BIMA N₂H⁺ 1–0 MAPPING OBSERVATIONS OF L183: FRAGMENTATION AND SPIN-UP IN A COLLAPSING, MAGNETIZED, ROTATING, PRESTELLAR CORE

The star formation process begins within molecular clouds, where the internal pressures (magnetic, turbulent, thermal, or otherwise), supporting dense regions against collapse, lose out to gravity. The densest regions are known as prestellar cores (Ward-Thompson et al. 1994), which are the gravitationally bound cores of clouds that are believed to be undergoing star formation, but as yet show no signature of an embedded protostellar source (Ward-Thompson et al. 2007). Recent studies have revealed the presence of extremely faint protostellar objects inside a small number of cores that were previously believed to be prestellar (Crapsi et al. 2005b; Bourke et al. 2006). These appear to be very young protostars embedded deep within these cores.

However, the manner of the collapse from a prestellar core to a core that contains a protostar is still a matter of some debate; for a review, see Ward-Thompson (2002). Gravitational collapse, rotation, turbulent motions, and fragmentation may all play a role. Detailed study of the densest prestellar cores that do not contain a protostar is necessary to understand the evolution of prestellar cores (Jessop & Ward-Thompson 2001).

The centrally condensed, high column density prestellar core L183, which is also known as L134N, is part of a loose group of high latitude (b ∼ 37°) dense molecular cores associated with the surface of the local bubble at a distance of 110 ± 10 pc (Franco 1989). It has been extensively studied with a number of different spectroscopic species (Lee et al. 2001; Dickens et al. 2000; Pagani et al. 2005; Requena-Torres et al. 2007; Pagani et al. 2007) and in the far-infrared and submillimeter continuum (Juvela et al. 2002; Ward-Thompson et al. 2002; Lehtinen et al. 2003; Pagani et al. 2003; Kirk et al. 2005; Kauffmann et al. 2008). It also appears to be more chemically evolved than other prestellar cores (Dickens et al. 2000) and displays a number of signatures common to cores that are believed to be close to forming a protostar (Lee et al. 2001; Crapsi et al. 2005a). Thus, L183 is a good candidate in which to study prestellar evolution.

In this paper we present new N₂H⁺ observations performed with the Berkeley–Illinois–Maryland Array (BIMA) interferometer. N₂H⁺ is one of the preferred tracers for cold dense molecular material as it is well correlated with submillimeter dust emission and is less prone to gas phase depletion than other molecular species within prestellar cores (Jessop & Ward-Thompson 2001; Bergin et al. 2002; Tafalla et al. 2004), and it has been widely used to survey prestellar cores in general and L183 in particular (Dickens et al. 2000; Caselli et al. 2002; Crapsi et al. 2005a; Pagani et al. 2007).

In Section 2 we describe the configuration of the BIMA array and in Section 3 we show channel and integrated maps of N₂H⁺ intensity. In Section 4 we show the results of fitting various line parameters, and we discuss their implications in Section 5.

BIMA³ was used to map the 3.22 mm (93 GHz) J = 1 → 0 transition of N₂H⁺ (Caselli et al. 1995; Daniel et al. 2006) toward the peak of the L183 submillimeter dust emission.⁴ BIMA was a 10-element millimeter-wave interferometer operated at the Hat Creek Radio Observatory in Northern California (Welch et al. 1996).

BIMA was used in its compact C-array (∼ 2–27 kλ) and ultracompact D-array (∼ 2–9 kλ) configurations. Table 1 shows the individual observation dates, durations, and configurations. A total duration of 50 hr (∼ 6 tracks) was spent in the C-Array configuration and a total duration of 18 hr (∼ 2 tracks) was spent in the D-Array configuration.

Our observations are only sensitive to angular scales between the FWHM of the primary beam (given by the effective diameter of an individual dish) and that of the synthesized beam (given by the uv sampling and weighting). The individual antennae have a physical diameter of 6.2 m and an effective primary beam FWHM of 1.7 arcmin at 3.22 mm. The synthesized-beam FWHM for the combined observations was 13 × 10 arcsec with a semimajor axis position angle of −14°. The beam was elongated due to the declination of the source. Therefore, these observations are sensitive to structure between 10 and 100 arcsec. This corresponds to a spatial scale of 0.005–0.05 pc at 110 pc.

The BIMA correlator was set to mode 5 to give 512 channels across a bandwidth of 6.25 MHz. The channel width was 12.2 kHz (0.039 km s⁻¹). The N₂H⁺ lineset was placed in the first sideband while the 96.4 MHz C³⁴S 2-1 transition was in the opposite sideband, but no signal was detected in the second sideband. Phase and amplitude calibration were conducted against the radio source 1549+026. The assumed flux of 1549+026 at our wavelength was 1.5 Jy (BIMA Calibrator List⁵). The data were reduced using the MIRIAD software package and the visibilities were Fourier transformed into the image plane using natural weighting in order to maximize the sensitivity (Sault et al. 1995). The theoretical per channel rms was 0.121 K.

The total intensity of emission from the seven hyperfine components of N₂H⁺ toward the L183 core is shown in Figure 1, integrated over velocity. The first solid contour is at 2.46 Jy beam⁻¹, which is also the separation between subsequent contours. For this data set the scaling between Jy beam⁻¹ and kelvin is 1.07 K (Jy beam⁻¹)⁻¹.

Zoom In Zoom Out Reset image size

The integrated emission shown in Figure 1 displays the characteristic north–south L183 filament as seen in submillimeter dust emission, infra-red extinction, and single dish line maps of N₂H⁺ 1–0 (e.g., Pagani et al. 2003; Kirk et al. 2005; Crapsi et al. 2005a). The integrated N₂H⁺ intensity shows three bright peaks (north, south, and west). For reference we will refer to these features as L183 N₂H⁺ (L183-N) 1, 2, and 3 in order of their peak intensity. Their positions are listed in Columns 2 and 3 of Table 2 and are shown in Figure 1 and in subsequent figures as labeled open stars.

Table 2. Table of Positions and Hyperfine Fit Parameters Toward the Three Peaks in the Integrated N₂H⁺ Map (Figure 1)

Source	R.A. (2000)	Decl. (2000)	Vlsr (km s−1)	ΔV (km s−1)	τ	TMB (K)	ΔVNT (km s−1)	Vrms [km s−1]	Tex (K)	$\mathcal {G}$ (km s−1 pc−1)	PA$_{\mathcal {G}}$ (°)
L183-N1	15h54m09.5s	−02°52'15''	2.256 ± 0.003	0.123 ± 0.006	0.37 ± 0.08	1.2 ± 0.4	0.064	0.054	6.5 ± 2.0	6.5 ± 5.8	−74 ± 23
L183-N2	15h54m09.1s	−02°52'50''	2.415 ± 0.004	0.127 ± 0.008	2.1 ± 0.6	0.5 ± 0.2	0.071	0.025	3.3 ± 1.8	1.0 ± 2.4	—
L183-N3	15h54m07.3s	−02°52'31''	2.588 ± 0.005	0.108 ± 0.010	0.36 ± 0.18	0.63 ± 0.45	0.025	0.090	4.8 ± 3.8	8.9 ± 8.0	−57 ± 60

Notes. Column 1 lists the source name. Columns 2 and 3 list the central position of each source. Columns 4–7 list the CLASS hfs parameters—radial velocity (V_lsr), linewidth (ΔV), total optical depth summed across all seven hyperfine components (τ), and the main beam temperature (T_MB)—for a fit to the spectra at the positions listed in Columns 2 and 3. Column 8 lists the nonthermal linewidth (ΔV_NT) based on the assumption that T_K = 7 K (see Section 4.2). Column 9 lists the rms of the radial velocity (V_rms) computed in a 5 by 5 pixel box. Column 10 lists the excitation temperature (T_ex). Columns 11 and 12 list the magnitude and direction, and their standard deviations, of the local velocity gradient (see Section 5).

Download table as:  ASCIITypeset image

To examine the kinematic structure of L183 we extracted channel maps of N₂H⁺ intensity across the 101–012 component as 0.039 km s⁻¹ slices between 2.15 and 2.72 km s⁻¹. This component is shown as it is isolated in velocity from the other components and thus is free from blending between components. These channel maps are shown in Figure 2. As the velocity increases, the emission first appears in the northeast, before moving to the south, and finishing in the northwest. These components roughly correspond to the three emission peaks seen in Figure 1. Thus, L183-N1 and L183-N3 are separated spatially and kinematically from each other, but are both connected kinematically to L183-N2.

Zoom In Zoom Out Reset image size

The central section of the SCUBA 850 μm map from Kirk et al. (2005) is shown as a gray-scale image in Figure 3. The SCUBA continuum emission peaks close to L183-N2 and has a northward extension between L183-N1 and 3. This extension is seen to continue on a larger scale in the SCUBA image (Kirk et al. 2005). The L183-N2 peak is not exactly coincident with the SCUBA 850-μm emission peak. The peak in the continuum emission is in fact coincident with the N₂H⁺ emission at 2.44–2.48 km s⁻¹.

Zoom In Zoom Out Reset image size

The data shown in Figure 1 use a pixel scale of 1.87 arcsec per pixel in order to show a smooth intensity distribution. For the analysis of the spectral lines the data were spatially averaged onto a 5.6 arcsec pixel grid in order to produce pixels that are approximately Nyquist sampled and to generate a factor of 3 increase in signal-to-noise ratio. The spectrum for each pixel that had a 2σ or greater feature at ∼ 2.5 km s⁻¹ (the approximate location of the 101–012 hyperfine component) was then extracted and piped to the hyperfine structure (hfs) routine from the CLASS data reduction package⁶ (Hily-Blant 2000; Buisson et al. 2005). The hfs routine was used to simultaneously fit all seven hyperfine components giving best fits for the radial velocity V_lsr, the linewidth ΔV, the total optical depth τ, and the product of the main beam temperature and optical depth (T_MB × τ). For the fitting it was assumed that all hyperfine components had the same linewidth and excitation temperature.

Table 2 lists the fit parameters at the positions of the three peaks identified from the map of integrated emission. The spectra for these positions are shown in Figure 4. We note that the errors on the fitting of V_lsr and Δ_V are narrower than the channel width (0.039 km s⁻¹) due to multiple channels over the line that contribute to the fitting.

Zoom In Zoom Out Reset image size

4.1. Kinematics

Figure 5 shows the variation of the local standard of rest velocity (V_lsr) across the L183 molecular core. It shows the same triple-lobed structure as shown in Figure 2, but now we see that the most extreme velocities are toward the outside of each of the three lobes. The velocities at the position of each of the three L183-N sources are listed in Column 4 of Table 2. There is a 0.33 km s⁻¹ separation between L183-N1 and L183-N3.

Zoom In Zoom Out Reset image size

There have been previous observations of N₂H⁺ toward L183 using the single-dish IRAM 30 meter telescope. Its beamwidth at 3.22 mm is ∼ 30 arcsec. Pagani et al. (2007) measured a value of V_lsr = 2.367 km s⁻¹ toward R.A.(2000) = 15^h54^m085, decl.(2000) = −02°52'48''. This is in good agreement with the values we measure in that area. Crapsi et al. (2005a) measured a value of V_lsr = 2.413 km s⁻¹ toward R.A.(2000) = 15^h54^m084, decl.(2000) = −02°52'23''. This again is similar to the values we measure exactly at that point. The literature value for the V_lsr of the whole core of L183 (2.415 km s⁻¹) matches closely that of L183-N2. Hence we take 2.415 km s⁻¹ as the systemic velocity of the L183 core and we see that the velocities of L183-N1 and L183-N3 are velocity structures overlaid on this, of up to ±0.17 km s⁻¹.

To examine the velocity gradient across L183 we averaged the gradient vectors between a single pixel and its nearest neighbors with the condition that each pixel must have at least seven neighbors with valid measurements of V_lsr. A vector map of the local velocity gradient is overlaid on the map of V_lsr in Figure 5.

Figure 5 also shows a black-line indicating the direction of the plane-of-the-sky magnetic field as observed by Crutcher et al. (2004). This is approximately orthogonal to the large-scale velocity gradient (see Section 5). The magnetic field may therefore play an important role in the evolution of this prestellar core (cf. Holland et al. 1996). We discuss further the relevance of the velocity gradient to previous observations in Section 5.

4.2. Linewidths

Figure 6 shows the remainder of the results from the hfs fitting. Panel (a) shows the linewidth variation across L183. The linewidths measured from the hfs fitting are narrow. The mean linewidth in Table 2 is only 0.119 km s⁻¹ and this is almost within the 1σ dispersion of each measurement. The literature values for N₂H⁺ linewidths in L183 on peak are in the range 0.195–0.25 km s⁻¹ (Caselli et al. 2002; Crapsi et al. 2005a; Pagani et al. 2007), which in turn are on the narrow side of the average linewidth of 0.250 ± 0.07 km s⁻¹ for the 31 starless cores surveyed by Crapsi et al. (2005a).

Zoom In Zoom Out Reset image size

A ridge of broader linewidths of about 0.14 km s⁻¹ occurs between L183-N2 and the northern part of the core containing L183-N1 and 3. The broadest linewidth along this ridge occurs between L183-N1 and L183-N3. This structure can be attributed to blending between the separate velocities of L183-N1 and 3.

A typical observed spectral line can be decomposed into nonthermal ΔV_NT and thermal ΔV_T components where the total linewidth is ΔV² = ΔV²_NT + ΔV²_T. The thermal line width is given by $\Delta V_{T}=c_{s} \sqrt{8\ln 2}$. The sound speed is given by c²_s = kT_K/μm_H where T_K is the kinetic temperature, m_H is the mass of a hydrogen atom, k is the Boltzmann constant, and μ is the mean particle mass in units of atomic mass number (for N₂H⁺ this is 29).

From the above we can estimate the nonthermal component from knowledge of the local sound speed, but we first need an estimate of the kinetic temperature. We have previously estimated the dust temperature in L183, using a graybody fit to the extended infrared and submillimeter emission (in a 150 arcsec diameter aperture) centered on the dust peak, and found a value of 10 K (Ward-Thompson et al. 2002). If this were also the kinetic temperature of the gas, then the thermal linewidth would be 0.126 km s⁻¹, and the nonthermal linewidth would be zero across the core.

However, radiative transfer studies have shown that the extinction within prestellar cores can support very low temperatures at the center (as low as ∼7 K) with the temperature increasing radially outward (Evans et al. 2001; Zucconi et al. 2001; Stamatellos & Whitworth 2003). In this scenario the dust observations give a mean temperature, but the temperature in the densest central parts, as probed by the current observations, is usually lower (Stamatellos & Whitworth 2003).

Another way to calculate the temperature is to use the relative strength of different molecular transitions. Daniel et al. (2007) and Pagani et al. (2007) recently did this for higher quantum number transitions of N₂H⁺ and its deuterated isotopologue. The Pagani et al. (2007) model found an internal temperature of 7 K. The Daniel et al. (2007) model found an internal temperature of 8K within a central radius of 20 arcsec. So we adopt the kinetic temperature of T_K = 7 K. From this we calculate a thermal linewidth of ΔV_T = 0.105 km s⁻¹ and the nonthermal linewidths listed in Column 8 of Table 2. A kinetic temperature of T_K = 8 K would give a ΔV_T = 0.112 km s⁻¹.

Alternatively, we can calculate V_rms, the plane of the sky variation of V_lsr, under the assumption that it is produced by ΔV_NT. This will be a measure of ΔV_NT, but will not necessarily be equal to it. We calculate V_rms by calculating the root-mean-square of the radial velocity in a 5 × 5 pixel box centered on each of the three peaks (∼30 × 30 arcsec). These values are listed in Column 9 of Table 2.

The value of V_rms for L183-N2 is lower than for the other two sources. This might be expected from a study of Figure 5, in which it can be seen that the values of V_lsr appear fairly uniform across L183-N2. In the case of L183-N1 the mean values of ΔV_NT and V_rms only differ by a few ms⁻¹. The variation is slightly greater for L183-N2 and 3, but for all three sources the mean of ΔV_NT and V_rms is about 0.05–0.06 kms⁻¹.

Particularly narrow linewidths are measured along the southern edge of the L183-N2 core. To the south of L183-N2 the total linewidth drops to 81 ± 6 m s⁻¹. If this were purely thermal it would imply a temperature of T_K = 4 K. However, we note that this is the region of highest optical depth (see next section), and so we may not be sampling the whole cloud along this line of sight.

4.3. Optical Depths

4.4. Temperatures

The MIRIAD data set was calibrated in Jy beam⁻¹. For the CLASS hfs fitting a conversion factor of 1.07 Jy beam⁻¹ K⁻¹ was used to covert the extracted spectra into units of main beam temperature T_MB. Figure 6(c) shows the distribution of fitted T_MB across the core. It is notable that the pattern of T_MB does correlate exactly with that of the integrated intensity. The intensity valley that separates L183-N3 from L183-N1 and 2 is more pronounced. This corresponds to the northwestern edge of the emission in the 2.44 km s⁻¹ channel map and the southeastern edge of emission in the 2.60 km s⁻¹ channel map.

The CLASS hfs routine returns the parameter χ = T_MB × τ. Thus the quoted T_MB error contains error contributions from the fit of χ and τ. These uncertainties also reflect the difficulty in fitting the height of the line when there are only a few channels across the peak. We can calculate the excitation temperature from the equation

$$\begin{equation} T_{MB} = \frac{\chi }{\tau } = (T_{\rm ex} - T_{\rm bg}) (1 - e^{-\tau }) \end{equation} \tag{ 1 }$$

where T_bg is the background temperature and is assumed to be equal to the cosmic microwave background (2.73 K).

Figure 6(d) shows a map of T_ex. The values for T_MB and T_ex at the L183-N peaks are listed in Columns 7 and 10 of Table 2. Both temperatures show a distinct increase in the region of L183-N1 where the excitation temperature is approximately equal to our assumed kinetic temperature. The CLASS hfs fitting assumes an equal excitation temperature for all hyperfine components, however, radiative transfer modeling has shown that the temperatures of the separate components can vary by up to 15% (Caselli et al. 1995; Pagani et al. 2007; Daniel et al. 2007). This effect is negligible for our purposes as we are primarily interested in the kinematics and we are in a low optical depth regime.

4.5. Velocity Gradients

4.6. Comparison with Single Dish Observations

The left-hand panel of Figure 7 shows the 5.6 arcsec pixel data smoothed spatially to the resolution of a 30 m single antenna telescope. Our BIMA FWHM at 3.22 mm was 13 × 10 arcsec, the FWHM of a 30 m telescope (e.g., the IRAM dish) at the same wavelength is 27 arcsec. We use the term "30m-Beam" to note when we are referring to the smoothed equivalent of a 30 m single-dish beam. The smoothed map was multiplied by a factor a 5.6, the ratio of the 30m-Beam to BIMA beam integrals, in order to give a map calibrated in Jy/30m-Beam. The conversion factor between Jy/30m-Beam and K becomes 0.191 K/(Jy/30m-Beam). The smoothed map shows that the multiple peaks we detect with the BIMA interferometer would not have been resolved in integrated emission if the observations had been conducted with a 30 m diameter single dish telescope.

Zoom In Zoom Out Reset image size

While the L183-N1, 2, and 3 substructure may not have been seen in integrated intensity maps, it could theoretically have been resolved in channel maps. However, the Pagani et al. (2007) observations were along an west-east slice through L183-N2 and do not comprise a fully sampled map. Furthermore, the Dickens et al. (2000) and Caselli et al. (2002) N₂H⁺ observations of L183 were taken with the FCRAO which has an even lower resolution (50 arcsec) beam at this wavelength. Nevertheless, Figure 2 of Dickens et al. (2000) shows an offset in the distribution of N₂H⁺ in the 2.4–2.6 km s⁻¹ and 2.6–2.8 km s⁻¹ channels that is comparable to us seeing higher velocity material to the west of the core. Caselli et al. (2002) show a comparable vector diagram, albeit over a larger area, but they do not show channel maps.

The (12'', 0'') position from Figure 1 of Pagani et al. (2007) (15^h54^m093 −02°52'48'') is within the same 5.6 arcsec pixel as our L183-N2 position. We repeated the hfs fitting for the smoothed data at that location. The results are shown in the right-hand panel of Figure 7. The fitted V_lsr for the smoothed data is 2.410 ± 0.002 km s⁻¹ which is comparable to that fitted to the unsmoothed data.

The linewidth Δ_V becomes 0.111 ± 0.004 km s⁻¹ which is narrower than the value fitted to the unsmoothed data. L183-N2 is at the systemic velocity of L183, is extended spatially, and shows no velocity structure. Therefore spatially smoothing the wings of the emission (e.g., channels 2.29 and 2.52 km s⁻¹ in Figure 2) would move flux out of the beam whereas smoothing the extended emission in channels near the line center (e.g., channels 2.40 and 2.44 km s⁻¹ in Figure 2) would bring more flux into the beam. This differential effect will naturally cause the linewidth to narrow. The opacity becomes τ = 3.2 ± 0.6, which is higher than for the unsmoothed data, but is still substantially lower than reported for single-dish observations.

The line temperature for our smoothed L183-N2 position was fitted as T_MB = 0.4 ± 0.1 K which agrees with our fit to the unsmoothed data at that position. However, inspection of Figure 7 shows that the fit has been unable to recover the full extent of the line temperature. Daniel et al. (2007) and Pagani et al. (2007) respectively published spectra toward L183 with T*_A and T*_R values of 1.5–2.0 K. The (12'', 0'') position of Pagani et al. (2007) had a peak T*_R ∼ 1.8 K. Even allowing for the BIMA main beam efficiency (83%; Lugten 1995) our line temperatures are only about a quarter of those found in the single dish observations. This agrees with the result of our optical depth fitting to the unsmoothed data. A substantial fraction, possibly 75%, of the extended N₂H⁺ emission from L183 is being resolved away by the interferometric process. It is this process that has revealed the clumps buried within the extended L183 core.

We have studied the prestellar core L183 with high-resolution BIMA interferometric observations, and have seen that the core breaks up into three separate clumps, which we have labeled L183-N1, 2, and 3. The clump L183-N2 is roughly coincident with the densest central part of L183. Furthermore, the velocity of L183-N2 is exactly that of the systemic velocity of the whole of L183. We therefore deduce that L183-N2 is at the center of L183, and that the L183 core has formed two further fragments, L183-N1 and 3.

The virial mass for a uniform density, turbulent, rotating molecular cloud threaded by a uniform magnetic field can be written as $M_{\rm vir} = M_{\Delta V}+M_{B}+M_{\mathcal {G}}$ where the three terms on the right-hand side represent the turbulent and thermal virial mass, the magnetic critical mass, and the kinematic mass of the system from Kepler's third law. This can be written as

$$\begin{equation} M_{\rm vir} = 210\Delta V^{2} r + 9.3Br^{2} + 233\mathcal {G}^{2}r^{3} \; M_{\odot } \end{equation} \tag{ 2 }$$

where ΔV is the linewidth in km s⁻¹, r is the radius of the cloud in parsec, B is the magnetic field strength in μG, and $\mathcal {G}$ is the velocity gradient due to rotation in km s⁻¹ pc⁻¹. Here we have used the expression for M_ΔV from MacLaren et al. (1988). Crutcher et al. (2004) used the Chandrasekhar–Fermi method to estimate a field strength of B = 80 μG toward L183. The approximate radius of L183-N1 and L183-N2 is r ∼ 0.01 pc and the radius of L183-N3 is r ∼ 0.005 pc. Based on these values and the parameters listed in Table 2 we estimate M_vir ∼ 0.1 M_☉ for L183-N1 and 2 and M_vir ∼ 0.03 M_☉ for L183-N3.

For a virialized system the virial mass M_vir will be balanced by the actual mass of the fragment M_clump plus a surface term M_P due to a confining medium. The material we have resolved away will act as this surface pressure term. By equating an expression for the energy in a confining isothermal pressure to the standard formula for gravitational potential energy it can be shown that $M_{P} = 316r^{2}\sqrt{n_{4}T_{10}}$, where n₄ is the number density of the confining medium in units of 10⁴ cm⁻² and T₁₀ is the gas temperature in units of 10 K. Using n₄ = 100 from Kirk et al. (2005) and our assumed gas temperature of 7 K we estimate M_P ∼ 0.3 M_☉ for L183-N1 and 2 and M_P ∼ 0.07 M_☉ for L183-N3. In each case the value of M_P is 2–3 times greater than the value of M_vir. Even without estimating the actual mass of these fragments, we can see that they are each pressure confined and will probably each collapse to form a protostar.

Spitzer archival data of L183 was checked for the presence of a young stellar object (YSO). No YSO candidates were identified toward L183 based on the selection criteria (Harvey et al. 2007; Rebull et al. 2007; Harvey et al. 2008) used by the Spitzer c2d (Evans et al. 2003) and Gould Belt (L. Allen et al. 2009, in preparation) legacy surveys. There are no Two Micron All Sky Survey sources within the L183 fields shown in this paper. We therefore confirm that L183 and each of the L183-N clumps are prestellar in nature and do not contain embedded YSOs.

The large-scale structure of L183 is elongated approximately north–south at a position angle of ∼−20°. Crutcher et al. (2004) gave the position angle of the long-axis of the SCUBA emission as ∼−15°. These position angles are consistent with the elongation orientation of L183-N2. We have examined the detailed velocity structure of the core and found that L183-N1 and 3 each have velocities slightly offset from that of L183-N2—one positive and one negative. Furthermore, there is a velocity gradient in the general direction from L183-N1 to L183-N3 of ∼ 6–9 km s⁻¹ pc⁻¹, over a separation of ∼ 0.016 pc (see Table 2). We refer to this hereafter as the small-scale velocity gradient.

The mean direction of the small-scale velocity gradient across the entire core is −60° ± 62°. However, the mean direction for "high-velocity vectors" (defined as those vectors with V_lsr greater than 0.1 km s⁻¹ away from the systemic velocity, i.e., |V_lsr − 2.415|>0.1 km s⁻¹) is −72° ± 2°. The high-velocity vectors have a more coherent direction than vectors with a velocity close to the systemic velocity. The mean direction for the subset of high-velocity vectors in L183-N1 is −68° ± 30° (nine vectors), and for L183-N3 is −79° ± 5° (five vectors). These directions are measured over the same areas used in Section 4.5.

Therefore the axis of rotation for the small-scale gradient is +22° ± 30° for L183-N1 and +11° ± 5° for L183-N3, assuming that the axis is orthogonal to the direction of the velocity gradient. We note that the axis of rotation given by the mean of all 14 high-velocity vectors in L183-N1 and N3 is +18° ± 19°. This is approximately parallel to the direction of the magnetic field across L183 (+28° ± 2°, Crutcher et al. 2004), but is at an angle of ∼33° to the core long axis. This is indicating that perhaps the magnetic field may be playing an important role in the evolution of L183. It also indicates that, for L183, the axis of rotation shares the same offset from the orientation of the cloud as is found for the magnetic fields in other prestellar cores (Crutcher et al. 2004; Kirk et al. 2006; Ward-Thompson et al. 2009).

We therefore interpret our data of L183 as a rotating elongated filamentary core, whose rotation axis roughly coincides with its magnetic field direction. L183-N2 is on-axis, and therefore essentially stationary. L183-N1 and 3 lie either side of the rotation axis and have similar magnitude line-of-sight velocities in opposite directions. The small-scale velocity gradient therefore tells us about the rotation near the center of L183.

On the larger scale of the L183 prestellar core as a whole, previous single-dish work by Crapsi et al. (2005a) has found a core 0.09 × 0.05 pc in extent; see their Figure 15. These authors also saw a velocity gradient roughly parallel to the gradient we observed, but on a larger spatial scale. They see a gradient of ∼ 1.4 km s⁻¹ pc⁻¹, over a scale of ∼ 0.05 pc. We refer to this hereafter as the large-scale velocity gradient.

The ratio of the spatial scales of the large-scale and small-scale velocity gradients is 0.05/0.016 ∼ 3. The ratio of the velocity gradients themselves is (6–9)/1.4 ∼ 4–6. Given the uncertainties in all of the foregoing estimates and calculations, the agreement of these two ratios to within a factor of two is very interesting. This is exactly what would be predicted by the law of conservation of angular momentum if L183 were collapsing and spinning up as it collapses.

We therefore hypothesize that all of our data of L183 can be explained in terms of a collapsing, fragmenting elongated core, that is rotating about the magnetic field direction and spinning up as it collapses. We note finally that our hypothesis is consistent with the single-dish data of both Crapsi et al. (2005a) and Pagani et al. (2007), whose large-scale velocity vectors also appear to increase in magnitude away from the center of L183. Thus our hypothesis is also consistent with previous data of L183.

We have used BIMA to make deep N₂H⁺ (1–0) maps of the prestellar core L183. The interferometric process resolved away the majority of the extended N₂H⁺ emission and revealed three clumps which we have labeled L183-N1, 2 & 3. L183-N2 is at the established systemic velocity of L183 and is coincident with the position of the peak of the dust emission. We detect a velocity gradient across the core. The direction of the gradient is consistent with previously published lower resolution studies, but its magnitude is larger than previously seen. We interpret this pattern as the spin-up and fragmentation of a rotating and collapsing prestellar core.

We thank the anonymous referee for providing constructive comments that have improved the contents of this paper. Work for this paper was conducted in part with NSF post-doctoral support at the University of Illinois from grant NSF AST 02-28953, and in part with the support of the UK Science and Technology Facilities Council (STFC) via the Cardiff Astronomy Rolling Grant. BIMA was operated by the Berkeley–Illinois–Maryland Association with NSF and University Funding. The BIMA facilities have since been merged with the OVRO array to form the present CARMA array supported at the University of Illinois by University funding and by grant NSF AST 05-40459.

Facilities: BIMA ()