DISTRIBUTIONS OF CARBON-CHAIN MOLECULES IN L1527

Observations of IRAS 04368+2557 in L1527 (d = 137 pc; Torres et al. 2007) were carried out with the PdBI on 2007 August 27 and 2008 April 16 in the D and C configurations of the array, respectively. The field center is (α₂₀₀₀, δ₂₀₀₀) = (04^h39^m5389,  26°03'110), which is the IRAS position of L1527. The observed lines are listed in Table 1. We used the 3 mm SIS mixer receivers, whose typical system temperature was about 80–120 K. For the line observation, a narrowband correlator was used, where six windows with 20 MHz bandwidth each were assigned to the molecular lines. The velocity resolution is 0.14 km s⁻¹ at 85 GHz. For the continuum observation, we also used a narrowband correlator, where two windows with 320 MHz bandwidth were assigned. Phase and amplitude calibrations were obtained by observing 0507+179 or 0415+379 every 25 minutes. The bandpass calibration was carried out on 0507+179 and 0415+379, and the absolute flux density scale was derived from 3C454.3. The data calibration was performed in the antenna-based manner, and uncertainties are less than 10%. The continuum image was produced by averaging line-free channels. The line maps were obtained by cleaning line images after subtracting the continuum directly from the visibilities. The primary beam (half-power beam width, HPBW) is 57'' at 85 GHz.

To obtain the short-spacing data, single-dish observations were carried out toward L1527 with the IRAM 30 m telescope at Pico Veleta in 2008 September. The lines listed in Table 1 were observed. The 3 mm SIS receivers (A/B100) were used as front ends, whose system noise temperatures ranged from 80 to 160 K. The beam size of the telescope is 29'' at 85.6 GHz. The main beam efficiency is 0.82. The telescope pointing was checked every hour by observing nearby continuum sources and was maintained to be better than 4''. The backend was an autocorrelator, VESPA. We set the individual bandwidth and resolution to be 20 MHz and 20 kHz, respectively. The frequency resolution corresponds to a velocity resolution of 0.07 km s⁻¹ at 85.6 GHz. The observations were made in the frequency switching mode with a frequency offset of 3.57 MHz. Since the line width is much narrower (∼0.5 km s⁻¹) than the baseline ripple (∼80 km s⁻¹), we readily subtracted the baseline for the 20 km s⁻¹ span with the second- or third-order polynomial. We made the 5 × 5 map around the IRAS position (field center) with the grid spacing of 141. Note that the short-spacing data of HC₅N was obtained but with insufficient signal-to-noise ratios and were therefore not considered any further.

Figure 1 shows the 3.5 mm continuum emission map observed with the PdBI, where the 1σ noise level is 0.19 mJy beam⁻¹. The half-power synthesized beam size and the position angle (as measured from east of north) of the beam are 376 × 357 and 26°, respectively, where the natural weighting is employed. We detected a continuum emission from the protostar with a high signal-to-noise ratio. The peak flux is 22 mJy beam⁻¹. The peak position is determined to be (α₂₀₀₀, δ₂₀₀₀) = (04^h39^m5387,  26°03'097). It is shifted from the field center by Δα = −03, Δδ = −13. The peak position is consistent with previous reports (e.g., Maury et al. 2010).

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The continuum emission is hardly resolved in the present observation. However, we can see a slight elongation along the southwest–northeast direction in the two lowest contours in Figure 1. This direction is almost parallel to the elongation of the infalling envelope proposed by Ohashi et al. (1997). Such an elongation is also seen in the distribution of c-C₃H₂, as described later (Section 3.3).

According to our single-dish observations with the Nobeyama 45 m telescope (Sakai et al. 2008a), carbon-chain molecules are distributed over the 40'' scale around the protostar. Hence, the short-spacing data are indispensable to recover the resolved-out components in the interferometric observation. We therefore combined the single-dish data taken with the IRAM 30 m telescope covering the 71'' × 71'' area to the interferometric data except for the HC₅N (J = 32–31) line.

Figure 2 shows the velocity channel maps of the CCH (N = 1–0,  J = 3/2–1/2,  F = 1–1) line, which is one of the weakest hyperfine components of the N = 1–0 transition. We chose this line in order to avoid saturation effects seen for the stronger components in the single-dish observation (Sakai et al. 2008a, 2009a, 2010). Nevertheless, the line was detected with more than 5σ confidence level over the velocity range of 1.5 km s⁻¹. Around the systemic velocity of the L1527 core (5.9 km s⁻¹), the emission is extended over the size of 30''–40'', as already reported by Sakai et al. (2008a). The size corresponds to the radius of 2000–3000 AU from the protostar. However, the emitting region becomes compact around the protostar position for the blueshifted and redshifted components. The 6.65 km s⁻¹ component is distributed toward north of the protostar position, while the 5.44 km s⁻¹ component is toward south of the protostar position. The blueshifted and redshifted components are distributed within a radius of 500–1000 AU from the protostar, where the redshifted component is stronger and more compact than the blueshifted one.

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According to the C¹⁸O (J = 1–0) observation with the Nobeyama Millimeter Array by Ohashi et al. (1997), an infalling envelope with a radius of 2000 AU exists around the protostar, whose rotation axis is almost along the east–west direction. Ohashi et al. (1997) reported that the north part of the envelope is redshifted, while the south part is blueshifted. This trend is consistent with the above result for CCH; hence, the CCH distribution traces a part of the infalling envelope seen in the C¹⁸O distribution. Note that the systemic velocity of the infalling envelope is reported to be 5.6–5.7 km s⁻¹ by Ohashi et al. (1997), which is slightly different from that of the L1527 core (5.9 km s⁻¹). The difference will be discussed later (Section 3.3).

Figure 3 shows the velocity channel maps of the C₄H (N = 9–8,  F₁) line. As in the case of CCH, the emission of C₄H was detected with more than 5σ confidence level over the velocity range of 1.5 km s⁻¹. Judging from the spatial extent of the distribution, the systemic velocity of the core is around 5.9 km s⁻¹ for C₄H, which is consistent with the CCH result. The emitting region becomes compact around the protostar position for the blueshifted and redshifted components. In contrast to CCH, the blueshifted component is stronger and more compact in C₄H than the redshifted component. The origin of the difference between CCH and C₄H is puzzling. Nevertheless, the result indicates that C₄H certainly exists even within the radius of 500–1000 AU from the protostar. In the single-dish observation with the Nobeyama 45 m telescope, the line width is found to be significantly broader toward the protostar than toward the surrounding positions, which is interpreted by the infalling motion (Sakai et al. 2008a). The association of C₄H with the protostar is now confirmed directly by resolving the spatial distribution.

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Figure 4 shows the velocity channel maps of the c-C₃H₂ (4_3,2–4_2,3) line. The c-C₃H₂ molecule is not a carbon-chain molecule in a strict sense. However, it can be regarded as a related molecule, because its production is related to carbon-chain molecules (e.g., Thaddeus et al. 1985). Again, the emission was seen with more than 5σ confidence level over the velocity range of 1.4 km s⁻¹. The size of the distribution is 1000–2000 AU in radius from the protostar at the velocity channel around 5.9 km s⁻¹, smaller than those of CCH and C₄H. The compact distribution can be seen both in the blueshifted and redshifted components, as in the case of CCH and C₄H. The upper state energy of this line is 29 K, and the critical density is about 3 × 10⁶ cm⁻³. Therefore, this line preferentially traces a denser part of the infalling envelope than the CCH and C₄H lines, which have lower critical densities (∼10⁵ cm⁻³) because of the small dipole moments. Therefore, the c-C₃H₂ line is the most useful probe among the observed lines to explore the inner part of the rotating infalling envelope, which will be described in Section 4.

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As for the HC₅N(J = 32–31) line, we cannot present meaningful channel maps because of the poor signal-to-noise ratio. Therefore, only the integrated intensity map will be presented later.

Figure 5 shows the integrated intensity maps of the CCH and C₄H lines superposed on a gray-scale image of the 3.5 mm continuum emission. The distributions of CCH and C₄H have a size similar to the X-shaped structure seen in the ¹³CO (J = 1–0) interferometric observation by Ohashi et al. (1997). The X-shaped structure is thought to be the infalling envelope around the protostar, although the fact that this structure is seen at the blueshifted velocity range in ¹³CO is probably due to the resolved-out problem around the systemic velocity. The present result suggests that CCH and C₄H would be distributed over the infalling envelope. On the other hand, it should be noted that their intensities become slightly weaker toward the dust continuum peak (see Section 4.3). This may indicate that the gas-phase abundances of the carbon-chain molecules decrease within a radius of 300–600 AU from the protostar.

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In the present study, the high excitation line of HC₅N (J = 32–31,  E_u = 67 K) was observed. Since the detection is marginal, we only present the integrated intensity map in Figure 6. For this molecule, we did not add the short-spacing data due to a lack of single-dish data with a good signal-to-noise ratio, and the fraction of the emission that is resolved out could not be recovered. This fraction is significant: by comparing the flux observed with the PdBI with the flux from the single-dish observation toward the center position, about 60% of the single-dish flux is found to be missing. Nevertheless, a ring-like structure around the protostar whose size is similar to the FWHM size of the CCH and C₄H distributions can be seen for HC₅N.

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Figure 7 shows the integrated intensity map of the c-C₃H₂ (4_3,2–4_2,3) line. The distribution of this line is more centrally concentrated than those of the CCH and C₄H lines, as indicated by the velocity channel maps. The overall distribution is not similar to the ¹³CO (J = 1–0) distribution, but rather to the C¹⁸O (J = 1–0) distribution (Ohashi et al. 1997). It also resembles the distribution of H¹³CO⁺ (J = 1–0) observed with the Nobeyama Millimeter Array by Saito et al. (2001). The central part of the distribution shows an elliptical shape whose major axis is slightly tilted from the north–south direction. This tilt is similar to that found in the two lowest contours of the 3.5 mm dust continuum map (Figure 1). It is also seen in the velocity channel map of H¹³CO⁺ (Saito et al. 2001). The intensity dip near the protostar position can marginally be recognized.

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Since the c-C₃H₂ (4_3,2–4_2,3) line traces the denser part of the envelope structure better than the CCH and C₄H lines, we here investigate its velocity structure in more detail. We prepare the position–velocity (PV) diagram along the dashed gray lines perpendicular to each other, as shown in Figure 7. We choose the cuts approximately along with the major and minor axes of the elliptical shape structure.

The right panel shows the PV diagram nearly along the north–south direction. The velocity width within 10'' from the protostar is apparently wider than that outside it. When we focus on the lowest two contours, they have a diamond shape. Moreover, a slight velocity gradient can be seen within the 10'' region: the southern part is a bit blueshifted, whereas the northern part is a bit redshifted. According to Ohashi et al. (1997), the PV diagram shows a diamond shape without a velocity gradient, when the protostellar envelope is just infalling without rotation. On the other hand, the PV diagram tends to show the velocity gradient as the contribution of the rotation increases (see Figure 10 of Ohashi et al. 1997). The peak positions are different from the model, probably reflecting the molecular distribution within the envelope, but the overall structure of the PV diagram such as a diamond shape with a velocity gradient seems to be qualitatively consistent with the model of the infalling envelope with rotation by Ohashi et al. (1997).

The top panel of Figure 7 shows the PV diagram nearly along the east–west direction. The velocity width is largest within the central area of 10'', and less outside it. A slight velocity gradient can be seen as in the case of the major axis. The velocity gradient along the minor axis may reflect the detailed geometry of the infalling envelope including the inclination angle. Thus, we can say that c-C₃H₂ exists in the inner part of the infalling envelope (probably a disk component) even within a radius of 500 AU from the protostar.

The systemic velocity of the c-C₃H₂ line is estimated from the PV diagram to be 5.65 km s⁻¹. This is consistent with the systemic velocity of C¹⁸O in the interferometric observation but is shifted by 0.25 km s⁻¹ from the systemic velocities of the L1527 core measured by the single-dish observations (∼5.9 km s⁻¹). Judging from our CCH and C₄H channel maps, the ambient gas obviously has a velocity of 5.9 km s⁻¹. Therefore, the systemic velocity of the infalling envelope seen in c-C₃H₂ and C¹⁸O seems different from the systemic velocity of the parent core. The origin of the difference is still an open question. However, it is comparable to or less than the typical velocity width seen in dense cores (∼0.5 km s⁻¹; Benson & Myers 1989). If the parent core had two or more subcomponents within it like TMC-1 (e.g., Dickens et al. 2001) and the star formation took place in one of them, the difference might be possible.

Since the emission of the c-C₃H₂ (4_3,2–4_2,3) line mainly comes from the inner part of the infalling envelope, its distribution is investigated in more detail. We prepare the integrated intensity profiles along the three cuts passing through the protostar position: south–north, east–west, and northeast–southwest (Figures 8(a)–(c)). In all the three cuts, a steep increase of the intensity can be seen at a radius of about 1000 AU from the protostar. The intensity profiles along the three cuts are similar to one another except for small shoulders appearing in the east–west cut and the south–north cut. The slight difference originates from the asymmetric structure of the real protostellar envelope. Hence, we employ an approximately spherical model of the protostellar envelope reported by Jørgensen et al. (2002), based on the one-dimensional radiative transfer code DUSTY (Ivezić & Elitzur 1997), to interpret the profiles.

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Figures 8(d) and (e) show the variation of the H₂ column density and the gas kinetic temperature, respectively, as a function of the distance from the protostar on the basis of the model for L1527 (Jørgensen et al. 2002). First we compare the temperature distribution with the integrated intensity profile of c-C₃H₂ along the three cuts. As seen in Figure 8(e), the radius at which the intensity shows the steep increase corresponds to the temperature range of 20 to 30 K. In order to assess whether this steep increase originates from the excitation condition or not, we simulate the expected intensity of the 4_3,2–4_2,3 line taking account of the density and temperature structures of the model. For this purpose, we numerically integrate the radiative transfer equation, where the constant fractional abundance of c-C₃H₂ relative to H₂ (2.7 × 10⁻⁹) is employed and the ortho–para ratio is assumed to be 3. The fractional abundance is chosen so as to explain roughly the intensity of the outer envelope (∼2000 AU). Since the critical density of the 4_3,2–4_2,3 line is as high as 3 × 10⁶ cm⁻³, its emission is not thermalized to the gas kinetic temperature in most parts of the cloud. Hence, the excitation temperature and the level population are evaluated by the statistical equilibrium calculation. The detail of the simulation is given in the Appendix. As shown in Figure 9(a), the calculated intensity profile of c-C₃H₂ as a function of the radius does not show the steep increase feature at about 1000 AU that is seen in the observed intensity variation. Therefore, the steep increase of the intensity is not an excitation effect but is due to the abundance jump inside the radius of the steep increase.

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Figures 9(b)–(d) show the simulations for which the abundance of c-C₃H₂ is enhanced by a factor of 10 in the inner region where the temperature is above 20, 25, and 30 K, respectively. Here, the synthesized beam size is taken into account. In this case, the steep increase of the intensity at about 1000 AU is successfully reproduced, and the steep increase position of the intensity approaches the protostar for the higher threshold temperature. On the other hand, the intensity dip toward the protostar position seen in the observed intensity profiles (Figures 8(a)–(c)) is not reproduced in the simple abundance jump model, where the abundance is uniformly enhanced within a certain radius from the protostar. This means that the dip is not the excitation effect. This point will be discussed in Section 4.3.

According to Jørgensen et al. (2002), the power index, α, of the H₂ density distribution (n ∝ r^−α) is 0.6, showing a flat density distribution. It is different from the power indices (1.0–1.6) for the other protostars. However, we confirm that the steep increase around 1000 AU could not be reproduced without the abundance jump, even if the power index were 1.6. Therefore, our conclusion for the abundance jump is robust against the assumption of the H₂ density profile.

It is most likely that the abundance enhancement is a consequence of WCCC. As mentioned briefly in the introduction, carbon-chain molecules and their related molecules are regenerated in the vicinity of the protostar triggered by evaporation of CH₄ from dust grains (Sakai et al. 2008a). The sublimation temperature of CH₄ is around 25 K (Collings et al. 2004), and the steep abundance jump of carbon-chain molecules is expected around this temperature. Since c-C₃H₂ can also be produced from CH₄ by WCCC, it would also show the abundance jump. Our observation indicates that this is the case; hence, it is the confirmation of WCCC on the basis of the spatial distribution of its product.

For hot corino sources, it is known that some organic molecules like H₂CO and CH₃OH show a similar abundance jump within a certain radius from the protostar (e.g., Ceccarelli et al. 2000; Schöier et al. 2002; Maret et al. 2005). This is due to evaporation of mixed ices for T>100 K at the radius of about 150 AU. The radius is much closer to the protostar than that of the abundance jump for carbon-chain molecules in WCCC, which is triggered by evaporation of CH₄ from dust grains at around 25 K with subsequent gas-phase reactions.

We also prepare the intensity profiles of the CCH and C₄H lines along the three cuts (Figure 10). In contrast to the c-C₃H₂ case, the distributions are slightly extended to the outer part, where the temperature is below 20 K. In particular, the intensity profile of the C₄H line along the south–north direction is rather gentle, and a steep increase cannot be recognized very clearly. Since the distributions of these two lines have a size similar to the X-shaped structure reported by Ohashi et al. (1997), the variation of the east–west direction may reflect the structure. Nevertheless, the positions of the two peaks near the center mostly correspond to the temperature of 20–30 K, according to the protostellar envelope model for L1527. Therefore, the regeneration by the WCCC mechanism also seems to contribute to the CCH and C₄H distributions.

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The extended feature of the CCH and C₄H distributions is consistent with the single-dish observation with the Nobeyama 45 m telescope (Sakai et al. 2008a). WCCC is triggered by the evaporation of CH₄ from grain mantles, which was produced by the hydrogenation reactions of the carbon atom on grains. To make this mechanism work, a parent core has to be contracted with a relatively short timescale close to the free fall case, so that the carbon atom can deplete onto dust grains before it is converted to CO in the gas phase (Sakai et al. 2009a, 2009b). In such a fast contraction case, carbon-chain molecules would partly survive even after the onset of star formation. This effect would contribute to the rather extended distributions of CCH and C₄H, which tend to obscure the effect of regeneration by WCCC. Note that this effect is less significant for the c-C₃H₂ line because it traces a denser part than the C₂H and C₄H lines due to the higher critical density. Another possibility for the extended feature is that the extended region would instantaneously be heated up by protostellar activities such as the shock by outflows. In any case, the essential point of the chemistry of L1527 would be a combination between the regeneration by WCCC and the remnant effect of carbon-chain molecules produced in the starless-core phase (Sakai et al. 2009a, 2009b).

When we focus on the vicinity of the protostar position, a small dip can be seen in the integrated intensity maps of all the molecules observed in the present study. The size of the dip is as small as 300–600 AU, which is much smaller than the hole of carbon-chain molecules (a few 1000 AU to 10000 AU) seen toward starless cores in the late evolutionary stage (e.g., Ohashi et al. 1999; Aikawa et al. 2001). In particular, the dip can clearly be seen in all the three cuts for c-C₃H₂ shown in Figures 8(a)–(c). This indicates that the abundance of c-C₃H₂ significantly decreases in the innermost part within a radius of 300–600 AU from the protostar. Note that the dip structure is slightly different from cut to cut, which may reflect the asymmetry of the protostellar envelope.

We examine the central dip with the radiative transfer calculation. Figures 11(a)–(c) show the results where the abundance jump is assumed to occur in the temperature range from 25 to 40 K, from 25 to 35 K, and from 25 to 30 K, respectively. As seen in Figure 11(a), the dip is not seen, if the temperature range for the abundance jump is from 25 to 40 K. Hence, the dip is sensitive to the upper range of the abundance jump. On the other hand, the steep increase position of the intensity is to the lower range as already shown in Figures 9(b)–(d).

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Figure 12 shows the comparison of the south–north and northeast–southwest profiles with the simulations. Since the east–west and south–north profiles of c-C₃H₂ are similar in the depth of the dip and position of the steep increase, only one of the two profiles is used for comparison with the radiative transfer calculation, while the northeast–southwest profile is investigated separately. As shown in Figure 12 (left panel), the south–north profile is well reproduced for the abundance jump from 25 to 30 K, except for a shoulder at the south side probably caused by asymmetry of the distribution. The observed peak intensity can be explained, when the fractional abundance of c-C₃H₂ in the 25–30 K region is 2.7 × 10⁻⁸. When the range of the abundance jump is changed by 1 K, the calculated profile is changed significantly. Hence, the above range is carefully chosen by adjusting the temperature range by 1 K. Figure 12 (right panel) shows the results for the northeast–southwest cut, which has the largest dip among the three cuts. In this case, the best-fit result is obtained for the abundance jump from 22 to 27 K, where the fractional abundance is assumed to be the same as that for the south–north cut. Note that the optical depth is estimated to be less than 0.8 for the case of the abundance jump from 22 to 27 K according to the simulation. Since the line is not optically very thick, it is unlikely that the intensity dip toward the protostar originates from the self-absorption effect. If the dip were due to self-absorption effect, one would expect a redshifted dip and asymmetric peaks due to infall effects in the spectrum toward the continuum peak. The absence of such an absorption feature further supports the fact that the dip is not an optical depth effect (see the top or the right panel of Figure 7).

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In our simulation, we assume a spherical cloud. However, the temperature range for the abundance jump looks different between the south–north cut (and hence east–west cut since these profiles are similar) and the northeast–southwest cut, as seen in the top panels of Figure 12. The south–north and east–west cuts give the highest temperature range among the three cuts, whereas the northeast–southwest cut gives the lowest temperature range. Although these two temperature ranges differ from each other by only a few Kelvin, the calculated profiles are significantly different. We also try to use the intermediate temperature range, 23–29 K, to reproduce both profiles with a single temperature range as much as possible. The result is shown in the bottom panels of Figure 12. The observed profiles are not well reproduced in comparison with the case using the different temperature range for each profile. In particular, the separation of the two peaks around the dip seems inconsistent with the observed profiles. The real cloud structure would be asymmetric, owing to the formation of a protostellar disk or to the disruption of a part of the parent core by outflows. The difference may reflect the asymmetric distribution of the gas around the protostar. Hence, the model should be used under this limitation. Nevertheless, the result that an abundance jump occurs around 20–30 K is consistent with the mechanism of WCCC, because the temperature range is close to the sublimation temperature (25 K) of CH₄ (Collings et al. 2004).

The intensity profiles of CCH and C₄H are also explained by the abundance jump from 20 K to 30 K, although the steep increase and the dip structure are not as clear as the c-C₃H₂ case due to overlapping of the extended component. As in the case of the c-C₃H₂ line, we also simulate the intensity profiles of the CCH and C₄H lines by the radiative transfer calculation based on the protostellar envelope model for L1527 by Jørgensen et al. (2002). Since the critical densities of the CCH and C₄H lines are lower than the cloud density, the local thermodynamic equilibrium (LTE) condition is assumed to calculate the level population. The peak abundance of CCH is thus estimated to be 2 × 10⁻⁷, which is even higher than the corresponding abundance in TMC-1 ((0.5–1) × 10⁻⁷). This further supports the regeneration of CCH. Furthermore, the peak abundance of C₄H is roughly estimated to be (3–4) × 10⁻⁸, which is higher than that derived from the single-dish observation (6 × 10⁻⁹; Sakai et al. 2009a). The difference seems to originate from the beam dilution effect of the single-dish measurement. On the other hand, this abundance is almost comparable to that found in TMC-1 (3 × 10⁻⁸; Irvine et al. 1987).

The above results mean that the abundances of carbon-chain molecules are enhanced in a relatively narrow range of the temperature. The abundances of carbon-chain molecules and their related molecules show immediate increase after the evaporation of CH₄ at its sublimation temperature of about 25 K, as shown in the dynamical model by Aikawa et al. (2008). Once CH₄ is evaporated from grain mantles, it will not be further supplied in the higher temperature region. On the other hand, carbon-chain molecules produced in the gas phase from CH₄ will be destroyed by the gas-phase reactions or will be depleted onto dust grains in the innermost part, although such an abundance decrease is not predicted in the model by Aikawa et al. (2008). In the simulation (Figures 11 and 12), we arbitrarily assume the fractional abundance of c-C₃H₂ for the innermost part to be the same as that for the outer envelope (2.7 × 10⁻⁹). Hence, it is still important to quantitatively determine how abundant carbon-chain molecules remain around the protostar. However, our observation has little sensitivity for this estimation mainly due to insufficient spatial resolution; the size of the dip and the abundance couple with each other. Observations with higher spatial resolution further resolving the dip structure are needed to decouple them.

We consider the following three reasons for deficiency of carbon-chain molecules toward the innermost region. The first possibility is the photodissociation by the radiation of the protostar. Since the protostar in L1527 is low mass and in the class 0 to class I phase, it is not certain whether the photon energy is enough to destroy c-C₃H₂, CCH, and C₄H. Nevertheless, the visual extinction at the 600 AU position from the protostar is estimated to be 10 mag on the basis of the protostellar envelope model for L1527. Hence, photodissociation may work to some extent. Even if the photodissociation process is effective in the innermost part, molecules in the midplane of the protostellar disk will survive, giving the blueshifted and redshifted line components. The second is the gas phase destruction of carbon-chain molecules. Hassel et al. (2008) reported on the basis of their chemical model calculation that carbon-chain molecules are not destroyed even if the temperature goes up to 200 K. However, destruction processes of carbon-chain molecules are not well established, and they might have to be considered more carefully. The third is the depletion onto dust grains. The adsorption timescale is about 10³ yr, while the timescale to pass through the dip to the protostar is estimated to be 2 × 10³ yr. If the gas is in the midplane of the protostellar disk and is centrifugally supported, the timescale to pass through the dip will be much longer. Hence, depletion can occur.

In the above discussions, the central dip is ascribed to the abundance decrease. However, it is not necessary, if the H₂ density of the innermost part is significantly lower than that expected from the density distribution of the outer region. In other words, the central dip might reflect a change in the physical structure near the protostar, for instance, formation of the protostellar disk. It should be noted that the centrifugal radius for the 0.5 km s⁻¹ rotation velocity around the 0.1 M_☉ protostar is 300 AU, which is close to the size of the dip. Furthermore, it is known that IRAS 04368+2557 is resolved into two components separated by 24 AU (017; Loinard et al. 2002). Such a substructure might affect the density distribution.

We should emphasize that CCH, C₄H, and c-C₃H₂ certainly exist in the blueshifted and redshifted components associated with the protostar. Such carbon-chain molecules remaining near the protostar will be brought into the protoplanetary disks through the protostellar disks. In particular, if a portion of these molecules are depleted onto dust grains, they will likely survive on these grains as the protostellar disk evolves toward the protoplanetary disk phase. Understanding of chemical evolution from the protostellar disks to the protoplanetary disks is an important target in the future and detailed characterization of the WCCC sources is one of the indispensable first steps. For this purpose, observations of the innermost part with the high critical density lines at the submillimeter wave region will be essential.