source
stringlengths 1
2.05k
⌀ | target
stringlengths 1
11.7k
|
|---|---|
While the formation and growth of a massive black hole undoubtedly requires a gas-dynamical. dissipative process. our results suggest that an NSC can be assembled nondissipatively. by the collisionless migration of star clusters. and thus. the observed agreement could be a coincidence. | While the formation and growth of a massive black hole undoubtedly requires a gas-dynamical, dissipative process, our results suggest that an NSC can be assembled nondissipatively, by the collisionless migration of star clusters, and thus, the observed agreement could be a coincidence. |
Although the well-studied NSC host galaxy M33 does not contain a central massive black hole (Merrittetal.2001:Gebhardtetal.2001.andreferences therein).. another one. 44395. does (Filippenko&Ho2003).. and still others contain AGNs (Sethetal.2008). | Although the well-studied NSC host galaxy M33 does not contain a central massive black hole \citep[][and references
therein]{Merritt:01b,Gebhardt:01}, another one, 4395, does \citep{Filippenko:03}, and still others contain AGNs \citep{Seth:08}. |
. Since AGNs and the growth of a central black hole require gas inflow into the center of the galaxy. NSC growth from migrating disk clusters would not generally be accompanied. with black hole growth. because disk clusters contribute stellar mass without augmenting the black hole mass (although gas inflow may be enhanced by the migrating clusters. see. e.g.. Goodman&Rafikov2001:Chang2008)). | Since AGNs and the growth of a central black hole require gas inflow into the center of the galaxy, NSC growth from migrating disk clusters would not generally be accompanied with black hole growth, because disk clusters contribute stellar mass without augmenting the black hole mass (although gas inflow may be enhanced by the migrating clusters, see, e.g., \citealt{Goodman:01,Chang:08}) ). |
This suggests that the central black hole mass in bulgeless disks should not be correlated with the mass of the NSC and that any pseudobulge that is present. unless the migrating clusters independently synthesize massive (or intermediate-mass) black holes which. in this case. they would deliver to the center to merge to form à more massive central black hole (Elmegreenetal.2008).. | This suggests that the central black hole mass in bulgeless disks should not be correlated with the mass of the NSC and that any pseudobulge that is present, unless the migrating clusters independently synthesize massive (or intermediate-mass) black holes which, in this case, they would deliver to the center to merge to form a more massive central black hole \citep{Elmegreen:08a}. . |
clensity of more distant clusty galaxies is available. | density of more distant dusty galaxies is available. |
There is an upper limit to the surface density of sources at mmm (Wilner Wright 1997): counts at p/m (Smail et 11997: Bareer ct 1998: LIolland et 11998: Hughes et 11998: Barecr et 119998: Blain et 1999b: Eales et 11999): upper imits (Small et 11997: Darger et 11998). and a new count (Blain et 22000) at pm: 175-jm.LSO counts rom Ixawara et ((1998) and Puget et ((1999): 95-5. counts from Wawara et ((1998): and 7- and 15-5 counts rom an extremely deep.50 image of Abcll22390 CATieri et 11999). which yields counts that are even. deeper han those determined in blank-Bield survevs by Oliver et ((1997). Aussel ct al. ( | There is an upper limit to the surface density of sources at mm (Wilner Wright 1997); counts at $\mu$ m (Smail et 1997; Barger et 1998; Holland et 1998; Hughes et 1998; Barger et 1999a; Blain et 1999b; Eales et 1999); upper limits (Smail et 1997; Barger et 1998), and a new count (Blain et 2000) at $\mu$ m; $\mu$ m counts from Kawara et (1998) and Puget et (1999); $\mu$ m counts from Kawara et (1998); and 7- and $\mu$ m counts from an extremely deep image of 2390 (Altieri et 1999), which yields counts that are even deeper than those determined in blank-field surveys by Oliver et (1997), Aussel et al. ( |
1999) and Flores et (1999). | 1999) and Flores et (1999). |
If the values of the activity parameter at recishift ZOLO. (m),1. listed in Tablell are usec to estimate the counts of galaxies at SSO and yam. then the results unclerpredict the observed counts by a luge factor. | If the values of the activity parameter at redshift zero, $(F\sigma)_0^{-1}$, listed in 1 are used to estimate the counts of galaxies at 850 and $\mu$ m, then the results underpredict the observed counts by a large factor. |
The form of evolution of the merger ellicieney ος) is fixed. by the observed background radiation intensity. and so. keeping within the framework of our well-constrained. models. the value of the activity parameter (£m) at high. redshift must be allowed. to increase above its value at. recishift zero in order to account for the observations. | The form of evolution of the merger efficiency $x(z)$ is fixed by the observed background radiation intensity, and so, keeping within the framework of our well-constrained models, the value of the activity parameter $(F\sigma)^{-1}$ at high redshift must be allowed to increase above its value at redshift zero in order to account for the observations. |
This has the ellect of increasing the luminosity of high-redshift mergers. thus increasing the 175- ancl μι counts. | This has the effect of increasing the luminosity of high-redshift mergers, thus increasing the 175- and $\mu$ m counts. |
However. the background. radiation intensity and the Iow-redshift 60-74 counts remain unchanged. | However, the background radiation intensity and the low-redshift $\mu$ m counts remain unchanged. |
The form of evolution of the activity parameter (£e)1 that is required to explain the data is illustrated in SS. | The form of evolution of the activity parameter $(F\sigma)^{-1}$ that is required to explain the data is illustrated in 8. |
In Fig.sSs(a) the ratio of the model predictions and the observed counts at wavelengths of 175 and m (Ixzwara et al. | In 8(a) the ratio of the model predictions and the observed counts at wavelengths of 175 and $\mu$ m (Kawara et al. |
1998: Blain et 11999b respectively) are compared as a function of the activity parameter in the four mocdels listed in H1. | 1998; Blain et 1999b respectively) are compared as a function of the activity parameter in the four models listed in 1. |
Phe same value of the activity parameter cannot account for the observed counts at both wavelengths simultaneously. and the value required to explain the low-redshift. 60-//mi counts is dillerent. from. either. | The same value of the activity parameter cannot account for the observed counts at both wavelengths simultaneously, and the value required to explain the low-redshift $\mu$ m counts is different from either. |
The value of the activity. parameter required. to. [it the 60-.. 175- and S50-//imi counts increases monotonically. | The value of the activity parameter required to fit the 60-, 175- and $\mu$ m counts increases monotonically. |
Because the median redshift of the galaxies contributing to the counts ab these redshifts is expected to increase monotonically. in SS(b) we present the ratio of the model predictions aud he observed. counts as a function of a parameter ps that describes a simple form of exponential redshift evolution of he activity. parameter. The exponential form provides a reasonable fit to the data. out. is only one example of a whole family of potential unctions. | Because the median redshift of the galaxies contributing to the counts at these redshifts is expected to increase monotonically, in 8(b) we present the ratio of the model predictions and the observed counts as a function of a parameter $p_\sigma$ that describes a simple form of exponential redshift evolution of the activity parameter, The exponential form provides a reasonable fit to the data, but is only one example of a whole family of potential functions. |
Phe important feature is that the function chosen ο represent the activity parameter (£m)+ increases rapiclly with increasing recshift. | The important feature is that the function chosen to represent the activity parameter $(F\sigma)^{-1}$ increases rapidly with increasing redshift. |
The zero-redshift value of the activity parameter (ho)? is lixed by requiring that the low-reelshilt G0-j4 count prediction is in agreement with observations: see Ll. | The zero-redshift value of the activity parameter $(F\sigma)^{-1}_0$ is fixed by requiring that the low-redshift $\mu$ m count prediction is in agreement with observations; see 1. |
Phe values of the evolution. parameter pe that correspond. to the most reasonable fit for assumed. single | The values of the evolution parameter $p_\sigma$ that correspond to the most reasonable fit for assumed single |
the (2D) orbit (on sky) is not a circular edge-on orbit. but inclined and/or eccentric with LR 7329 D currently near the apastron. hence the small motion on sky. | the (2D) orbit (on sky) is not a circular edge-on orbit, but inclined and/or eccentric with HR 7329 B currently near the apastron, hence the small motion on sky. |
Phe orbital plane of Η 7329 could be in the line of sight («dcge-on like the debris disk aroundD LER. 7329 A). with Lh 7329 D currently near the largest angular separation from LII 7329 X. but with orbital motion mostly in the radial direction: such | The orbital plane of HR 7329 B could be in the line of sight (edge-on like the debris disk around HR 7329 A), with HR 7329 B currently near the largest angular separation from HR 7329 A, but with orbital motion mostly in the radial direction; such |
(1)) we can estimate AL.(final)=Ceetemex μη hence ecc:=Cétetina/te. Where € is a constant fucdec factor of order unitv that depends on the actual time evolution of the SFRs aud My, (Gud eg if it js time dependent). | \ref{eq:dotMstar}) ) we can estimate $M_*({\rm final})=\xi\epsilon_{\rm ff}/t_{\rm ff}\max(M_\H2)t_{\rm final}$ , hence $\epsilon_{\rm GMC}=\xi\epsilon_{\rm ff}t_{\rm final}/t_{\rm ff}$, where $\xi$ is a constant fudge factor of order unity that depends on the actual time evolution of the SFRs and $M_\H2$ (and $\epsilon_{\rm ff}$ if it is time dependent). |
We will estimate & for a simple tov model iu section 3.. | We will estimate $\xi$ for a simple toy model in section \ref{sect:model}. . |
Combining this result with equation (2)) aud (3)) we obtain: There are several wavs of creating large values of Hoare 0nd they correspond to the various terms iu this equation. | Combining this result with equation \ref{eq:GMCeps}) ) and \ref{eq:GMCeta}) ) we obtain: There are several ways of creating large values of $\eta_{\rm GMC}$ and they correspond to the various terms in this equation. |
First. eg could be time depeudenut. | First, $\epsilon_{\rm ff}$ could be time dependent. |
For instance. it could sinoothly increase as the cloud collapse advances or. alternatively. vary stochastically about some average value. | For instance, it could smoothly increase as the cloud collapse advances or, alternatively, vary stochastically about some average value. |
A second possibility is that some clouds may live for many free fall times. ne. fgi/fg Is large in a subset of (λος, | A second possibility is that some clouds may live for many free fall times, i.e. $t_{\rm final}/t_{\rm ff}$ is large in a subset of GMCs. |
The third factor in the third bracket in equation (1)) explain why jeje: can also be than €oaxrc. | The third factor in the third bracket in equation \ref{eq:GMCeps2}) ) explain why $\eta_{\rm GMC}$ can also be than $\epsilon_{\rm GMC}$. |
Finally. ον" can be boosted if the observed IT, muss is siguificautly less than maxoWy,). i.c. if GAICS lose (à one way or another) a large fraction of their molecular hydrogen over their lite time. | Finally, $\eta_{\rm GMC}$ can be boosted if the observed $\H2$ mass is significantly less than $\max(M_\H2)$, i.e. if GMCs lose (in one way or another) a large fraction of their molecular hydrogen over their life time. |
The latter seenario predicts that jesjo(7) should roughly scale κMy over the life time of CAICs. | The latter scenario predicts that $\eta_{\rm GMC}(t)$ should roughly scale $\propto{}M_\H2^{-1}$ over the life time of GMCs. |
Au observational sample of an of GMCs shows this trend (Muay90101. | An observational sample of an of GMCs shows this trend \citep{2010arXiv1007.3270M}. |
However. this trend cau also be produced by a selection effect based on stellar 1iass. e.g. selecting CAICs with M.2μμ excludes values of Hoare that are ΙΕΤΕ see equation (3)). | However, this trend can also be produced by a selection effect based on stellar mass, e.g. selecting GMCs with $M_* > M_{*,{\rm limit}}$ excludes values of $\eta_{\rm GMC}$ that are smaller than $M_{*,{\rm limit}}/M_\H2$, see equation \ref{eq:GMCeta}) ). |
In fact. ἁπαπαν(2010). is selecting clouds based on ionizing hunuinositioes. which roughly corresponds to selecting clouds based the stellar mass formed within the last [1 Myr. | In fact, \cite{2010arXiv1007.3270M} is selecting clouds based on ionizing luminosities, which roughly corresponds to selecting clouds based the stellar mass formed within the last 4 Myr. |
Such a selection effect explains why a different study of ~1011, GMCs find auch lower efficiencies Ladaetal.(2010)... | Such a selection effect explains why a different study of $\sim{}10^5 M_\odot$ GMCs find much lower efficiencies \cite{2010arXiv1009.2985L}. |
The existence of the selection effect is an argunient against or in favor of an evolving €g. rather it shows that the CAICs with large values of Hoare: In the sample of Murray.(2010). are likely a heavily biased. subset. | The existence of the selection effect is an argument against or in favor of an evolving $\epsilon_{\rm ff}$, rather it shows that the GMCs with large values of $\eta_{\rm GMC}$ in the sample of \cite{2010arXiv1007.3270M} are likely a heavily biased subset. |
In the case that eg is. in fact. a non evolving quantity aud the measured laree values of year are driven by changing molecular gas masses. we can dnake a rather generic prediction. | In the case that $\epsilon_{\rm ff}$ is, in fact, a non evolving quantity and the measured large values of $\eta_{\rm GMC}$ are driven by changing molecular gas masses, we can make a rather generic prediction. |
The similarity of the sealing with GAIC (x My) of nac: on the oue haud. aud the lower boundary of the region excluded by the discussed selection. effect. on the other hand. iuplies that the observed GMCS with large values of Hoare: Should have rather similar IT» masses inax(AMg,). | The similarity of the scaling with GMC $\propto{}M_\H2^{-1}$ ) of $\eta_{\rm GMC}$ , on the one hand, and the lower boundary of the region excluded by the discussed selection effect, on the other hand, implies that the observed GMCs with large values of $\eta_{\rm GMC}$ should have rather similar $\H2$ masses $\max(M_\H2)$. |
The tov model that we discuss in section 3 predicts max(Mq,)~108LOTAL... | The toy model that we discuss in section \ref{sect:model} predicts $\max(M_\H2)\sim{}10^6-10^7 M_\odot$. |
We note that this scenario explains rather naturally the absence of iiassive (Gm10° AL.) CMCS with high values of yore. | We note that this scenario explains rather naturally the absence of massive $\gtrsim{}10^6$ $M_\odot$ ) GMCs with high values of $\eta_{\rm GMC}$. |
A different issue can arise df one compares star formation rates and Πο masses in order to estimate ἐμτμ Via equation (1)). | A different issue can arise if one compares star formation rates and $\H2$ masses in order to estimate $\epsilon_{\rm ff}/{t_{\rm ff}}$ via equation \ref{eq:dotMstar}) ). |
For example. let us assiunoe that we measure SFRs aud Ty masses within small (Z100 pe) apertures around peaks of CO cinission (tracing the IL, mass) and peaks of Πα οσο (tracing star formation rates). see e.g. Schrubaetal.(2010)... | For example, let us assume that we measure SFRs and $\H2$ masses within small $\lesssim{}100$ pc) apertures around peaks of CO emission (tracing the $\H2$ mass) and peaks of $H\alpha$ emission (tracing star formation rates), see e.g. \cite{2010arXiv1009.1651S}. |
If we observe that CO peaks have lower SFRs at giveu IL lass colpared with peaks of Πα enüssion. does this iuplv a time-varving eg/fg? | If we observe that CO peaks have lower SFRs at given $\H2$ mass compared with peaks of $H\alpha$ emission, does this imply a time-varying $\epsilon_{\rm ff}/{t_{\rm ff}}$? |
The answer to that question depends ou the way the SFRs are measured. | The answer to that question depends on the way the SFRs are measured. |
SERs that are derived from ΓΓα cinission are effectively averaged over the past 5-10 Myr. which nueht well be a significant action of the lite time of the molecular cloud. | SFRs that are derived from $H\alpha$ emission are effectively averaged over the past 5-10 Myr, which might well be a significant fraction of the life time of the molecular cloud. |
For SFRs hat are based on Πα μυ cinission this averaging iue span would be even longer. | For SFRs that are based on $H\alpha$ $24\mu{}m$ emission this averaging time span would be even longer. |
The star formation cficieucies per free-fall time that are estimated frou such a time averaged SER will be zinall initially (uo stars have (en formed over most of the time averaging interval simply because the GAIC has ouly formed receutlv). | The star formation efficiencies per free-fall time that are estimated from such a time averaged SFR will be small initially (no stars have been formed over most of the time averaging interval simply because the GMC has only formed recently). |
The neasured SFRs will increase until the age of the GMC is siuilar to the averaging time span. | The measured SFRs will increase until the age of the GMC is similar to the averaging time span. |
In additiou. the Il amass of the cloud might evolve (possibly decrease) cading to an additional increase in the appareut value of eg/ty with time. | In addition, the $\H2$ mass of the cloud might evolve (possibly decrease) leading to an additional increase in the apparent value of $\epsilon_{\rm ff}/{t_{\rm ff}}$ with time. |
If the followine three conditions are satisfied. a difference in the measured SER per imceasured IL» mass can provide strong evidence for a tinic-varviug star formation effücienev per free fall time. | If the following three conditions are satisfied, a difference in the measured SFR per measured $\H2$ mass can provide strong evidence for a time-varying star formation efficiency per free fall time. |
First. the averaging times of the SFRs need to be small compared to ages of the observed clouds. | First, the averaging times of the SFRs need to be small compared to ages of the observed clouds. |
Second. the observable IT, reservoirs need to be close to πας.) aud. finally. the free fall times of the clouds need to be known. | Second, the observable $\H2$ reservoirs need to be close to $\max(M_\H2)$, and, finally, the free fall times of the clouds need to be known. |
A recent study that measures SET with reasonably short averaging times (2 Myr. Ladactal.2010)) estimates star formation efficiencies per free fall time of the order of 2! for most clouds in the sample. with the scatter mostly driven bv the mass of molecular eas of relatively low deusity (o<10! 7) that does not participate in the star formation. | A recent study that measures SFRs with reasonably short averaging times (2 Myr, \citealt{2010arXiv1009.2985L}) ) estimates star formation efficiencies per free fall time of the order of $2\%$ for most clouds in the sample, with the scatter mostly driven by the mass of molecular gas of relatively low density $n<10^4$ $^{-3}$ ) that does not participate in the star formation. |
We will now discuss a tov model iu order to both excinplity the poiuts mace im section 2.. but also to provide a framework iu which we can make some quantitative predictions. | We will now discuss a toy model in order to both exemplify the points made in section \ref{sect:pitfalls}, but also to provide a framework in which we can make some quantitative predictions. |
We should stress that the statements made in the previous section are completely eeneric aud do not depend on the specific assuuption that go iuto the model that we are goine to present. | We should stress that the statements made in the previous section are completely generic and do not depend on the specific assumption that go into the model that we are going to present. |
Our model is alinost iusultiuelv simples and. given that. our aiu is not to reproduce the full complexity in the evolution of GMCS or even. to be consistent with auv available observation. | Our model is almost insultingly simple, and, given that, our aim is not to reproduce the full complexity in the evolution of GMCs or even, to be consistent with any available observation. |
On the other haud the model offers a pragluatic approach to the mass evolution of CAICs and iav be casily generalized to facilitate more complex scenarios. | On the other hand the model offers a pragmatic approach to the mass evolution of GMCs and may be easily generalized to facilitate more complex scenarios. |
The iusatzof the model is to supplement equation (1)) with an equivalent equation that describes the evolution of the II» mass: The extra term àM. is motivatedby. assuuing that stellar feedback is limitine the life time of molecular clouds. e.g. via photo-donization. thermal pressure or radiation pressure (Williams&Melee1997:Muay 2010).. | The ansatzof the model is to supplement equation \ref{eq:dotMstar}) ) with an equivalent equation that describes the evolution of the $\H2$ mass: The extra term $\alpha{}M_*$ is motivatedby assuming that stellar feedback is limiting the life time of molecular clouds, e.g. via photo-ionization, thermal pressure or radiation pressure \citep{1997ApJ...476..166W, 2010ApJ...709..191M, 2010arXiv1008.2383L}. |
This feedback. should therefore couple to theformed stellar mass via some eficiency factor a that sets the time scale for the | This feedback should therefore couple to theformed stellar mass via some efficiency factor $\alpha$ that sets the time scale for the |
above test. to our. observations. namely burs and null sequences. returns values Z5; 60. verifving this conclusion of a non-random undermixing. | above test to our observations, namely burst and null sequences, returns values $\simlt$ –60, verifying this conclusion of a non-random “undermixing”. |
Reearding nul periodicity. we find only a suggestion in our observations of a very long periodicitv[αν too long in relation to their tota length for it to be significant. | Regarding null periodicity, we find only a suggestion in our observations of a very long periodicity—far too long in relation to their total length for it to be significant. |
Thus the bursts and nulls of B1944|17 can be regarde as falling into two categories: a) short bursts or nulls of some 1-7 £P, that show a roughly random: distribution. and b) medium to lone bursts or nulls (720 7) tha can occasionally persist for several hundred: pulses and. are patently non-random. | Thus the bursts and nulls of B1944+17 can be regarded as falling into two categories: a) short bursts or nulls of some 1-7 $P_1$ that show a roughly random distribution, and b) medium to long bursts or nulls $>$ 20 $P_1$ ) that can occasionally persist for several hundred pulses and are patently non-random. |
We will elaborate. further on. this distinction in 8&5. | We will elaborate further on this distinction in 5. |
We here investigate the properties. of the four modes identified by DCTII. | We here investigate the properties of the four modes identified by DCHR. |
Following their convention. we refer to the three drift modes as A-C. and the final burst mode as 7D. | Following their convention, we refer to the three drift modes as A-C, and the final burst mode as “D”. |
The defining eharacteristies of modes A-D are the same at both P and L band. as are their frequencies of occurrence. | The defining characteristics of modes A-D are the same at both P and L band, as are their frequencies of occurrence. |
The four modes can be readily distinguished. by eve due to their unique subpulse structures and intensities. as shown in Figure 4.. | The four modes can be readily distinguished by eye due to their unique subpulse structures and intensities, as shown in Figure \ref{colourPS}. |
The transitions between moces occur on a time scale of less than oneae... there are typically no observable "transitions" between modes. | The transitions between modes occur on a time scale of less than one, there are typically no observable “transitions” between modes. |
We lind that within a sequence of 107 pulses there is a high probability of linding at least one occurrence of each mocle. | We find that within a sequence of $10^3$ pulses there is a high probability of finding at least one occurrence of each mode. |
I is interesting that this pulsar. which displavs an almost overwhelming variety of behaviors. is quite reliable in how often it. does SO. | It is interesting that this pulsar, which displays an almost overwhelming variety of behaviors, is quite reliable in how often it does so. |
As seen in the colour polarization cisplav of Fig. 4.. | As seen in the colour polarization display of Fig. \ref{colourPS}, |
the stars mode changes are usually punctuated. by nulls. though there are some combinations of mode changes that characteristically occur adjacent to one another. | the star's mode changes are usually punctuated by nulls, though there are some combinations of mode changes that characteristically occur adjacent to one another. |
As modes A and B exhibit subpulse drifting. they can best be characterized by their /% and. P; values. where £ is defined as the separation of subpulses within a period. anc Py is the separation between drift bands at a fixed. pulse phase. | As modes A and B exhibit subpulse drifting, they can best be characterized by their $P_2$ and $P_3$ values, where $P_2$ is defined as the separation of subpulses within a period, and $P_3$ is the separation between drift bands at a fixed pulse phase. |
Aloce € characteristically displavs an organized vet stationary subpulse structure. | Mode C characteristically displays an organized yet stationary subpulse structure. |
Lastly. we classify those PSs which show no organized subpulse structure as modo D: it is worth noting that mode D is significantly weaker than the others. | Lastly, we classify those PSs which show no organized subpulse structure as mode D; it is worth noting that mode D is significantly weaker than the others. |
Table 2 eives [5 and P. values for modes A. D. and €. The A mode is characterized. by prominent intervals of remarkably precise drifting subpulseswhich is paradoxical considering the stars otherwise unpredictable and discontinuous behavior. | Table \ref{modes} gives $P_2$ and $P_3$ values for modes A, B, and C. The A mode is characterized by prominent intervals of remarkably precise drifting subpulses—which is paradoxical considering the star's otherwise unpredictable and discontinuous behavior. |
Mode. X ids. unique in its regularitv and is marked bv its negativelv-drifting bancs with a roughly 14-74 £5. | Mode A is unique in its regularity and is marked by its negatively-drifting bands with a roughly $P_1$ $P_3$. |
This 14-27) 2; feature can be seen in an Ir of a PS that includes all the moces. indicating its dominance (Weltevrede 2006. 2007).ILwo bright. central subpulses are usually seen in mode A. At L band. weak subpulses on the outer edges of the profile turn on and. olf with a period that is comparable to mode A's £5: see Figure 6.. | This $P_1$ $P_3$ feature can be seen in an lrf of a PS that includes all the modes, indicating its dominance (Weltevrede 2006, 2007).Two bright, central subpulses are usually seen in mode A. At L band, weak subpulses on the outer edges of the profile turn on and off with a period that is comparable to mode A's $P_3$; see Figure \ref{modeA_modfold}. |
Mode A always appears in bursts having durations of more than 15 periods: however. usually these bursts are even longer. typically some GO 100 2, and. remarkably. these A-mocde intervals are very rarely interrupted by nulls. | Mode A always appears in bursts having durations of more than 15 periods; however, usually these bursts are even longer, typically some 60 – 100 $P_1$ —and, remarkably, these A-mode intervals are very rarely interrupted by nulls. |
The drifting subpulses of mode. D. are. visibly less ordered. than those of mode A: however they are clearly structured and negativelv-drifting. | The drifting subpulses of mode B are visibly less ordered than those of mode A; however they are clearly structured and negatively-drifting. |
P; is approximately half that of mode A at both L. and. P band: although clue | $P_3$ is approximately half that of mode A at both L and P band; although due |
resolution image is resolved out iu this image. | resolution image is resolved out in this image. |
One max suspect that the diffuse EHI cussion (particularly that seen in between the three clumps iu the low resolution niap) is not real but is the result of beam smieaxiug. | One may suspect that the diffuse HI emission (particularly that seen in between the three clumps in the low resolution map) is not real but is the result of beam smearing. |
To check for this possibility. the individual channel maps iu the 25«25. data cube were inspected. | To check for this possibility, the individual channel maps in the $25^{''}\times25^{''}$ data cube were inspected. |
In the channel maps. the peak of the diffuse euission in the ceutral region of the galaxv occurs at a differeut beloceutiic velocity than peak velocities of nearby EHI chuups. coutrary to what one would expect from beam smearing. | In the channel maps, the peak of the diffuse emission in the central region of the galaxy occurs at a different heliocentric velocity than peak velocities of nearby HI clumps, contrary to what one would expect from beam smearing. |
As a further confirmation of this. the clean componcuts from the 25425 resolution data cube were couvolyed with a sinaller restoring beam of 10.«10. to eenerate a new data cube, | As a further confirmation of this, the clean components from the $25^{''}\times25^{''}$ resolution data cube were convolved with a smaller restoring beam of $10^{''}\times10^{''}$, to generate a new data cube. |
The diffuse euission is visible iu the channel maps in this cube. contrary to what would have been expected in case the diffuse cussion was cutirely due to beam simeariue (n which case the clean components would have been restricted to the three clumps). | The diffuse emission is visible in the channel maps in this cube, contrary to what would have been expected in case the diffuse emission was entirely due to beam smearing (in which case the clean components would have been restricted to the three clumps). |
As can be seen in Fig. 2.. | As can be seen in Fig. \ref{fig:ov}, |
each HT chump is associated with a chuup of optical emission. | each HI clump is associated with a clump of optical emission. |
However. for each chump. the peak optical enission is generally offset from the peak of the III cmiussion. | However, for each clump, the peak optical emission is generally offset from the peak of the HI emission. |
The Πα iaage of ITodgeetal.(1989) shows that the optical clips also cuit copious amounts of IIo and are heuce regions of ou going star formation. | The $\alpha$ image of \cite{hodge89} shows that the optical clumps also emit copious amounts of $\alpha$ and are hence regions of on going star formation. |
Iu addition to the bright clumps. diffuse optical eiissiou is also seen in Fig. 2.. | In addition to the bright clumps, diffuse optical emission is also seen in Fig. \ref{fig:ov}. |
The optical emission has a πιο higher ellipticitv than the IIT emission aud the position angeles of the optical and ITE major axis can also be seen to be different. | The optical emission has a much higher ellipticity than the HI emission and the position angles of the optical and HI major axis can also be seen to be different. |
Quantitativelv. cllipse fittine to the outermost contours of the 107.&38" aud «425 resolution III monent maps (lich are less distorted by the oxesence of the III chumps in the iuncr regious) gives a position anele of 7T7E5 degrees aud au inclination (assuimiug the oeintrinsic shape of the III disk to be circular) of 2843 degrees. | Quantitatively, ellipse fitting to the outermost contours of the $''\times38''$ and $^{''}\times25^{''}$ resolution HI moment maps (which are less distorted by the presence of the HI clumps in the inner regions) gives a position angle of $\pm$ 5 degrees and an inclination (assuming the intrinsic shape of the HI disk to be circular) of $\pm$ 3 degrees. |
The values obtained from the two different resolution maps agree to within the error bars. | The values obtained from the two different resolution maps agree to within the error bars. |
Ou the other hand. these values are cousiderablv different from those obtained from eclipse fitting to the optical isophotes. which vields a position angle of 38.1 degrees aud an inclination of 57.7 degrees respectively (De Vaucouleurs Moss 1983). | On the other hand, these values are considerably different from those obtained from ellipse fitting to the optical isophotes, which yields a position angle of 38.4 degrees and an inclination of 57.7 degrees respectively (De Vaucouleurs Moss 1983). |
We return to lis issue in Sect. 3.3.. | We return to this issue in Sect. \ref{ssec:discuss}. |
The velocity field derived from the 25<25" resolution data cube is shown iu Fie. l.. | The velocity field derived from the $25^{''}\times 25^{''}$ resolution data cube is shown in Fig. \ref{fig:mom1}. . |
This velocity field is iu reasonable agreement (albeit of better quality) with that obtained by Carignanctal.(1990). | This velocity field is in reasonable agreement (albeit of better quality) with that obtained by \cite{carignan90}. |
. The velocity field shows closed coutours and is. to zeroth order. consisteut with a velocity field that would be produced by a rotating disk with au approximately north southkinematical major axis. | The velocity field shows closed contours and is, to zeroth order, consistent with a velocity field that would be produced by a rotating disk with an approximately north south kinematical major axis. |
Subsets and Splits
No community queries yet
The top public SQL queries from the community will appear here once available.