% PSAMPLE2.TEX -- PASP Conference Proceedings macro package tu % PSAMPLE2.TEX -- PASP Conference Proceedings macro packag \documentstyle[11pt,paspconf,epsf]{article} \documentstyle[11pt,paspconf,epsf]{article} \markboth{}{} \markboth{}{} \setcounter{page}{1} \setcounter{page}{1} \begin{document} \begin{document} \title{ZAMS O Stars} \title{ZAMS O Stars} \author{Margaret M.\ Hanson\altaffilmark{1}} \author{Margaret M.\ Hanson\altaffilmark{1}} \affil{Steward Observatory, University of Arizona, Tucson, AZ \affil{Steward Observatory, University of Arizona, Tucson, \altaffiltext{1}{Hubble Fellow} \altaffiltext{1}{Hubble Fellow} \begin{abstract} \begin{abstract} The combination of large local extinction and fast evolution h The combination of large local extinction and fast evoluti in few direct observations of very young massive stars. For in few direct observations of very young massive stars. F these and other reasons, formation scenarios for the most mass these and other reasons, formation scenarios for the most are poorly constrained. However, new observations made at ne are poorly constrained. However, new observations made a wavelengths can better penetrate the large local extinction in wavelengths can better penetrate the large local extinctio star-forming regions than traditional optical studies. Star- star-forming regions than traditional optical studies. S which contain the youngest massive stars are called which contain the youngest massive stars are called Ultra-Compact H\,{\sc ii} (UC~H\,{\sc ii}) regions. I outline Ultra-Compact H\,{\sc ii} (UC~H\,{\sc ii}) regions. I out classification system for OB stars which allows direct detecti | classification system for OB stars which may allow direct the central ionizing stars of some UC~H\,{\sc ii} regions. Th | the central ionizing stars of some UC~H\,{\sc ii} regions. demonstrated in observations of G29.96$-$0.02, where we have m demonstrated in observations of G29.96$-$0.02, where we ha first direct identification of the central ionizing star in an first direct identification of the central ionizing star i \end{abstract} \end{abstract} % Keywords should be included, but they are not printed in the % Keywords should be included, but they are not printed in \keywords{massive stars, young stellar objects, UC~H\,{\sc ii} \keywords{massive stars, young stellar objects, UC~H\,{\sc \section{Introduction} \section{Introduction} What does it mean for a star to be `zero age'? Certainly such What does it mean for a star to be `zero age'? Certainly is applicable to low-mass stars because they pass through a pr is applicable to low-mass stars because they pass through evolutionary stage to one where the star suddenly begins core evolutionary stage to one where the star suddenly begins c The ignition of hydrogen in its core has a quite sudden and dr The ignition of hydrogen in its core has a quite sudden an further evolution of the star. This change is reflected in a further evolution of the star. This change is reflected i path on the HR diagram. A `zero-age' main-sequence (ZAMS) st path on the HR diagram. A `zero-age' main-sequence (ZAMS minimum radius, its maximum mass (for single-star evolution), minimum radius, its maximum mass (for single-star evolutio (or hottest effective temperature), and its central core posse (or hottest effective temperature), and its central core p This definition, however, breaks down when one considers stars This definition, however, breaks down when one considers s For high-mass stars, contraction occurs exceedingly quickly. For high-mass stars, contraction occurs exceedingly quickl above $M_{\ast} \simeq 10 M_{\odot}$, their Kelvin-Helmholtz ( above $M_{\ast} \simeq 10 M_{\odot}$, their Kelvin-Helmhol {$\tau_{\rm KH} \propto GM_{\ast}^2 / R_{\ast} L_{\ast}$,} bec {$\tau_{\rm KH} \propto GM_{\ast}^2 / R_{\ast} L_{\ast}$,} (formation) timescale. The accretion, or gravitational collap (formation) timescale. The accretion, or gravitational co is the time it takes for the total mass of the star to collaps is the time it takes for the total mass of the star to col molecular cloud and onto the central star, $\tau_{\rm acc} \pr molecular cloud and onto the central star, $\tau_{\rm acc} The pre-main-sequence phase seen in low-mass stars is the obse The pre-main-sequence phase seen in low-mass stars is the having a long contraction timescale, $\tau_{\rm KH} \sim $ few having a long contraction timescale, $\tau_{\rm KH} \sim $ compared to their accretion timescale, {$\tau_{\rm acc} \simeq compared to their accretion timescale, {$\tau_{\rm acc} \s However, the very short contraction time of massive stars lead However, the very short contraction time of massive stars important consequence: they do not pass through any observab important consequence: they do not pass through any obse phase. Instead, a high-mass star will pass directly from a hy phase. Instead, a high-mass star will pass directly from mass-accreting, protostar, in which it is radiative throughout mass-accreting, protostar, in which it is radiative throug to a hydrogen-burning ZAMS star once it has accreted more than to a hydrogen-burning ZAMS star once it has accreted more \ga 10 M_{\odot}$, no matter what its {\it final}\/ mass may b \ga 10 M_{\odot}$, no matter what its {\it final}\/ mass m Thus, all stars with masses greater than $\sim$10$M_{\odot}$, Thus, all stars with masses greater than $\sim$10$M_{\odot O stars, will have begun burning hydrogen well before they eve O stars, will have begun burning hydrogen well before they It would appear that not only do the most massive stars not pa It would appear that not only do the most massive stars no pre-main-sequence phase, but they also never appear as true ZA pre-main-sequence phase, but they also never appear as tru \section{Formation Timescales for Massive Stars} \section{Formation Timescales for Massive Stars} In stating that there is no pre-main-sequence phase in high-ma In stating that there is no pre-main-sequence phase in hig have assumed the accretion timescale will exceed the contracti have assumed the accretion timescale will exceed the contr If by some means, a $60M_{\odot}$ star can accrete to its fina If by some means, a $60M_{\odot}$ star can accrete to its 1000~years, it {\it would}\/ pass through a pre-main-sequence 1000~years, it {\it would}\/ pass through a pre-main-seque albeit a very short one ($\tau_{\rm KH} \simeq 10\,000$ yr). albeit a very short one ($\tau_{\rm KH} \simeq 10\,000$ yr accretion would require a mass infall rate of $\dot M \simeq 6 accretion would require a mass infall rate of $\dot M \sim M_{\odot}$ yr$^{-1}$, which is not necessarily outside the ra M_{\odot}$ yr$^{-1}$, which is not necessarily outside th possibilities (Churchwell 1997, and see below). Let us consi possibilities (Churchwell 1997, and see below). Let us c alternate case, where the accretion to a final $60M_{\odot}$ r alternate case, where the accretion to a final $60M_{\odot fairly long time. If the rate of mass infall onto the growing fairly long time. If the rate of mass infall onto the gro star is slow, say the canonical rate for solar-mass stars $(\d star is slow, say the canonical rate for solar-mass stars M_{\odot}$ yr$^{-1}$; Stahler 1994), then the accretion timesc M_{\odot}$ yr$^{-1}$; Stahler 1994), then the accretion ti $60M_{\odot}$ star will be greater than its main-sequence life $60M_{\odot}$ star will be greater than its main-sequence will have burned nearly all its core hydrogen and be at the e will have burned nearly all its core hydrogen and be at t main-sequence lifetime before it has even reached its final ma main-sequence lifetime before it has even reached its fina effects of applying such an accretion rate to the formation of effects of applying such an accretion rate to the formatio has recently been studied by Bernasconi \& Maeder (1996). The has recently been studied by Bernasconi \& Maeder (1996). obvious effect is it creates a hard limit on how massive a sta obvious effect is it creates a hard limit on how massive a under such conditions. Other interesting effects of a low under such conditions. Other interesting effects of a lo mass-infall rate outlined by Bernasconi \& Maeder include: hig mass-infall rate outlined by Bernasconi \& Maeder include: be seen emerging from their parental clouds already evolved; a be seen emerging from their parental clouds already evolve traditionally used, theoretical ZAMS, which assumes $\tau_{\rm | traditionally used, theoretical ZAMS, which assumes $\tau_ is incorrect at high masses. | \tau{\rm KH}$, is incorrect at high masses. Is there any observational evidence that allows us to determin Is there any observational evidence that allows us to dete timescale or the mass accretion rate for massive stars? timescale or the mass accretion rate for massive stars? Garmany et~al.\ (1982) noted a ``rarity'' of the youngest ($<2 Garmany et~al.\ (1982) noted a ``rarity'' of the youngest stars among Galactic star populations. They suggested that ma stars among Galactic star populations. They suggested tha were still enshrouded in their natal molecular clouds for 1 to were still enshrouded in their natal molecular clouds for may have created the misperception that massive stars are stil may have created the misperception that massive stars are for 1 to 2 Myr. In fact, it only says the youngest stars are for 1 to 2 Myr. In fact, it only says the youngest stars obscured,}\/ $A_V \ga$ several magnitudes, to be undetected at obscured,}\/ $A_V \ga$ several magnitudes, to be undetecte Clearly, accreting, young massive stars will be heavily exting Clearly, accreting, young massive stars will be heavily ex but the reverse logic, that stars with large extinctions stars | but the reverse logic, that stars with large extinctions a true. An O~star may have long ago stopped accretion and compl true. An O~star may have long ago stopped accretion and c outer circumstellar envelope, and yet be completely obscured f outer circumstellar envelope, and yet be completely obscur obscuration may be due to large line-of-sight extinction, or, obscuration may be due to large line-of-sight extinction, due to patches of molecular material from the cloud in which t due to patches of molecular material from the cloud in whi It should be noted that a more complete study of OB stars in b It should be noted that a more complete study of OB stars and the Magellanic Clouds by Massey et~al.\ (1995) was unable and the Magellanic Clouds by Massey et~al.\ (1995) was una if the lack of very young massive stars wasn't just due to a if the lack of very young massive stars wasn't just due t in their data. Such selection effects are likely present in t in their data. Such selection effects are likely present The Carina region (Trumpler 14 \& 16) contains some of the mos The Carina region (Trumpler 14 \& 16) contains some of the stars known in our Galaxy. These stars have inferred masses stars known in our Galaxy. These stars have inferred mass $\sim$100$M_{\odot}$, based on both quantitative $\sim$100$M_{\odot}$, based on both quantitative spectroscopic analysis (Kudritzki 1980; Puls et~al.\ 1996) and spectroscopic analysis (Kudritzki 1980; Puls et~al.\ 1996) comparisons with evolutionary tracks (Massey \& Johnson 1993). comparisons with evolutionary tracks (Massey \& Johnson 19 stars also have ages on the order of 1 million years. Massey stars also have ages on the order of 1 million years. Mas (1997, 1998) find numerous stars within (1997, 1998) find numerous stars within the R136 core to have masses in excess of 120$M_{\odot}$ and a the R136 core to have masses in excess of 120$M_{\odot}$ a of 1 or maybe 2 million years. The existence of young, massiv of 1 or maybe 2 million years. The existence of young, ma fairly solid lower limit for the mass infall rate, $\dot M \ga fairly solid lower limit for the mass infall rate, $\dot M M_{\odot}$ yr$^{-1}$, which is already significantly higher th M_{\odot}$ yr$^{-1}$, which is already significantly highe canonical values determined for solar-mass star. canonical values determined for solar-mass star. Thus, direct observations of very young, very massive stars te Thus, direct observations of very young, very massive star rates are not the same for all stellar masses. This is furthe rates are not the same for all stellar masses. This is fu theoretical considerations. For an isothermal sphere the mass theoretical considerations. For an isothermal sphere the during `inside out' collapse, $\dot M$, during `inside out' collapse, $\dot M$, is given by $\sigma^3/G$, where $\sigma$ is the sound speed (S is given by $\sigma^3/G$, where $\sigma$ is the sound spee In high-mass cores, the sound speed is dominated by a nontherm In high-mass cores, the sound speed is dominated by a nont component (closely related to the magnetic Alfv\'{e}n speed) a component (closely related to the magnetic Alfv\'{e}n spee observed to be higher in high-mass cores than low-mass cores ( observed to be higher in high-mass cores than low-mass cor \& Fuller 1992; Caselli \& Myers 1995). This relation between \& Fuller 1992; Caselli \& Myers 1995). This relation betw and sound speed leads to a higher mass-infall rate, with more- and sound speed leads to a higher mass-infall rate, with m and the formation of more-massive stars. and the formation of more-massive stars. %XXX Margaret - these hyphens are inserted to distinguish betw | %XXX okay! %XXX that are more massive, and the alternative interpretatio < %XXX massive clumps/stars. Is this okay? < This relation has one further This relation has one further important effect. While stars may span a range of $\ga100$ in important effect. While stars may span a range of $\ga100 their formation timescales, over such a mass range, spans mayb | their formation timescales, over such a mass range, span m (Myers \& Fuller 1993). That is to say, {\it all stars,}\/ re | (Myers \& Fuller 1993). This suggests, {\it all stars,}\/ mass, form on the order of a few hundred thousand years (Adams | mass, form on the order of $\tau_{\rm acc} \approx 10^5$ y Further evidence for a high mass-infall rate in high-mass star Further evidence for a high mass-infall rate in high-mass studies of molecular outflows. Bipolar molecular outflows rep studies of molecular outflows. Bipolar molecular outflows star formation in low-mass stars (Shu, Adams \& Lizano 1987). star formation in low-mass stars (Shu, Adams \& Lizano 198 high-mass star-forming regions. These flows have been found t high-mass star-forming regions. These flows have been fou kinetic energy than those from low-mass objects (Shepherd \& C kinetic energy than those from low-mass objects (Shepherd these outflows are similar to those seen in low-mass regions; these outflows are similar to those seen in low-mass regio higher. It can be argued that the mass found in the outflows higher. It can be argued that the mass found in the outfl has recently accreted onto the young massive star and has been has recently accreted onto the young massive star and has (Churchwell 1997). These mass-outflow rates, exceeding $\dot M (Churchwell 1997). These mass-outflow rates, exceeding $\d suggest enormous infall rates for the most massive stars at ea suggest enormous infall rates for the most massive stars a \section{Formation of a High-Mass Stars} \section{Formation of a High-Mass Stars} It has been a decade now since a well-developed theory for the It has been a decade now since a well-developed theory for formation of low-mass stars was first presented (Shu, Adams \& formation of low-mass stars was first presented (Shu, Adam In that time, an immense amount of observational and theoretic In that time, an immense amount of observational and theor the area of star formation has strengthened and refined this the area of star formation has strengthened and refined th model. As currently understood in this `accretion scenario', model. As currently understood in this `accretion scenari star formation begins when high-density cores form inside of m star formation begins when high-density cores form inside clouds and become dynamically unstable. As the core collapses clouds and become dynamically unstable. As the core colla a small hydrostatic object, the protostar, forms at the centre a small hydrostatic object, the protostar, forms at the ce collapse. Material which falls non-radially begins to collect collapse. Material which falls non-radially begins to col rotational equator of the collapse and develops into a circums rotational equator of the collapse and develops into a cir Surrounding the protostar and its circumstellar disk is a more Surrounding the protostar and its circumstellar disk is a infalling envelope of gas and dust, which serves as the reserv | infalling envelope of gas and dust, which serves as the re which the star is fed. As infall continues, the protostar bec | the star and disk system is fed. As infall continues, the massive and more luminous. Eventually the protostar develops massive and more luminous. Eventually the protostar devel and bipolar outflow. and bipolar outflow. How much of this theory for the formation of How much of this theory for the formation of low-mass stars can be applied to high-mass star formation? low-mass stars can be applied to high-mass star formation? Massive stars acquire large luminosities very early in their f Massive stars acquire large luminosities very early in the Their radiation pressure may serve to reverse further infall ( Their radiation pressure may serve to reverse further infa Cassinelli 1986) and disrupt the surrounding circumstellar mat Cassinelli 1986) and disrupt the surrounding circumstellar serving as a mass reservoir. Despite these physical differenc serving as a mass reservoir. Despite these physical diffe high- and low-mass stars, observational evidence suggests that high- and low-mass stars, observational evidence suggests mechanisms involved in the formation of massive stars may be s mechanisms involved in the formation of massive stars may similar to those in low-mass star formation. similar to those in low-mass star formation. For example, outflows are now known to exist in high-mass star For example, outflows are now known to exist in high-mass regions just as in low-mass star-forming regions (Shepherd \& regions just as in low-mass star-forming regions (Shepherd Churchwell 1996). What drives outflows from either low- or hi Churchwell 1996). What drives outflows from either low- o objects (YSOs) is not entirely understood, but they may requir objects (YSOs) is not entirely understood, but they may re enhanced equatorial material (disks). The formation of a dis enhanced equatorial material (disks). The formation of a is a direct and necessary result of the collapse of a dense mo is a direct and necessary result of the collapse of a dens with any net angular momentum (Terebey, Shu \& Cassen 1984). with any net angular momentum (Terebey, Shu \& Cassen 1984 around low-mass young stars have been directly observed (Lay e around low-mass young stars have been directly observed (L In the high-mass YSOs S106 and BN, there In the high-mass YSOs S106 and BN, there is strong evidence that high-density material is closely surro is strong evidence that high-density material is closely s these stars, and both show bipolar winds, perhaps driven throu these stars, and both show bipolar winds, perhaps driven t rarefied polar regions of the young star's circumstellar envel rarefied polar regions of the young star's circumstellar e In fact, arguments can be made In fact, arguments can be made that disks are perhaps required in order to continue feeding m that disks are perhaps required in order to continue feedi young massive star (Jijina \& Adams 1996). Spherical infall i young massive star (Jijina \& Adams 1996). Spherical infa prohibited due to the large radiation pressure of the prohibited due to the large radiation pressure of the young massive star (Wolfire \& Cassinelli 1986). However, it young massive star (Wolfire \& Cassinelli 1986). However, easy for infall to occur to the disk, first. Then, as the dis easy for infall to occur to the disk, first. Then, as the instabilities develop which create episodes of catastrophic ac instabilities develop which create episodes of catastrophi the star from the disk (Adams, Ruden \& Shu 1989). the star from the disk (Adams, Ruden \& Shu 1989). At this very early stage of evolution, the young, massive star At this very early stage of evolution, the young, massive a fairly dense, opaque envelope. If the mass infall rate is n a fairly dense, opaque envelope. If the mass infall rate (see discussion in Churchwell 1997), the young star will ioniz (see discussion in Churchwell 1997), the young star will i ultra-compact H\,{\sc ii} (UC~H\,{\sc ii}) region inside this ultra-compact H\,{\sc ii} (UC~H\,{\sc ii}) region inside t detected at radio wavelengths (Habing \& Israel 1979). The nu detected at radio wavelengths (Habing \& Israel 1979). Th regions found in the Galaxy through radio and mid-infrared sur regions found in the Galaxy through radio and mid-infrared lasts on the order of 10$^5$~years (Wood \& Churchwell 1989; C lasts on the order of 10$^5$~years (Wood \& Churchwell 198 they have linear diameters $\le$ 0.1 pc (Churchwell 1993). As they have linear diameters $\le$ 0.1 pc (Churchwell 1993). expands beyond this diameter, mass infall will have ceased, as expands beyond this diameter, mass infall will have ceased its circumstellar environment become further de-coupled and se its circumstellar environment become further de-coupled an of the UC~H\,{\sc ii} region phase, {$\tau \simeq 10^5$~yr}, a | of the UC~H\,{\sc ii} region phase, {$\tau_{\rm UC} \simeq accretion phase in high-mass stars,}\/ though we cannot rule o | {\it the end of the accretion phase in high-mass stars.} have ended considerably sooner than this time. However, it is | out that accretion could have ended considerably sooner th that one will find the youngest of the most massive stars just | to say $\tau_{\rm acc} \le \tau_{\rm UC}$. It is at the e > UC}$, that one will find the youngest of the most massive their birth cocoons. their birth cocoons. \begin{figure} \begin{figure} %\plotfiddle{EPSFILE}{VSIZE}{ROT}{HSF}{VSF}{HTRANS}{VTRANS} %\plotfiddle{EPSFILE}{VSIZE}{ROT}{HSF}{VSF}{HTRANS}{VTRANS %\plotfiddle{spec.eps}{7cm}{0}{65}{55}{-140}{-10} %\plotfiddle{spec.eps}{7cm}{0}{65}{55}{-140}{-10} \plotfiddle{sed.eps}{7.5cm}{0}{65}{65}{-140}{-10} \plotfiddle{sed.eps}{7.5cm}{0}{65}{65}{-140}{-10} %\plotone{sed.eps} %\plotone{sed.eps} \caption{\small The spectral energy distribution (SED) of the \caption{\small The spectral energy distribution (SED) of from Watson et~al.\ (1997). }\label{sedfig} from Watson et~al.\ (1997). }\label{sedfig} \end{figure} \end{figure} \section{Ultra-Compact H\,{\sc ii} Regions} \section{Ultra-Compact H\,{\sc ii} Regions} Thus far, all information about UC~H\,{\sc ii} regions has bee Thus far, all information about UC~H\,{\sc ii} regions has emission detected from the gas and dust surrounding the centra emission detected from the gas and dust surrounding the ce In Fig.~\ref{sedfig} we show the spectral energy distribution In Fig.~\ref{sedfig} we show the spectral energy distribut for the UC~H\,{\sc ii} region G29.96$-$0.02, from Watson et~al for the UC~H\,{\sc ii} region G29.96$-$0.02, from Watson e regions look nearly identical on such a plot. UC~H\,{\sc ii} regions look nearly identical on such a plot. UC~H\,{\sc are dominated by free-free emission longward of a few mm ($S_{ are dominated by free-free emission longward of a few mm ( \propto \nu^{-0.1}$), due to high-density ($n_{\rm e} \geq 10^ \propto \nu^{-0.1}$), due to high-density ($n_{\rm e} \geq ionized gas. UC~H\,{\sc ii} regions are further distinguished ionized gas. UC~H\,{\sc ii} regions are further distingui in the far- and mid-infrared, consistent with a very cold Plan in the far- and mid-infrared, consistent with a very cold ($T \simeq 25$--30 K) from a massive dust shell. Their high | ($T \simeq 25$--30 K) from a massive, $\ga 10^4 M_{\odot}$ > (Churchwell et al.\,1990). Their high luminosities and unique mid-infrared flux distributions make t luminosities and unique mid-infrared flux distributions ma to detect using IRAS broad-band flux ratios (colours) (Wood \& to detect using IRAS broad-band flux ratios (colours) (Woo 1989; Kurtz et~al.\ 1994). 1989; Kurtz et~al.\ 1994). \begin{figure} \begin{figure} %\plotfiddle{EPSFILE}{VSIZE}{ROT}{HSF}{VSF}{HTRANS}{VTRANS} %\plotfiddle{EPSFILE}{VSIZE}{ROT}{HSF}{VSF}{HTRANS}{VTRANS \plotfiddle{prefig2.ps}{3cm}{0}{35}{25}{-140}{+148} \plotfiddle{prefig2.ps}{3cm}{0}{35}{25}{-140}{+148} \plotfiddle{prefig4.ps}{3cm}{0}{35}{25}{-141}{+122} \plotfiddle{prefig4.ps}{3cm}{0}{35}{25}{-141}{+122} \caption{\small a) On top is the observed spectrum of G45.12+0 \caption{\small a) On top is the observed spectrum of G45. (1996). The solid line at the bottom shows the estimated comp (1996). The solid line at the bottom shows the estimated flux, based on the strength of the Pa$\beta$ line. flux, based on the strength of the Pa$\beta$ line. b) The lower spectrum, also from Lumsden \& Puxley (1996), sh b) The lower spectrum, also from Lumsden \& Puxley (1996), with the free-free component removed, and the line of sight ex with the free-free component removed, and the line of sigh corrected. The blue sloping line represents a simple model fo corrected. The blue sloping line represents a simple mode the starlight while the red sloping line is for a grey-body la the starlight while the red sloping line is for a grey-bod the local dust. The underlying star in this UC~H\,{\sc ii} r the local dust. The underlying star in this UC~H\,{\sc i to detect directly.}\label{g45} to detect directly.}\label{g45} \end{figure} \end{figure} The three rightmost points plotted in the SED of G29 (Fig.~\re The three rightmost points plotted in the SED of G29 (Fig. the near-infrared {\it J, H,} and $K$ bands. The thermal tem the near-infrared {\it J, H,} and $K$ bands. The thermal to heat dust to radiate at these wavelengths is near the dust to heat dust to radiate at these wavelengths is near the d temperature. For this reason, this is the wavelength region w temperature. For this reason, this is the wavelength regi directly from the star's photosphere may first be detected. H directly from the star's photosphere may first be detected dust emission is fading in this wavelength range, dust extinct dust emission is fading in this wavelength range, dust ext increasing. At 2~microns, dust extinction is reduced by more increasing. At 2~microns, dust extinction is reduced by m 10 from the visible ($A_V / A_K \simeq 10$). At 1~micron, the 10 from the visible ($A_V / A_K \simeq 10$). At 1~micron, factor of 3. When one is considering sightlines to stars with factor of 3. When one is considering sightlines to stars of magnitudes of visible extinction, the difference between ob of magnitudes of visible extinction, the difference betwee 1~micron becomes substantial. It was for this reason that we 1~micron becomes substantial. It was for this reason that developed a spectral-classification system at 2~microns (Hanso developed a spectral-classification system at 2~microns (H %XXXHanson, Rieke \& Conti 1996). | Hanson, Conti \& Rieke 1996). The setback of a 2-micron cl %XXX? Hanson Conti & Reike in ref list < Hanson et~al.\ 1996). < The setback of a 2-micron classification system is < that hot dust, at a temperature of $T > 1500$ K, may still dom that hot dust, at a temperature of $T > 1500$ K, may still the youngest systems (Hanson et al.\ 1997; see also Conti \& B the youngest systems (Hanson et al.\ 1997; see also Conti Also, if the young star Also, if the young star sits within a high-density H\,{\sc ii} region, free-free emiss sits within a high-density H\,{\sc ii} region, free-free e background against which the star will be more difficult to de background against which the star will be more difficult t Contamination from free-free emission is illustrated in the Contamination from free-free emission is illustrated in th near-infrared spectrum of the UC~H\,{\sc ii} region G45.12+0.1 near-infrared spectrum of the UC~H\,{\sc ii} region G45.12 (1996). In Fig.~\ref{g45}a) they show that approximately half (1996). In Fig.~\ref{g45}a) they show that approximately the 1- to 2-micron range is attributed to free-free emission. the 1- to 2-micron range is attributed to free-free emissi In Fig.~\ref{g45}b), also taken from Lumsden \& Puxley, they h In Fig.~\ref{g45}b), also taken from Lumsden \& Puxley, th free-free component from the flux and corrected for extinction free-free component from the flux and corrected for extinc flux is now rising to the left). However, they still find add flux is now rising to the left). However, they still find flux longward of 1.5~microns. They attribute this emission to flux longward of 1.5~microns. They attribute this emissio dust. At wavelengths shorter than 1.5~microns, the extinctio dust. At wavelengths shorter than 1.5~microns, the extin overwhelming, yet, longward of 2~microns, the free-free and th overwhelming, yet, longward of 2~microns, the free-free an emission dominate the continuum. This wavelength region, lyin emission dominate the continuum. This wavelength region, near-infrared $H$- and $K$-band, was what Peter Conti and I us near-infrared $H$- and $K$-band, was what Peter Conti and as ``between a rock and a hard place.'' Regardless, this spect as ``between a rock and a hard place.'' Regardless, this s remains the best available for directly detecting the stellar remains the best available for directly detecting the stel deeply embedded young stars, including the deeply embedded ion deeply embedded young stars, including the deeply embedded UC~H\,{\sc ii} region. It should be possible, depending on the UC~H\,{\sc ii} region. It should be possible, depending on dust emission and the line-of-sight extinction, to detect the dust emission and the line-of-sight extinction, to detect sources of {\it many}\/ UC~H\,{\sc ii} regions, but {\it not a sources of {\it many}\/ UC~H\,{\sc ii} regions, but {\it n require near-infrared spectroscopy with relatively high resolu require near-infrared spectroscopy with relatively high re and signal-to-noise (${\rm S/N} > 100$), in particular to over and signal-to-noise (${\rm S/N} > 100$), in particular to mentioned above. mentioned above. \begin{figure} \begin{figure} %\plotfiddle{EPSFILE}{VSIZE}{ROT}{HSF}{VSF}{HTRANS}{VTRANS} %\plotfiddle{EPSFILE}{VSIZE}{ROT}{HSF}{VSF}{HTRANS}{VTRANS \plotfiddle{comp.eps}{3.5cm}{0}{85}{85}{-140}{-20} \plotfiddle{comp.eps}{3.5cm}{0}{85}{85}{-140}{-20} \caption{\small a) The Br$\gamma$ image of G29.96$-$0.02 with \caption{\small a) The Br$\gamma$ image of G29.96$-$0.02 w ionizing source indicated. b) The 2-cm continuum image of G2 ionizing source indicated. b) The 2-cm continuum image o smoothed to resolution of the Br$\gamma$ image, showing the po smoothed to resolution of the Br$\gamma$ image, showing th the ionizing source. In the near-infrared image the stellar so the ionizing source. In the near-infrared image the stella resolved separately from the strong free-free emission peak to resolved separately from the strong free-free emission pea Figure taken from Watson et al.\ (1997).} \label{g29image} Figure taken from Watson et al.\ (1997).} \label{g29image} \end{figure} \end{figure} \subsection{The Central Ionizing Star of G29.96$-$0.02} \subsection{The Central Ionizing Star of G29.96$-$0.02} G29.96$-$0.02 (G29) is a well studied ultra-compact H\,{\sc ii G29.96$-$0.02 (G29) is a well studied ultra-compact H\,{\s mapped in the radio by Wood \& Churchwell (1989) and found to mapped in the radio by Wood \& Churchwell (1989) and found remarkable cometary structure (see Fig.~\ref{g29image}). remarkable cometary structure (see Fig.~\ref{g29image}). About 20\% of all UC~H\,{\sc ii} regions mapped thus far show About 20\% of all UC~H\,{\sc ii} regions mapped thus far s structure (Kurtz et~al.\ 1994 and references therein). UC~H\, structure (Kurtz et~al.\ 1994 and references therein). UC frequently lie behind huge line-of-sight extinctions, renderin frequently lie behind huge line-of-sight extinctions, rend unobservable even at near-infrared wavelengths. However, the unobservable even at near-infrared wavelengths. However, extinction toward G29 is not so prohibitive, with `only' $A_V extinction toward G29 is not so prohibitive, with `only' $ ($A_K \approx 2.5$). Watson et~al.\ (1997) obtained deep {\it ($A_K \approx 2.5$). Watson et~al.\ (1997) obtained deep imaging of the G29 field. They carefully {\it removed}\/ the imaging of the G29 field. They carefully {\it removed}\/ flux from the free-free emission in the $H$- and $K$-band imag flux from the free-free emission in the $H$- and $K$-band more clearly revealed a fairly bright, `blue' star, located at more clearly revealed a fairly bright, `blue' star, locate of the cometary structure seen in the radio and Br$\gamma$ emi of the cometary structure seen in the radio and Br$\gamma$ The position of the star is shown as a white cross in Fig.~\re The position of the star is shown as a white cross in Fig. \& b). When accurate $J, H,$ and $K$-band \& b). When accurate $J, H,$ and $K$-band colours are obtained, intrinsically blue stars ($T_{\rm eff} \ colours are obtained, intrinsically blue stars ($T_{\rm ef separated from cool giants and supergiants ($T_{\rm eff} < 500 separated from cool giants and supergiants ($T_{\rm eff} < colour--colour plot, regardless of line-of-sight extinction (s colour--colour plot, regardless of line-of-sight extinctio et~al.\ 1997). After correcting for extinction and estimate et~al.\ 1997). After correcting for extinction and esti the near-infrared magnitudes of the central source in G29 was | the near-infrared magnitudes and colours of the central so to be consistent with an OB star (Watson et~al.\ 1997). to be consistent with an OB star (Watson et~al.\ 1997). \begin{figure}[t] \begin{figure}[t] %\plotfiddle{EPSFILE}{VSIZE}{ROT}{HSF}{VSF}{HTRANS}{VTRANS} %\plotfiddle{EPSFILE}{VSIZE}{ROT}{HSF}{VSF}{HTRANS}{VTRANS \plotfiddle{spec.eps}{7cm}{0}{65}{55}{-140}{-10} \plotfiddle{spec.eps}{7cm}{0}{65}{55}{-140}{-10} \caption{$K$-band spectra of G29.96$-$0.02. At the top is the \caption{$K$-band spectra of G29.96$-$0.02. At the top is spectrum of the bright nebular peak to the SW of the stellar s spectrum of the bright nebular peak to the SW of the stell nebular corrections; `Stellar~C' is over-corrected (note the o nebular corrections; `Stellar~C' is over-corrected (note t features of He\,{\sc i} at 2.058~$\mu$m and Br$\gamma$ at 2.16 features of He\,{\sc i} at 2.058~$\mu$m and Br$\gamma$ at spectroscopic standards from Hanson et~al.\ (1996) are also sh spectroscopic standards from Hanson et~al.\ (1996) are als \end{figure} \end{figure} An important morphological aspect of the G29 region is that th An important morphological aspect of the G29 region is tha bright nebular knot detected in the radio and infrared is {\it bright nebular knot detected in the radio and infrared is on the star}\/ (Fig.\ 3). If this had been the case, the spec on the star}\/ (Fig.\ 3). If this had been the case, the be dominated by free-free and thermal dust emission, as shown be dominated by free-free and thermal dust emission, as s spectrum in Figure~\ref{g45}. Despite the angular distance be spectrum in Figure~\ref{g45}. Despite the angular distanc and the bright nebular knot, strong nebular features were stil and the bright nebular knot, strong nebular features were at the position of the star in its 2-$\mu$m spectrum. Figure~ at the position of the star in its 2-$\mu$m spectrum. Fig Watson \& Hanson 1997) shows the raw, uncorrected spectrum of Watson \& Hanson 1997) shows the raw, uncorrected spectrum source, labeled `Stellar A'. The spectrum labeled `Nebular' w source, labeled `Stellar A'. The spectrum labeled `Nebula the Br$\gamma$ peak and scaled to match the nebular features i the Br$\gamma$ peak and scaled to match the nebular featur `Stellar~A' spectrum. The spectrum labeled `Stellar B' is our `Stellar~A' spectrum. The spectrum labeled `Stellar B' is at removing the nebular contamination using the scaled nebular at removing the nebular contamination using the scaled neb free-free emission is contributing about 20\% of the flux at 2 free-free emission is contributing about 20\% of the flux `Stellar~C' shows the stellar spectrum over-corrected for nebu `Stellar~C' shows the stellar spectrum over-corrected for Even with its over-corrected nebular component, `Stellar~C' st Even with its over-corrected nebular component, `Stellar~C a broad emission feature at $\sim$2.115$\mu$m, which is N\,{\s a broad emission feature at $\sim$2.115$\mu$m, which is N\ from the star. Further confirmation that we have detected the from the star. Further confirmation that we have detected photosphere is provided in the strong stellar C\,{\sc iv} feat photosphere is provided in the strong stellar C\,{\sc iv} and He\,{\sc ii} at 2.1885$\mu$m. The ionizing star of G29 lo and He\,{\sc ii} at 2.1885$\mu$m. The ionizing star of G2 kO5-O6 ($K$-band classification; see Hanson et~al.\ 1996). Ty kO5-O6 ($K$-band classification; see Hanson et~al.\ 1996). spectra are shown as comparisons in Fig.~\ref{g29spectra}. spectra are shown as comparisons in Fig.~\ref{g29spectra}. \subsection{Previous Analyses of G29.96$-$0.02} \subsection{Previous Analyses of G29.96$-$0.02} Simpson et~al.\ 1995 and Afflerbach et~al.\ 1997 have modeled Simpson et~al.\ 1995 and Afflerbach et~al.\ 1997 have mode and found best fits with ionizing stars that have effective t and found best fits with ionizing stars that have effecti we now {\it know}\/ the temperature of the stellar ionizing s we now {\it know}\/ the temperature of the stellar ionizi with an effective temperature between 41\,000 and 44\,000 K. with an effective temperature between 41\,000 and 44\,000 38\,000 K (an O8 star) and there is evidence, too, that the c 38\,000 K (an O8 star) and there is evidence, too, that t is too hot (Harries \& Hilditch 1997; Hubeny 1997). However, is too hot (Harries \& Hilditch 1997; Hubeny 1997). Howe Perhaps G29 is not a good choice for mid-infrared analyses? Perhaps G29 is not a good choice for mid-infrared analyses morphology which could be due to motion of the star relative t morphology which could be due to motion of the star relati shock (Van Buren \& Mac\,Low 1992), or it may have entered a shock (Van Buren \& Mac\,Low 1992), or it may have entere 1996). A successful analysis of G29 may require a more comp 1996). A successful analysis of G29 may require a more account the region's non-spherical geometry and explores a la account the region's non-spherical geometry and explores Improved stellar-atmosphere models which properly predict vari Improved stellar-atmosphere models which properly predict high-metallicity ionizing stars, particularly over the ultrav high-metallicity ionizing stars, particularly over the ul relative ionizations of the observed elements, are also criti relative ionizations of the observed elements, are also c Clearly, having direct detections of the central ionizing sou Clearly, having direct detections of the central ionizing be useful. be useful. Cometary morphologies are seen in about 20\% of all UC~H\,{\sc Cometary morphologies are seen in about 20\% of all UC~H\, et~al.\ 1994). Other morphologies seen in UC~H\,{\sc ii} regi et~al.\ 1994). Other morphologies seen in UC~H\,{\sc ii} the central star may be currently breaking out of its cocoon. the central star may be currently breaking out of its coco able to obtain spectra of the central ionizing sources for all able to obtain spectra of the central ionizing sources for sources at 2$\mu$m? I am currently exploring this idea. The sources at 2$\mu$m? I am currently exploring this idea. are difficult. The radio maps have high spatial resolution, w are difficult. The radio maps have high spatial resolutio separation between star and H\,{\sc ii} region appear deceptiv separation between star and H\,{\sc ii} region appear dece of the cometaries, the stellar source may be visible at 2~micr of the cometaries, the stellar source may be visible at 2~ spatially blended ($\la 1''$) with its nearby dense envelope o spatially blended ($\la 1''$) with its nearby dense envelo for observations from the ground. Either adaptive optics, to for observations from the ground. Either adaptive optics, the central point source from its bright extended cocoon, or s the central point source from its bright extended cocoon, at extremely high at extremely high signal-to-noise, to detect the diluted stellar photospheric fe signal-to-noise, to detect the diluted stellar photospheri will be needed. In most cases, both techniques may be needed. will be needed. In most cases, both techniques may be nee \section{Final Comments} \section{Final Comments} It would appear that `zero age' is perhaps a misnomer when use It would appear that `zero age' is perhaps a misnomer when very young massive stars ($M > 40 M_{\odot}$). Unless $\dot very young massive stars ($M > 40 M_{\odot}$). Unless $\ M_{\odot}$ yr$^{-1}$ (which may certainly be the case), massiv M_{\odot}$ yr$^{-1}$ (which may certainly be the case), ma burning hydrogen, ionizing H\,{\sc ii} regions, creating wind burning hydrogen, ionizing H\,{\sc ii} regions, creating w {\it evolving}\/ since well before they achieve their final ma {\it evolving}\/ since well before they achieve their fina `zero age', or ZAMS, perhaps better represents that stage in e `zero age', or ZAMS, perhaps better represents that stage when the star, having reached its final mass, is seen at its m when the star, having reached its final mass, is seen at i main-sequence effective temperature and minimum radius. This main-sequence effective temperature and minimum radius. T be applied to ZAMS stars of all masses. This said, however, t be applied to ZAMS stars of all masses. This said, howeve theoretical ZAMS, based on instantaneous hydrogen ignition at theoretical ZAMS, based on instantaneous hydrogen ignition not be the same as the observed ZAMS at very high masses. The not be the same as the observed ZAMS at very high masses. using the above definition, probably lies to cooler effective using the above definition, probably lies to cooler effect and its exact position depends on the timescale $\tau_{\rm acc and its exact position depends on the timescale $\tau_{\rm timescale is thus critical to the development of a coherent th timescale is thus critical to the development of a coheren massive-star formation, yet it remains poorly constrained. massive-star formation, yet it remains poorly constrained. This is why observations of UC~H\,{\sc ii} regions are importa | Observations of UC~H\,{\sc ii} regions are important for c 2~decades, the study of these sources at radio and mm waveleng | accretion timescale in high mass stars. For more than 2~ information available on the early stages of massive-star form | of the gas and dust from these sources at radio and mm wav > the best information on the early stages of massive-star f However, recent instrumental advances now allow us to extend o However, recent instrumental advances now allow us to exte slightly shorter wavelengths. Through the use of deep near-in slightly shorter wavelengths. Through the use of deep nea high-resolution, near-infrared spectroscopy, we can search for high-resolution, near-infrared spectroscopy, we can search detect very young, very massive stars detect very young, very massive stars while they are still heavily buried within their birth cocoons while they are still heavily buried within their birth coc observations of massive stars at this young age, which had pre observations of massive stars at this young age, which had entirely inaccessible, will allow new questions to be answered entirely inaccessible, will allow new questions to be answ formation and early evolution of massive stars. For the vast formation and early evolution of massive stars. For the v UC~H\,{\sc ii} regions, their central star (or stars!) will re UC~H\,{\sc ii} regions, their central star (or stars!) wil regardless of the wavelength. However, I am optimistic that a regardless of the wavelength. However, I am optimistic th regions may be dispersed just enough that their very young cen regions may be dispersed just enough that their very young directly detected at 2$\mu$m. What is useful about detecting directly detected at 2$\mu$m. What is useful about detec powering UC~H\,{\sc ii} regions is then we know they are very powering UC~H\,{\sc ii} regions is then we know they are v from the approximate age of the UC~H\,{\sc ii} region phase ($ from the approximate age of the UC~H\,{\sc ii} region phas Determining the age of young O stars to an accuracy better tha Determining the age of young O stars to an accuracy better is simply not possible if we rely on evolutionary isochrones. is simply not possible if we rely on evolutionary isochron An important future goal is to directly detect the youngest of An important future goal is to directly detect the younges stars for subsequent quantitative spectroscopic analyses. Thi stars for subsequent quantitative spectroscopic analyses. require quantitative analyses to be extended to the near-infra require quantitative analyses to be extended to the near-i in an observational and a theoretical sense, this is only now in an observational and a theoretical sense, this is only I and others are currently obtaining high-resolution, high sig I and others are currently obtaining high-resolution, high spectroscopic observations of O stars in an effort to identify spectroscopic observations of O stars in an effort to iden lines through out the near-infrared wavelength region. Furth lines through out the near-infrared wavelength region. F are now being made to include near-infrared line syntheses in are now being made to include near-infrared line syntheses atmosphere codes (Najarro et~al.\ 1994; Schaerer et~al. 1996) atmosphere codes (Najarro et~al.\ 1994; Schaerer et~al. 1 analysis of ZAMS O~stars and their physical conditions at earl analysis of ZAMS O~stars and their physical conditions at stages promises to be an exciting new direction of research th stages promises to be an exciting new direction of researc better understand how the most massive stars form. better understand how the most massive stars form. \acknowledgments \acknowledgments I am grateful to Phil Puxley, Stan Lumsden, and Alan Watson, f | I am grateful to Phil Puxley, Stuart Lumsden, and Alan Wat use figures from their papers. My thoughts, ideas and work on use figures from their papers. My thoughts, ideas and wor and UC~H\,{\sc ii} regions have been inspired through discussi and UC~H\,{\sc ii} regions have been inspired through disc Ed Churchwell, John Bieging, and Fred Adams. Ed Churchwell, John Bieging, and Fred Adams. \begin{references} \begin{references} \reference Adams, F.\,C., Ruden, S.\,P., Shu, F.\,H., 1989, \a \reference Adams, F.\,C., Ruden, S.\,P., Shu, F.\,H., 1989 \reference Adams, F.\,C., Fatuzzo, M., 1996, XXX, 464, 256 | \reference Adams, F.\,C., Fatuzzo, M., 1996, \apj, 464, 25 %XXX Which journal? < \reference Afflerbach, A., Churchwell, E., Werner, M., 1997, \ \reference Afflerbach, A., Churchwell, E., Werner, M., 199 \reference Bernasconi, P.\,A., Maeder, A., 1996, \aap, 307, 82 \reference Bernasconi, P.\,A., Maeder, A., 1996, \aap, 307 \reference Caselli, P., Myers, P.\,C., 1995, \apj, 446, 665 \reference Caselli, P., Myers, P.\,C., 1995, \apj, 446, 66 \reference Churchwell, E., 1993, in XXX | \reference Churchwell, E., 1993, in {\it Massive Stars: Th (ASP Conf.\ Series, Vol.~35), eds.~XXX, p.~35 | medium,} (ASP Conf.\ Series, Vol.~35), eds.~J.\,P.~Cassine > \reference Churchwell, E., Wolfire, M.\,G., Wood, D.\,O.\, \reference Churchwell, E., 1997, \apj, 479, L59 \reference Churchwell, E., 1997, \apj, 479, L59 \reference Comeron, F., Torra, J., 1996, \aap, 308, 565 \reference Comeron, F., Torra, J., 1996, \aap, 308, 565 \reference Conti, P.\,S., Blum, R.\,D., 1997, in \reference Conti, P.\,S., Blum, R.\,D., 1997, in Boulder-Munich II: Properties of Hot, Luminous Stars Boulder-Munich II: Properties of Hot, Luminous Stars (ASP Conf.\ Series, Vol.~XXX), ed.~I.\,D.~Howarth (ASP, San Fr (ASP Conf.\ Series, Vol.~XXX), ed.~I.\,D.~Howarth (ASP, Sa \reference Drew, J., 1997, \reference Drew, J., 1997, in Boulder-Munich II: Properties of Hot, Luminous Stars in Boulder-Munich II: Properties of Hot, Luminous Stars (ASP Conf.\ Series, Vol.~XXX), ed.~I.\,D.~Howarth (ASP, San Fr (ASP Conf.\ Series, Vol.~XXX), ed.~I.\,D.~Howarth (ASP, Sa \reference Garmany, C.\,D., Conti, P.\,S., Chiosi, C., 1982, \ \reference Garmany, C.\,D., Conti, P.\,S., Chiosi, C., 198 \reference Habing, H.\,J., Israel, F.\,P., 1979, \araa, 17, 34 \reference Habing, H.\,J., Israel, F.\,P., 1979, \araa, 17 \reference Hanson, M.\,M., Conti, P.\,S., 1994, \apj, 423, L13 \reference Hanson, M.\,M., Conti, P.\,S., 1994, \apj, 423, \reference Hanson, M.\,M. Conti, P.\,S., Rieke, M.\,J., 1996, \reference Hanson, M.\,M. Conti, P.\,S., Rieke, M.\,J., 19 \reference Hanson, M.\,M., Howarth, I.\,D., Conti, P.\,S., 199 | \reference Hanson, M.\,M., Howarth, I.\,D., Conti, P.\,S., \reference Harries, T.\,J., Hilditch, R.\,W., 1997, \reference Harries, T.\,J., Hilditch, R.\,W., 1997, in Boulder-Munich II: Properties of Hot, Luminous Stars in Boulder-Munich II: Properties of Hot, Luminous Stars (ASP Conf.\ Series, Vol.~XXX), ed.~I.\,D.~Howarth (ASP, San Fr (ASP Conf.\ Series, Vol.~XXX), ed.~I.\,D.~Howarth (ASP, Sa \reference Hubeny, I., 1997, \reference Hubeny, I., 1997, in Boulder-Munich II: Properties of Hot, Luminous Stars in Boulder-Munich II: Properties of Hot, Luminous Stars (ASP Conf.\ Series, Vol.~XXX), ed.~I.\,D.~Howarth (ASP, San Fr (ASP Conf.\ Series, Vol.~XXX), ed.~I.\,D.~Howarth (ASP, Sa \reference Jijina, J., Adams, F.\,C., 1996, 462, 874 \reference Jijina, J., Adams, F.\,C., 1996, 462, 874 \reference Kudritzki, R.-P., 1980, \aap, 85, 174 \reference Kudritzki, R.-P., 1980, \aap, 85, 174 \reference Kurtz, S., Churchwell, E., Wood, D., 1994, \apjs, 9 \reference Kurtz, S., Churchwell, E., Wood, D., 1994, \apj \reference Lay, O., Carlstrom, J.\,E., Hills, R.\,E., Phillips \reference Lay, O., Carlstrom, J.\,E., Hills, R.\,E., Phil \reference Lumsden, S.\,L., Hoare, M.\,G., 1996, \apj, 464, 27 \reference Lumsden, S.\,L., Hoare, M.\,G., 1996, \apj, 464 \reference Lumsden, S.\,L., Puxley, P.\,J., 1996, \mnras, 281, \reference Lumsden, S.\,L., Puxley, P.\,J., 1996, \mnras, \reference Massey, P., Johnson, J., 1993, \aj, 105, 980 \reference Massey, P., Johnson, J., 1993, \aj, 105, 980 \reference Massey, P., et~al., 1995, \apj, 438, 188 | \reference Massey, P., Lang, C.\,C., Degioia-Eastwood, K., %XXX Please add authors < \reference Massey, P., Hunter, D., 1998, \apj, in press \reference Massey, P., Hunter, D., 1998, \apj, in press \reference Massey, P., Hunter, D., 1997, \reference Massey, P., Hunter, D., 1997, in Boulder-Munich II: Properties of Hot, Luminous Stars in Boulder-Munich II: Properties of Hot, Luminous Stars (ASP Conf.\ Series, Vol.~XXX), ed.~I.\,D.~Howarth (ASP, San Fr (ASP Conf.\ Series, Vol.~XXX), ed.~I.\,D.~Howarth (ASP, Sa \reference Myers, P.\,C., Fuller, G.\,A., 1992, \apj, 396, 631 \reference Myers, P.\,C., Fuller, G.\,A., 1992, \apj, 396, \reference Myers, P.\,C., Fuller, G.\,A., 1993, \apj, 402, 635 \reference Myers, P.\,C., Fuller, G.\,A., 1993, \apj, 402, \reference Najarro, F., et~al.\ 1994, \aap, 285, 573 | \reference Najarro, F., Hillier, D.\,J., Kudritzki, R.\,P. %XXX Please add authors | Drapatz, S., Geballe, T.\,R., 1994, \aap, 285, 573 \reference Palla, F., Stahler, S.\,W., 1993, \apj, 418, 414 \reference Palla, F., Stahler, S.\,W., 1993, \apj, 418, 41 \reference Puls, J., et~al., 1996, \aap, 305, 171 | \reference Puls, J., Kudritzki, R.-P., Herrero, A., Pauldr %XXX Please add authors | Lennon, D.\,J., Gabler, R., Voels, S.\,A., Vilchez, J.\,M. \reference Schaerer, D., de Koter, A., Schmutz, W., Maeder, A. \reference Schaerer, D., de Koter, A., Schmutz, W., Maeder \reference Shepherd, D.\,S., Churchwell, E., 1996, 457, 267 \reference Shepherd, D.\,S., Churchwell, E., 1996, 457, 26 \reference Shu, F.\,H., 1977, \apj, 214, 488 \reference Shu, F.\,H., 1977, \apj, 214, 488 \reference Shu, F.\,H., Adams, F.\,C., Lizano, S., 1987, \araa \reference Shu, F.\,H., Adams, F.\,C., Lizano, S., 1987, \ \reference Simpson, J.\,P., et~al., 1995, \apj, 444, 721 | \reference Simpson, J.\,P., Colgan, S.\,W.\,J., Rubin, R.\ %XXX Please add authors | 1995, \apj, 444, 721 \reference Stahler, S.\,W., 1994, \pasp, 106, 337 \reference Stahler, S.\,W., 1994, \pasp, 106, 337 \reference Terebey, S., Shu, F.\,H, Cassen, P., 1984, \apj, 28 \reference Terebey, S., Shu, F.\,H, Cassen, P., 1984, \apj \reference Van Buren, D., Mac\,Low, M.-M., 1992, \apj, 394, 53 \reference Van Buren, D., Mac\,Low, M.-M., 1992, \apj, 394 \reference Watson, A.\,M., et~al., 1997, \apj, in press | \reference Watson, A.\,M., Coil, A.\,L., Shepherd, D.\,S., %XXX Please add authors < \reference Watson, A.\,M., Hanson, M.\,M., 1997, \apj, in pres \reference Watson, A.\,M., Hanson, M.\,M., 1997, \apj, in \reference Wolfire, M.\,G., Cassinelli, J.\,P., 1986, \apj, 31 \reference Wolfire, M.\,G., Cassinelli, J.\,P., 1986, \apj \reference Wood, D., Churchwell, E., 1989, \apjs, 69, 831 \reference Wood, D., Churchwell, E., 1989, \apjs, 69, 831 \end{references} \end{references} \bigskip \bigskip \centerline{\bf Discussion} \centerline{\bf Discussion} \medskip \medskip \noindent{\bf Kaper:} I'm interested in the bow shocks; \noindent{\bf Kaper:} I'm interested in the bow shocks; I have a colleague at ESO who models these with hydro codes, a I have a colleague at ESO who models these with hydro code shock as soon as you have an O-star wind blowing against a mol shock as soon as you have an O-star wind blowing against a cloud. If you add some relative velocity you create a bow sho cloud. If you add some relative velocity you create a bow you think you get such a shock? you think you get such a shock? \noindent{\bf Hanson:} You just described it! But there's | \noindent{\bf Hanson:} You just described it! There's problem of how you get ultracompacts to live as long as the st | how you get ultracompacts to live as long as the statistic suggest. About 5 years ago one of the best suggestions, by Ma | About 5 years ago one of the best suggestions, by Mac\,Low was simply that if you move the star relative to the cloud you was simply that if you move the star relative to the cloud high-density environment much longer, even though the cocoon m high-density environment much longer, even though the coco away. Now I think that's actually happening in some situation away. Now I think that's actually happening in some situa certainly doesn't explain all of them. | doesn't explain all of them. \noindent{\bf Kaper:} But then it doesn't necessarily put a c \noindent{\bf Kaper:} But then it doesn't necessarily put age of the stars? age of the stars? \noindent{\bf Hanson:} I can't say anything about the age of \noindent{\bf Hanson:} I can't say anything about the ag thing you can say now is, what is the timescale for it to leav thing you can say now is, what is the timescale for it to That's probably on the order of millions of years. | For small ($\le 0.1$ pc), dense ($n_e \ge 10^5$ cm$^{-3}$) > the timescale must be very short (few times 10$^5$ yr). F > created by runaway-OB stars, they might be maintained for \noindent{\bf Drew:} There's an alternative view, the `bliste | \noindent{\bf Drew:} There's an alternative view for G29. model. Recent high-resolution work by Stuart Lumsden and Melv model. Recent high-resolution work by Stuart Lumsden and if anything, come down in favour of this, rather than bow shoc if anything, come down in favour of this, rather than bow \noindent{\bf Hanson:} Yes, there the density's not quite the | \noindent{\bf Hanson:} Yes, there the density's not the s round, and you get blow-out on one side. In fact, my co-autho | round, and you get blow-out on one side. In fact, my co-a believes in `champagne'. Since I don't work on the models mys | Alan Watson, prefers the `champagne' model. Since I don't accept what the bow-shock people say, which is that the veloci | myself, I've accepted what the bow-shock people say. The from radio recombination lines are in better agreement with th | from radio recombination lines are in better agreement wit with `champagne' -- I guess it hasn't been sorted out yet. Th with `champagne' -- I guess it hasn't been sorted out yet. bringing that up. bringing that up. \noindent{\bf Najarro:} Some of the shapes you've shown look \noindent{\bf Najarro:} Some of the shapes you've shown the `pistol' feature at the Galactic Centre, and most of us th the `pistol' feature at the Galactic Centre, and most of u nebula is the result of an LBV which erupted around a million nebula is the result of an LBV which erupted around a mill What is the mass of you nebulosity? At the Galactic Centre the | What is the mass of your nebulosity? At the Galactic Centr 10$M_\odot$. Can you measure the recombination-line radial vel 10$M_\odot$. Can you measure the recombination-line radial and compare them with the radial velocity of the star? | and compare them with the radial velocity of the star in G \noindent{\bf Hanson:} The velocities range from 85 to 115 km | \noindent{\bf Hanson:} Typical dust masses for UC~H\,{\sc across the region. Those velocities `smear' spatially across | $\ga 10^4 M_{\odot}$. The radial velocities range from 85 which is what supports the idea of a bow shock. | across the G29.96-0.02 region. Those velocities `smear' s > which is what supports the idea of a bow shock. The stell > not known. Since we've identified stellar features at 2~$ > be possible to obtain the stellar velocity. \noindent{\bf Maeder:} You've been speaking about the accret \noindent{\bf Maeder:} You've been speaking about the ac There are {\it two}\/ accretion rates: one from the molecular There are {\it two}\/ accretion rates: one from the molecu desk, and the second from the disk onto the star. It's not cl | disk, and the second from the disk onto the star. It's no you're observing, and they're not necessarily the same, as the you're observing, and they're not necessarily the same, as long delay between accretion onto the disk and accretion onto long delay between accretion onto the disk and accretion o \noindent{\bf Hanson:} What I'm told is that the situation is \noindent{\bf Hanson:} What I'm told is that the situatio the disk is much more massive than the central star. Over lon the disk is much more massive than the central star. Over timescales, therefore, at least half the (disk) accreted mass timescales, therefore, at least half the (disk) accreted m going onto the star -- you can't build up lots of mass into th going onto the star -- you can't build up lots of mass int substantially smaller accretion onto the star. The rates have substantially smaller accretion onto the star. The rates less the same. less the same. \end{document} \end{document}