ( H ) are a factorof order unity, the gas density, the non-thermal velocity dispersion in km s − , the position angledispersion of polarizations, and the molecular hydrogen number density, respectively (e.g., Crutcheret al. 2004; Ostriker et al. 2001).First, for the background large-scale fields we smoothed the Stokes Q and U maps with a nine-times larger beam (extended in both major and minor axes of the original beam by a factor ofthree), which is comparable to a half of the width across the continuum structure. This provides areasonable background field morphology (Pattle et al. 2017). As Figure 5 shows, the smoothed fields ASTEX Magnetic fields in L1448 IRS 2 − (thick gray contours in Fig. 5), for estimating the magnetic field strength. Themeasured dispersion is estimated as 10 ◦ (Fig. 5, right). Second, for estimating the number densityof H we utilized the dust continuum. The total continuum flux density of the area between 0.3 and6.5 mJy beam − , which is 18 . M T = F ν D /κ ν B ν ( T d ), where F ν , D , κ ν , B ν , and T d are the flux density, distance,mass absorption coefficient, blackbody radiation intensity, and dust temperature, respectively. Using F ν = 91 mJy, D = 300 pc, κ ν = 0 .
01 cm g − at 233 GHz (Ossenkopf & Henning 1994) assuminga gas-to-dust mass ratio of 100, and T d = 30 K (Kwon et al. 2009), the total mass is estimated tobe 0.08 M (cid:12) . In addition, assuming a cylinder with the profile of the continuum feature, the totalvolume would be 9 × cm . Therefore, we derive the volume density ρ ≈ . × − g cm − , whichcorresponds to n ( H ) ≈ . × cm − . We do not have an observational non-thermal linewidth, butit may be reasonable to adopt the trans-sonic velocity at 30 K: ∼ . − . These values result inthe magnetic field strength in the plane of the sky of about 640 µ G, with the relationship following: B P OS ≈ µ G (cid:16) n ( H )5 . × cm − (cid:17) . (cid:16) ∆ V . (cid:17)(cid:16) ◦ δφ (cid:17) . (1)Furthermore, we estimate the magnetic braking time scale of the presumed disk structure at thecenter, when the rotation velocity decreases by a half (e.g., Basu & Mouschovias 1994). The Alfv´enspeed follows the relationship, v A = B √ πρ = 0 .
43 km/s (cid:16) B µ G (cid:17)(cid:16) . × − g/cm ρ (cid:17) . . (2)In addition, the central mass surrounded by the inner thick gray contour in Figure 5 is estimatedas 0.05 M (cid:12) based on the total flux density of 55 mJy. This mass is rather uncertain: it could beoverestimated because the central region is warmer than the outer region and could be underestimatedbecause the very central region ( < a few au in radius) would be optically thick even in millimeterwavelengths. On the other hand, the central rotating structure could be much smaller than theinner region considered here. The same mass beyond the central area is extended up to the intensityof ∼ . − , which is about 0 . (cid:48)(cid:48) (150 au) away. When the Alfv´en wave reaches thispoint, the rotating mass tied up by the magnetic field is doubled so the rotation velocity becomesa half assuming angular momentum conservation. This timescale is calculated to be ∼ ∼ − M (cid:12) year − (e.g., Shu 1977; Dunham et al. 2014), the magnetic braking effect, which slows down 0.05 M (cid:12) in1700 years, dominates the system. Taken at face value, the estimated field strength is high enoughfor the magnetic field to brake the disk rotation efficiently. However, as we mentioned earlier, thepolarization orientations on the several tens au scale are indicative of a disk. If true, the existence ofa relatively large disk in the presence of a strong inferred magnetic field would point to a decouplingof the field from the bulk disk material, most likely through non-ideal MHD effects, which becomemore important at higher densities (e.g., Inutsuka et al. 2010; Krasnopolsky et al. 2011; Dapp et al.2012; Tomida et al. 2015; Tsukamoto et al. 2015; Zhao et al. 2018).2 Kwon et al.
Figure 5.
Intensity map in color scales. Magnetic fields of the original angular resolution are marked ingray and fields smoothed by nine times larger beams in area are in white. Refer to the text for the thick graycontour lines. In the right is the histogram of the magnetic field directions with respect to the backgroundfield. 5.
CONCLUSIONWe have detected a well-ordered polarization pattern toward the Class 0 YSO L1448 IRS 2, whoseinferred magnetic field presents the clearest hourglass morphology to date on 100 au scales: poloidalin the outer regions and rapidly switching to toroidal in the inner region. This can be interpretedas a toroidal magnetic field wrapped up by a rotating (disk) structure or by a self-scattered polar-ization pattern due to large grains in an inclined disk: either case supports a rotationally dominantstructure. Future high resolution molecular line observations are needed to investigate whether thereis a rotationally supported disk.We found four regimes with different slopes in the relationship between polarization fractions andintensities, which provide interesting constraints on grain alignment mechanisms. In addition, wedetected a clumpy depolarization strip, which is indicative of magnetically channelled protostellaraccretion flows which drag the field lines into a radially pinched configuration that, when combinedwith inclination effects, lowers the degree of polarization.Finally, we estimated the plane-of-sky magnetic field strength using the DCF technique and foundthat magnetic braking should be very efficient in the system, which is inconsistent with the stronghints of a central disk, protobinary, and observations of rotation. Therefore, the magnetic brakingcatastrophe based on simple ideal MHD simulations may not be so disastrous, at least in this source.Our observations emphasize that non-ideal MHD effects (and possibly turbulence) should be takeninto account, in order to fully understand the formation of disks at the early protostellar systems,We are grateful to ALMA staff for their dedicated work and an anonymous referee for helpfulcomments. W.K. thanks Thiem Hoang for fruitful discussions on grain alignments and polarizationmechanisms. W.K. was supported by Basic Science Research Program through the National ResearchFoundation of Korea (NRF-2016R1C1B2013642). Z.-Y.L. is supported in part by NSF AST-1313083and AST-1716259 and NASA NNX14AB38G and 80NSSC18K1095. This paper makes use of the fol-
ASTEX Magnetic fields in L1448 IRS 2
Facilities:
ALMA REFERENCES