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The HII Region KR 140 Spontaneous Formation of a High Mass Star

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THE HII REGION KR 140:SPONTANEOUS FORMATION OF A HIGH MASS STAR D.R.Ballantyne 1,C.R.Kerton 2and P.G.Martin Canadian Institute for Theoretical Astrophysics,University of Toronto,Toronto,ON,Canada M5S 3H8;ballanty,kerton and pgmartin@cita.utoronto.ca ABSTRACT We have used a multiwavelength data set from the Canadian Galactic Plane Survey (CGPS)to study the Galactic H II region KR 140,both on the scale of the nebula itself and in the context of the star forming activity in the nearby W3/W4/W5complex of molecular clouds and H II regions.From both radio and infrared data we have found a covering factor of about 0.5for KR 140and we interpret the nebula as a bowl-shaped region viewed close to face on.Extinction measurements place the region on the near side of its parent molecular cloud.The nebula is kept ionized by one O8.5V(e)star,VES 735,which is less than a few million years old.CO data show that VES 735has disrupted much of the original molecular cloud for which the estimated mass and density are about 5000M ⊙and 100cm ?3,respectively.KR 140is isolated from the nearest star forming activity,in W3.Our data suggest that KR 140is an example of spontaneous (i.e.,non-triggered)formation of,unusually,a high mass star.Subject headings:HII regions —ISM:individual (KR 140)—stars:formation 1.INTRODUCTION

Massive OB stars are almost always found in clusters.In fact,accumulating observational evidence suggests that most stars,regardless of mass,actually form as members of some kind of group,cluster,or association.While there has been a lot of work to understand the processes involved in the formation of a single star (e.g.,Shu,Adams,&Lizano 1987),theories of cluster formation are still in their infancy (see the recent reviews by Elmegreen et al.2000,and Clarke,Bonnell &Hillenbrand 2000).Nevertheless,many important issues have already been identi?ed.Foremost among these is the question of whether the formation of clusters,particularly ones with OB stars,are triggered by an external agent (Elmegreen 1992)such as expanding H II regions (Elmegreen &Lada 1977),or colliding molecular clouds (Loren 1976,1977;Scoville,Sanders &

Clemens1986;Usami,Hanawa,&Fujimoto1995).Despite triggered or sequential star formation being theoretically and intuitively appealing,a major problem is to determine unambiguously

a cause and e?ect relationship,because of the long time scales(and time lags)of the processes involved.There is observational evidence for regions that have had triggered star formation,both on large scales(e.g.,IC1396,Patel et al.1998)and on small scales(e.g.,IC1805,Heyer et al. 1996).However,there are other young star forming regions,like Taurus,where no evidence of a trigger can been found.These are often in modest-sized molecular clouds and contain only lower mass stars.

The Perseus Arm star forming regions W3/W4/W5(Westerhout1958)have been studied extensively over the last twenty years(e.g.,Lada et al.1978;Braunsfurth1983;Digel et al.1996; Normandeau,Taylor&Dewdney1997;Heyer&Terebey1998).They are often considered to be the archetypical examples of how the formation of massive star clusters can be triggered by the in?uence of other nearby clusters.For example,W3is thought to be have been triggered by the expansion of W4(Dickel1980;Thronson,Lada,&Hewagama1985;van der Werf&Goss1990) and there is evidence that the expansion of W5is also triggering star formation(Vall′e e,Hughes, &Viner1979;Wilking et al.1984).

In new high resolution multiwavelength(radio and mid-infrared)data of the W3/W4/W5 complex from the Canadian Galactic Plane Survey(CGPS;English et al.1998)we have identi?ed a star forming region containing a single O star.Figure1shows a1420MHz continuum image from the CGPS pilot project(Normandeau et al.1997).The circled area is the H II region in question,KR140(l=133.425?,b=0.055?;Kallas&Reich1980).This region appears to be completely separate from the vigorous star formation going on nearby in W3,although it is in the same Perseus arm molecular complex(see context in CO images in Heyer et al.1998).What is unusual is that this massive star seems to have been formed spontaneously.

In this paper we present and analyze the multiwavelength data on KR140in order to quantify the properties of this region of spontaneous massive star formation.As described in§2these data have su?cient resolution(1′)to resolve this H II region for the?rst time and reveal a fairly symmetrical structure.Based on?rst impressions,we thought that KR140might prove to be a “textbook”spherical H II region.Instead,we?nd that KR140is likely a bowl-shaped H II region (an example of a blister geometry:Israel1978,Yorke et al.1989;§3and§4.2.2,§6.4).We have found that the H II region is kept ionized by an O8.5V(e)star,VES735,and is at a distance of 2.3±0.3kpc from the Sun(Kerton et al.1999).

We analyze the ionized component of KR140in§4and make various estimates of the age in§5.The dust,molecular,and atomic components of KR140are examined in§6,§7,and§8, respectively.In§9we discuss star formation in KR140in the context of the Perseus arm and the possible accompanying cluster.

2.THE MULTIW A VELENGTH DATA

KR140was initially cataloged in a1420MHz radio continuum survey of the northern Galactic plane with the E?elsberg100m telescope by Kallas&Reich(1980).At the available resolution of9′KR140was barely resolved,with a reported diameter of about11′.Although KR140was measured in other subsequent single-dish surveys(§4.2)it had never been examined with a radio interferometer to provide the angular resolution necessary for the present study.

The data analyzed here include1420MHz(λ21cm)and408MHz(λ74cm)continuum images from the Dominion Radio Astrophysical Observatory(DRAO)Synthesis Telescope(Roger et al.1973;Veidt et al.1985).An H I21-cm line data cube is also available,with spectral resolution 2.64km s?1and channel spacing1.65km s?1.These DRAO data were obtained during the pilot project of the CGPS(Normandeau et al.1997),and made into a full8?×5?mosaic of the W3/W4/W5star forming regions.

We also made use of a CO(J=1→0,115GHz,λ2.60mm,spectral resolution0.98

km s?1)data cube from the Five Colleges Radio Astronomy Observatory(FCRAO)obtained in the complementary survey described by Heyer et al.(1998).To examine the dust components we processed Infrared Astronomical Satellite(IRAS)data to make HIRES mosaics at12,25,60,and 100μm.Independent processing of the entire Galactic plane data at60and100μm has been released as the IRAS Galaxy Atlas(IGA,Cao et al.1997)and we have completed a complementary project at12and25μm(Mid-Infrared Galaxy Atlas,MIGA,Kerton&Martin2000).

An important feature for subsequent analysis with these data sets is the common relatively high angular resolution achieved.The DRAO synthesis telescope images have a resolution

of1′×1.14′at1420MHz(the408MHz DRAO data have proportionately3.5times lower resolution),and the CO images are at a resolution of50′′(beam sampled).The HIRES images have non-circular beams of size about2′at100μm,1′at60μm,and somewhat less at12and 25μm.For detailed intercomparison,pairs or groups of images have been convolved to the same beam shape.

3.A MODEL FOR THE3-DIMENSIONAL MORPHOLOGY

Analysis of an H II region bene?ts from a knowledge of its three-dimensional structure,but often the observed two-dimensional morphology is too complex to interpret.Fortunately,this is not the case for KR140.Figure2shows the four HIRES images of KR140at12,25,60,and 100μm.As discussed in§6,the100μm image best describes the dust distribution in KR140. The100μm HIRES image(Figure2d)shows that most of the dust in the nebula has been swept into a shell-like structure and that we are observing a region which is close to being circularly (axially)symmetric.In the1420MHz image(Figure3)KR140is also seen to be fairly circular with a central depression.As highlighted by the contour map,there are three high intensity areas

–two“eyes”and a“mouth”.However,where the“nose”would be is a lower intensity area.The simplest interpretation appears to be a three-dimensional structure with a central hole(§5.2.2). We have identi?ed VES735as the exciting star of this H II region(Kerton et al.1999);it lies projected close to this lower intensity area of H II emission,though not coincident with it.

Since we have also classi?ed the exciting star,we can draw further conclusions about the geometrical structure of KR140from a global analysis(§5in Kerton et al.1999).Even after taking into account the emission from dust that was not observed by IRAS,we?nd that the total luminosity of warm dust in the nebula is far less then the bolometric luminosity of the ionizing star VES735and that the covering factor of the dust is only about0.4–0.5(see§6.4). Furthermore,the total radio?ux at1420MHz is lower than expected for an ionization bounded nebula surrounding VES735,with a similar implied covering factor(§4.2.2).This together with geometrical information to be discussed(e.g.,§4.1)implies that KR140is a bowl-shaped H II region(cf.Roger&Irwin1982),rather than a classical Str¨o mgren sphere,a not unexpected morphology when the O star is close to the edge of its parent molecular cloud(Yorke et al. 1983).This model is perhaps the simplest geometry that is consistent with the derived covering factors,but our data cannot rule out other,more complicated geometries,such as a broken shell or shredded and dissipated Str¨o mgren sphere.Observations of a champagne?ow(Tenorio-Tagle 1979),perhaps via Fabry-Perot imaging,could help establish that KR140is truly a blister region.

The circular symmetry observed in the dust shell(and in the ionized gas)implies that we are observing KR140almost face on(i.e.,the opening of the bowl is oriented almost along the line-of-sight).For example,see the simulated radio maps of the R2model by Yorke et al.(1983). The CO data and extinction measurements to VES735both suggest that the molecular cloud is behind the star,with the inferred opening toward us(§7.1).

4.KR140FROM THE PERSPECTIVE OF THE RADIO CONTINUUM

4.1.Size and Structure

The basic radio morphology appears to be that of a limb-brightened hemispherical shell(most of the free-free emission would be from the dense and thin ionization-bounded zone).The last contour shown(at T b=5.0K)in Figure3has a diameter of8.5′which for a distance d=2.3kpc corresponds to a physical diameter of about5.7pc.It is expected that neutral material could extend well beyond the H II region,and indeed there is both dust(§6.2)and CO gas(§7.1)around KR140.The free-free surface brightness falls o?most quickly(the radio contours are more tightly spaced)in the west-northwest part of the nebula.A sharper ionization front would occur in more dense(pre-existing)material.

Other internal detail might be interpreted as variations in distance between the star and the background ionization front,with lower surface brightness features originating from sectors further

from the star;such an interpretation of radio surface brightness has been used to construct a topological map of the ionized surface of the Orion molecular cloud behind the Trapezium(Wen& O’Dell1995).If we use the1420MHz surface brightness at the pixels near the projected position of VES735and the known properties of this O8.5V exciting star,the distance that the star must be from the back of the nebula turns out to be only2.14pc.Together with the above estimate of the diameter across the line of sight,this is consistent with the idea that KR140is similar to a hollowed-out hemispherical bowl.

4.2.Physical Properties

In this section we use the1420MHz data to derive a number of physical properties of the nebula.For ease of reference,Table1summarizes these properties,some of the properties of the exciting star,and other quantities that we later derive from the infrared and CO data.

4.2.1.Emission Measure,Radio Flux,Electron Density,and Ionized Mass

The KR140emission is optically thin at1420MHz with a peak brightness temperature about 10K.For an optically-thin nebula,the speci?c intensity(surface brightness)Iνand brightness temperature T bνfor thermal radiation are(Osterbrock1989):

Iν≡2ν2kT bν/c2=jν n i n e ds≡jνE,(1) where jνis the free-free emissivity,n i is the ion density,n e is the electron density,ds measures distance along the line of sight,and the integral E is called the emission measure.For jνof a pure hydrogen nebula(a good approximation;see§4.2.2for models which include ISM abundances), E=5.77×102T bν(T e/7500K)0.35(ν/1420MHz)2.1cm?6pc.The peak emission measure in the map is about6000cm?6pc and after background subtraction E =2050cm?6pc.After subtracting the local Galactic background we?nd the total?ux(Fν= Iνd?)to be2.35±0.05Jy at 1420MHz.This is in good agreement with the value2.3±0.2Jy reported by Kallas&Reich (1980).Becker,White,&Edwards(1991)detected this H II region in a6cm survey with the NRAO91m.Their source,designated[BWE91]0216+6053,had Fν=https://www.wendangku.net/doc/5010997715.html,ing the same telescope and frequency Taylor&Gregory(1983)and Gregory&Taylor(1986)record this source as GT0216+608in their survey;in their analysis they treated all sources as point sources and so ?nd systematically low?uxes for extended nebulae like this(648mJy for KR140).

In the optically thin limit the theoretical expectation for the spectral index for thermal radiation(de?ned as Iν∝να)isα≈?0.1(Oster1961;Gordon1988).The spectral index was measured between1420and408MHz.First,the distribution across the nebula was determined using L.Higgs’“specmap”routine(Zhang&Higgs1997),which evaluates the spectral index for the spatially variable component of the emission(this approach has the advantage of circumventing

uncertainties in background subtraction in each image);for this procedure,we convolved the 1420MHz image to the lower resolution of the408MHz image.The typical value ofαis

0.02±0.08.We con?rmed this by measuring the total background-subtracted?ux at408MHz as well.Considering there is a10%uncertainty in the?ux calibration in the408MHz pilot project data at the position of KR140(Taylor1999),the derived spectral index is in satisfactory agreement with theory.

Depending on details of the geometry,the appropriate pathlength for estimating the density would be of order the observed radius;with this choice we?nd an rms n e=27cm?3.The mass of radio-observable ionized gas(allowing for He)is163(27cm?3/n e)M⊙.For comparison,the exciting star of KR140has a mass of25M⊙(Kerton et al.1999).

4.2.2.Number of Ionizing Photons

In ionization equilibrium the number of ionizations(where H is the predominant species) that occur each second in the nebula is equal to the number of(H+)recombinations each second, both locally and globally.An ionization-bounded nebula occurs when there is su?cient material to intercept all of the ionizing photons(λ≤912?A),Q(H0),emitted by the star per second (Osterbrock1989).Formally,the spatially-integrated radio emission of an ionization-bounded nebula is an e?ective calorimeter for Q(H0):

fQ(H0)= n e n H+αB dV=αB Fνd2/jν(2) whereαB is the case B hydrogen recombination coe?cient of H+(≈3.3×10?13cm3s?1at 7500K;Storey&Hummer1995),dV is the volume element,and we have used dV=dsd?×d2 and Equation(1)to?nd the right hand side.Here f is the product of a number of correction factors including f ci,the all-important covering factor by an ionization front:if the star is only partially surrounded by gas then a fraction of the ionizing photons will escape and a lower Fνwill result.

Equation2is based on recombination and emissivity for a pure H nebula at a?xed temperature and so f is made up of a number of factors which modify the expected radio continuum?ux for a given Q(H0).We used the spectral synthesis code Cloudy(Ferland et al. 1998)to analyze the observations of KR140in the radio,particularly the correction factors which comprise f.We constructed a number of toy models for di?erential comparisons,using a35,500K blackbody with log[Q(H0)]=48.06and n e=50cm?3;see Table2.The di?erential results are appropriate for any low-density H II region heated by a late O-type star.

The temperature of the nebula enters in the equation through theαB and jνfactors,although the e?ect on the ratio is very https://www.wendangku.net/doc/5010997715.html,paring models1and5we estimate f T=0.99.

The e?ect of the addition of other elements to the nebula is to increase the total radio?ux. This e?ect is primarily due to the presence of He+increasing the e?ective free-free emissivity(He

has a tiny e?ect on the amount of ionization of H;e.g.,Osterbrock1989).Comparing models1 and4,or models2and5we estimate that f He=1.07.

Dust competes with the gas for the absorption of ionizing photons and thus,when present in the nebula,will reduce the observed radio Fνfor a given Q(H0).Comparing models1and2 (with grains)to model3(without grains)we derive f dust=0.78for the typical ISM grains used in model1and f dust=0.93for dark-cloud dust,similar to the dust found in Orion,used in model2. The dark-cloud dust is not as e?cient as typical ISM dust in absorbing short-wavelength photons and so has less of an e?ect on the emergent radio?ux.

The combination of these factors,except for the covering factor,leads to a factor of

f model=0.91±0.08dependin

g upon the grain composition(i.e.,f=0.91f ci).

In our study of the exciting star VES735(Kerton et al.1999)we used the measured radio ?ux at1420MHz(2.35±0.05Jy)to calculate log[fQ(H0)]=48.05±0.11,and then estimated f ci≈0.4?0.5based on the known spectral type of VES735(O8.5V(e);log[Q(H0)]～48.45, Panagia1973).We have not attempted to put any formal error estimate on the covering factor, but given the uncertainties involved,both in the stellar properties and in estimating the various f factors,this result is certainly consistent with the idea of KR140being an open bowl-shaped region.As demonstrated in§6.4,an analysis of the IR data leads to the same conclusion independently.

5.THE AGE OF KR140

The approach we have adopted here is?rst to date the H II region using data we have on the exciting star VES735.With that age in mind,we then investigate the dynamics of the region,the goal being to show that certain scenarios for the evolution of KR140,such as it being a blister region,are at least consistent with the age suggested from the exciting star.This approach is similar to that adopted by Dorland et al.(1986)in their study of the Rosette Nebula.

5.1.Stellar Content

The idea of using the stellar content of a H II region to measure the age of the nebulosity was ?rst attempted by Hjellming(1968).Basically one plots the evolutionary tracks of stars with various masses in the log(L/L⊙)vs.log T eff plane.One obtains T eff from the spectral type of the exciting star and log(L/L⊙)from the radio?ux(much like Q(H0))and then the position in the plane determines an age for the star and thus the H II region.Clearly the e?ectiveness of this technique depends strongly on the quality of the calibration between theoretical and observational quantities as well as the quality of stellar models,which have vastly improved in the thirty years since this technique was introduced.In the early work the primary result was to indicate whether

the H II regions were ionization or density bounded,and that many H II regions required additional, unobserved,sources of ionization.

Here we follow a slightly di?erent technique compared to Hjellming(1968)in order to avoid any uncertainties associated with the covering factor and structure of the nebula.We instead use the absolute magnitude(M V)as a measure of log(L/L⊙),which is possible because we have a good estimate of the distance to the star and the extinction along the line of sight(Kerton et al.1999).Figure4plots the stellar evolution models of Schaerer&de Koter(1997)for a20,25 and40M⊙star along with the value determined for VES735.The observed values for VES735 are consistent with a25M⊙star with an age of a few million years away from the ZAMS(see Figure4),where the age would be only of order105years.Of course,the spectral type alone gives us a simple upper limit to the age of KR140:for an O8.5V star the main sequence lifetime is ～6×106(Chie?et al.1998).

5.2.Dynamical Models:Spherically-symmetric

5.2.1.Str¨o mgren Sphere

The simplest description of a H II region is that of the formation and expansion of a ionized ball of pure hydrogen at constant temperature in an uniform medium of constant density(Str¨o mgren 1939).We summarize this only as a point of departure and contrast.The evolution of a Str¨o mgren sphere starts with a formation phase where the O star ionizes a region of space around it to the radius(R s)given by:

R s= 3Q(H0)

4r o 4/7,(4) where C II is the sound speed in the ionized medium,r o is the initial radius and r i is the radius at time t.With the densities quoted above the H II region will evolve from R s to R obs on a timescale of105years,which is improbably short.Increasing the initial density makes R s smaller and forces

=500cm?3will push the timescale to the pressure expansion stage to be longer.A value of n H

106years.However,this value is not consistent with our observations of the molecular material (§7),so the Str¨o mgren sphere is not an appropriate dynamical model for KR140.

5.2.2.Stellar Winds

In the radio image a local minimum is evident near the position of the central star (see Figure 1in Kerton et al.1999).One interpretation is that this is a wind blown bubble around the O8.5V(e)star.The apparent radius of the central hole is 1.4pc.We used the model of Castor et al.(1975)for the size of a circumstellar shell:

R =28 ˙M 6V 22000

2r o 2/5.(6)

This is based upon mass conservation between material being ejected in the ?ow and molecular material being eroded o?the cloud.One can envisage the evolution of a blister as consisting of three stages:the initial rapid formation stage,a pressure driven expansion stage,and ?nally a blow-out stage.The relative length of time of the latter two stages depends upon the distance of the star from the edge of the cloud and the density structure of the cloud.One very important point is that an O star very close (～1R s )to the edge of a cloud will very quickly develop a covering factor of ～0.5in order 105years and will maintain this covering factor over the lifetime of the O star (Yorke et al.1983,1989).We assume that the star formed very close to the edge of the cloud,thus ignoring the pressure driven expansion https://www.wendangku.net/doc/5010997715.html,ing Equation

6,the age of the region is of order a million years when R s =0.64pc.For the stellar properties of VES 735this requires n H 2～280cm ?3.This is somewhat encouraging as it does not require as vast a di?erence

between the properties of the observed molecular cloud and the putative initial conditions.

6.KR140IN THE INFRARED

6.1.Morphology and Emission Mechanisms

Stars that are forming and evolving in the ISM interact with the interstellar dust component by heating,redistributing,and possibly destroying it.The dust can be heated by at least three distinct mechanisms:direct radiation from the central star,reprocessed radiation from the ionized gas,and di?use radiation from the interstellar radiation?eld.Radiation pressure from the star acts on the dust in the ionized zone(Spitzer1978)which causes the dust(and gas)to be pushed away from the star.Some forms of dust like polycyclic aromatic hydrocarbons(PAHs) are destroyed in intense ultraviolet?elds.A morphological study of the infrared emission from KR140is therefore important to fully understand the energetics and e?ects on the environment.

As mentioned in§2,IRAS scans of KR140at12,25,60,and100μm were processed by the HIRES software to generate maps of about1′resolution.The new beam shapes are somewhat elliptical and so point sources will be visibly stretched;however,the larger scale morphology of the observed dust emission from the nebula will not be greatly a?ected by the asymmetric beam.

The intensity of the dust emission,I dust,at a particular frequency,ν,from a distance increment ds along the line of sight has the form

I dust=N dustπa2Bν(T dust)Qν(a,T dust)ds,(7) where N dust is the number density of grains,Bν(T dust)is the Planck function at a temperature T dust,a is the radius of the particles,and Qν(a,T dust)is the absorption e?ciency factor.

Figure2shows the four HIRES images of KR140,with the12and25μm images convolved to the1420MHz resolution,and overlaid by1420MHz contours(Figure2a&b).The12and 25μm emission is spread well outside the radio contours of KR140(this is taken up in§6.3). Features that are common to all four images in Figure2are the bright arcs on either side of the nebula(easiest to see in Figure2c)that extend outside the radio contours.We interpret this as the limb-brightened warm dust shell around KR140.Note that there could be cooler dust further out around the KR140complex that will not have been detected in the IRAS bands.We account for the energetics of this cooler dust in our models of KR140(§6.4).

6.2.Temperature and Column Density

Equation7can be used to calculate a mean dust temperature for each pixel of IR emission. Taking the intensity ratio for any two frequencies,ν1andν2,yields:

e hν1/kT?1

Iν

ν1 3+β

where T is the T dust of Eq.7,and the(ν2/ν1)3+βfactor is from the frequency-dependent part of Qν(a,T),whereβdepends on the type of dust.The most common components proposed

for interstellar dust are silicate and graphite(or some related carbonaceous material),which,at long wavelengths,haveβ≈2.This value ofβis close to what is generally observed in the ISM (Lagache et al.1998).

The60μm and100μm data were used to calculate the dust temperature map,as they are the bands where classical grain emission dominates.It is likely that non-equilibrium heating of very small dust grains(VSGs)contributes some of the observed60μm?ux and so the derived temperatures are probably slight overestimates of the true grain temperatures(Boulanger et al. 1988).The images were brought to the same resolution and background subtracted and the IPAC analysis program‘cttm’was used to compute a dust temperature map.An azmuthially-averaged radial cut of this map is shown in Figure5.The temperature distribution was sampled every1′in radius and every10?azimuthally.The temperatures in the central region are around31K.Nearer the edge of the H II region,the dust temperatures drop to around28K.However,a line of sight passing near the star also has cooler dust in the background and the foreground,which would lower the apparent temperature observed.

A calculation of dust temperature from?rst principles was done as a check on the empirical values output from‘cttm’.In this calculation,a single silicate or graphite dust grain with a radius of0.1μm(a typical interstellar size,see Kim,Martin&Hendry(1994))was placed at a distance of3.0pc from the center of the nebula,corresponding to a dust grain within the KR140dust shell.The luminosity of the O8.5V(e)exciting star is about105L⊙(Panagia1973).Making use of Planck-averaged absorption factors from Laor and Draine(1993),we?nd that ifτUV≈1,the calculated dust temperature is about28K,in agreement with the temperature deduced empirically. Note,however,that this calculation did not take into account the recombination-line photons emitted by the ionized gas.Furthermore,some of the free-free photons and collisionally-excited cooling lines are emitted not in the ultraviolet,but in the optical(Osterbrock1989),where the dust absorption e?ciency is somewhat lower.

The optical depthτνfrom dust is de?ned to be

τν= N dustπa2Qν(a,T dust)ds,(9) and can be calculated by dividing Iν(Equation7)by the Planck function Bν(T),assuming a constant T along the line of sight,adopted from the temperature map.Figure6displays a dust optical depth map at100μm.The values range from about0.0004in the middle of the nebula to about0.002in the bright northwest rim.These values show that the dust is transparent to its own100μm emission.The minimum in the center of the nebula and the ring-shaped appearance is most simply interpreted as limb brightening in a thick shell of dust with the highest column density along the northwest rim.The latter is the same region of the nebula where the1420MHz radio contours fall o?most steeply(§4.1),and where there is no CO emission(§7.1).

The optical depth in the ultravioletτUV can be gauged using extinction curves from the literature.We are interested in the radial as opposed to line of sight optical depth.Judging the thickness of the shell from one of the arcs in Figure6gives a path length of about7×1018cm. If we assume the dust is associated with gas at a molecular density of100cm?3,the molecular column density of this gas is about7×1020cm?https://www.wendangku.net/doc/5010997715.html,ing the extinction curves of Kim,Martin& Hendry(1994),we?nd thatτν≈2.8at1100?A(far in the ultraviolet),whereas at5500?A(in the optical),τν≈0.6.The estimated optical depths show that the dust in the arcs has a high enough radial optical depth to absorb most of the incident ultraviolet photons.However,since KR140 has a covering factor of0.4–0.5(§4.2.2and§6.4),the total infrared luminosity of the nebula will be correspondingly less than the bolometric luminosity of VES735(§6.4).

6.3.The12and25μm Emission

It has been known for over a decade that there is excess emission within the12μm IRAS passband(e.g.,Boulanger,Baud&van Albada1985).Onaka et al.(1996)found that more than 70%of the12μm di?use interstellar emission detected by IRAS is emitted in spectral features attributed to polycyclic aromatic hydrocarbons(PAHs;L′e ger&Puget1984;Allamamdola, Tielens&Barker1985).The12μm image of KR140is very instructive.There is very little 12μm?ux within the radio contours of KR140,implying that PAHs are destroyed there in the intense ultraviolet radiation?eld.However,there is a large amount of di?use?ux outside the radio contours,especially to the west(Figure2a).The PAH emission is a tracer of the photodissociation region around a nebula(Giard et al.1994;Bregman et al.1995;Fig.10).

Unique among the HIRES images of KR140,the25μm image(Fig.2b;Fig.7)shows a bright spot near the center of the nebula,on a transition between a radio peak(the“left eye”)and the deepest depression.The IRAS Point Source Catalog(Joint IRAS Science Working Group1988) lists this feature as IRAS02165+6053,and it also has been identi?ed with VES735(Bidelman 1988).Figure7shows that IRAS02165+6053and VES735are practically coincident.Dust closer to the star would tend to be warmer,contributing to the spot if not pushed away by radiation pressure.

The entire nebula looks hotter and more extended at25μm than it would be for equilibrium emission from normal-sized grains.Most of the25μm emission is probably contaminated by non-equilibrium emission from very small grains(VSGs;Sellgren1984)which have absorbed a UV photon and have had their temperatures instantaneously rise to～103K.Although these small grains make up a tiny fraction of the mass in the dust distribution,they make up a good fraction of the number distribution and absorb a signi?cant fraction of the near-ultraviolet radiation.This VSG emission is therefore a major source of uncertainty when analyzing the25μm image.3Note

that the“temperature-spiking”phenomenon also occurs outside the ionized zone,far from the exciting star.

6.4.Infrared Models

H II regions are some of the most luminous objects in the Galaxy when observed in the infrared,especially at wavelengths longer than60μm.In fact,if the dust shell covers4πsteradians around the exciting star,the infrared luminosity should be a good measure of the star’s bolometric luminosity.To account for the true extent of the dust,we can de?ne a covering factor of dust(f cd). Often it is simply assumed that L ir=L bol;however,this is not correct even for f cd=1.Some of the radiation from the star and nebula is at long enough wavelengths to avoid being absorbed by the https://www.wendangku.net/doc/5010997715.html,ing Cloudy we found that for models with f cd=1,log(L ir/L bol)～?0.1.

To calculate the integrated infrared?ux from KR140,the background in each of the four HIRES images was?tted and subtracted,and the total?ux from the nebula was measured in each band.In order to estimate L bol using these data,we used Cloudy to simulate the emission from large classical grains,roughly matching the?uxes at60and100μm.This approach allows us to account for emission from grains with a range of temperatures and thus unobserved emission at long wavelengths.First,we constructed models with f cd=1.We?nd log(L F IR)=37.90. Since the resulting spectrum misses most of the observed12and25μm?ux(which is caused by temperature spiking of VSG’s and the excitation of PAH molecules)we converted the observed12 and25μm?uxes to luminosities using tophat approximations to the IRAS passbands(Emerson 1988)and added the results to L F IR.With this addition we?nd log(L IR)=38.00.Correcting from L ir to L bol we?nd log(L bol/L⊙)=4.52.This is signi?cantly below what would be expected for any late O main sequence star(e.g.,Panagia1973).We interpret this low apparent L bol as being due to a covering factor of0.4–0.5.

We know that VES735is an O8.5V(e)star.Figure8demonstrates that one can reproduce the observed L ir using the appropriate stellar parameters and a covering factor of about0.5.A single temperature(T=28.25K)ν2Bνspectrum is shown for comparison;note the wider model curve caused by the range of dust temperatures contributing to the emission.

Using an alternative calibration of stellar parameters which includes wind-blanketed models (Kerton1999),we?nd f cd～0.5and f ci～0.7.These results are still consistent with a blister model for KR140.

7.KR140IN A MOLECULAR CLOUD

7.1.CO Signature

A slit spectrum of VES735allowed us to measure the radial velocity of the nebular Hαline within30′′of VES735to be?46±2.1km s?1with respect to the Local Standard of Rest[LSR] (Kerton et al.1999).The di?erential radial velocity between VES735and the nebular line was measured to be+2.0±2.2km s?1.

An expected e?ect of the evolution of the H II region is photodissociation of molecular gas both inside the H II region and in the immediate surrounding interstellar medium(ISM).The CO data traces the molecular gas content of the ISM,so an examination of the CO data cube ought to reveal a lack of emission near the ionized gas velocity of KR140.Indeed,a distinct CO hole was found within the radio contours of KR140over the velocity range?45.53km s?1to

?47.16km s?1(LSR),which corresponds to three channels of width0.813km s?1in the CO cube. At more negative velocities,CO emission from the parent molecular cloud?lls in the1420MHz contours.

Figure9shows the sum of these three channels overlayed with the1420MHz continuum contours.A well de?ned ring structure is clearly seen.Interestingly,the ring does not extend all the way around the nebula:there is no CO emission in the north-west.This is the same area where the1420MHz contours are falling o?more sharply(§4.1),there is a bright infrared arc (§6.1),and a bright H I feature(§8).The relationship between12μm emission,which is a good tracer of the PDR,and the CO emission is shown in Figure10.The north-west peak in the12μm emission corresponds to the region where there is no CO emission.To the east,as one moves away from the H II region the12μm peak occurs?rst followed by the peak in the CO emission.

7.2.Density and Mass of the Molecular Cloud

To estimate the mass of the molecular cloud we integrated the CO cube over the whole velocity range of the cloud(?45.5km s?1to?52.8km s?1),and measured the total surface brightness of the molecular cloud.We can use the empirical X factor to convert from CO surface brightness to molecular hydrogen column density.We have used X=(1.9±0.5)×1020cm?2(K km s?1)?1 (Strong&Mattox1996),where the quoted uncertainty is to take into account the range

of X values measured for a variety of di?erent clouds,to obtain a mean column density of

N(H2)=2.2×1021cm?2.Assuming a path length of7pc(the north-south radial extent)the

=100cm?3,typical of a giant molecular cloud(Blitz density of the molecular cloud is roughly n H

1993).

By integrating the column density over the face of the cloud,we estimate the mass of the cloud to be(4.4±1.6)×103M⊙.This mass includes a He correction factor of1.36,and the error bar includes the distance uncertainty.The CO images show that the parent molecular cloud of KR140has been greatly disrupted by the nebula and the exciting star VES735and so this mass

will be an underestimate of the cloud’s initial mass.

7.2.1.Estimating the Mass of the Original Molecular Cloud

One way to estimate the mass of the original molecular cloud would be to assume that

KR140is indeed a blister H II region as our data suggests.Yorke et al.(1989)have modeled blister H II regions with one exciting O star and found that mass loss rate of material through the blister is3?5×10?3M⊙yr?1.However,these authors modeled an O6star in a cloud with

n H

=500cm?3,which is not an accurate description of either VES735or its parental cloud. Unfortunately,Yorke et al.(1989)give no indication of how to scale their mass loss rate for

di?erent values of Q(H0)or n H

.A analytic estimate of the amount of mass lost from a blister H II region is given by Whitworth(1979).His Equation(41)has the desirable property that it agrees with the results of Yorke et al.(1989)for the values used in their models;therefore,it might have

the correct scaling for Q(H0)and n H

.However,his calculations were based on a cylindrical geometry which is not a very realistic model for a more bowl-shaped spherically symmetric region such as KR140.

In§5.3,we made use of the blister evolution formula given by Franco et al.(1994).These authors considered a simple symmetric blister region and were also able to estimate the cloud evaporation rate:

˙M≈πR2

S m p2n H

C II 1+5C II t

7.3.Interpretation of the Velocity

Stepping through the velocity channels in the CO data cube near the velocity of KR140 shows that at higher(least negative)velocities(～?40km s?1)there is no emission,then the CO emission comes in at the southern end of the H II region and spreads northward,before?lling in the1420MHz contours at a velocity of?48.78km s?1.Perhaps the best way to illustrate these data is by examining the cube in velocity-latitude space.Figure11shows such a CO image of KR140,averaged over the longitude range133.366?to133.477?.Note the“hole”at the velocity of the ionized gas and VES735.

The question arises as to the relative radial position of VES735and KR140with respect to the molecular gas.In general,it is di?cult to use only the CO radial velocity information to determine absolute distances to clouds or GMCs.However,in this direction,Galactic rotation causes radial velocity to become more negative with increasing distance.On the face of it,that would place the molecular cloud on the far side of the star and nebula.However,this is possibly too naive an interpretation because if the cloud is only10s of pc in radial extent(like its dimension in the plane of the sky),then the average shear would be too small to explain the large range in velocity(unless there were a large enhancement from a density wave).

To address the question of relative position,we made use of both the H I and CO data cubes to make an estimate of A V produced by gas with velocities out to?45km s?1.As in§7.2 we used the conversion factor X=(1.9±0.5)×1020cm?2(K km s?1)?1to convert I(CO)to N(H2).We integrated both data cubes to obtain the atomic column densities.At the position of VES735we found N(H I)=4.5×1021cm?2assuming the emission is optically thin,and N(H2) =3.1×1021cm?2,the latter mostly from local rather than Perseus arm gas.Summing the contributions of atomic and molecular hydrogen we obtain the total hydrogen column density,N H =2N(H2)+N(H I)=1.1×1022cm?2.The total visual extinction is computed using the standard conversion factor,A V=5.3×10?22mag cm?2(Bohlin,Savage&Drake1978).At the position of VES735we obtain A V=5.7±0.9,where most of the uncertainty comes from the uncertainty in the X factor(using the X factors of Digel et al.1996gives a lower A V).This value of A V compares favorably with the values around5.7derived by Kerton et al.(1999)using a number of methods(e.g.,B?V colour,Hαemission measure,DIBs).However,if the cloud’s molecular column density(summed over?45to?53km s?1)of N(H2)=2.2×1021cm?2(§7.2)is included in the extinction calculation,then the A V rises to8.0,which is quite inconsistent with the observed A V to VES735.We also made maps of the predicted extinction over the surface of the nebula for comparison with the extinction map derived by comparing Hαsurface brightness with radio emission(Kerton et al.1999).Again,inclusion of the extinction from gas in the molecular cloud produces more than two magnitudes too much extinction.This is strong evidence that both VES735and the ionized gas of KR140lie on the near side of the molecular cloud gas which has velocities?45to?53km s?1.

The extinction measurements combined with Figure11suggest that KR140could be a

blister H II region on the near side of the molecular cloud.This simple geometrical interpretation of KR140runs into di?culty if a systematic champagne?ow has developed with gas?owing away from the parent cloud at up to the sound speed of～10km s?1.At face value,our Hαvelocity implies that the ionized gas is redshifted with respect to the molecular gas,which means the H II region should be on the far side of the molecular cloud.However,it can be noted that Hαis considered to be a poor line for estimating the velocity?eld of champagne?ows(Israel1978;Yorke et al.1984).Modeling of line pro?les in champagne?ows also shows that the velocity?eld can be quite random and of low amplitude,depending on geometry(Yorke et al.1984).Recall that the geometry here is certainly not plane-parallel and probably more like a broken shell.Furthermore, our measurement sampled only gas within a projected distance of30′′of VES735.Fabry-Perot data of the whole nebula in a line other than a hydrogen line would be a good way to obtain a better picture of the velocity structure of the ionized gas in KR140.

The question then arises as to the origin of the radial velocity spread within the molecular cloud.If this is a single,gravitationally bound cloud,then this is just the virial velocity,and its mass can be estimated from

5Rσ2

M vir=

atomic gas.Since the velocity of the neutral atomic gas can be many times the typical channel width in the H I data sets velocity becomes a much poorer proxy for physical distance than in a CO data set,which is probing a species with lower velocity dispersion(heavier,cooler,and less turbulent).The situation is especially problematic for H II regions in the Galactic plane where one has to look through a large column of atomic gas towards the H II region.For KR140we are looking through the local arm and part way into the Perseus arm of the Galaxy.A preliminary reconnaissance of the H I cube con?rmed that the complex H I emission structure along the line of sight makes seeing any H I signal associated directly with KR140extremely di?cult.Nevertheless, some simple processing of the H I data cube does bring out some features associated with the region.

In order to exclude local H I emission(which is assumed to have a relatively smooth spatial structure)and to enhance the dynamic range of the resulting channel maps we constructed a median-subtracted data cube(Joncas et al.1992;Joncas et al.1985).In this technique a median spectrum is calculated for the data cube and then subtracted from each spectrum making up the cube.The resulting channel maps thus can contain both negative and positive values indicating deviations relative to the median base level.

Figure12shows channel maps of the median-subtracted cube over the velocity range?43.40 to?54.95km?1bracketing the ionized gas velocity of?46±2.0km s?1(Kerton et al.1999).We could not detect any features de?nitely associated with KR140in the channel maps outside of this velocity range.

Examining these maps we note the following three features.First,there is a noticeable de?cit of H I seen in velocity channels?46.70to?51.65km s?1outside of the H II region.There is excellent positional agreement between this de?cit in H I and the observed position of the CO emission.The de?cit can be simply interpreted as being caused by a lack of H I emission in the molecular material surrounding KR140.Second,we also see a drop in the H I emission occurring within the H II region in the?43.40and?45.05km s?1channels.This de?cit is most likely associated with the ionized gas in KR140as suggested by their spatial correspondence.Figure13presents spatially averaged H I and CO spectra for an area just outside of the H II region to the north-east and the area inside the H II region.The anticorrelation between H I emission and the presence of CO and H II is evident.Third,there is an enhancement of H I emission at(133.36?,0.1583?),seen best in the velocity channel?46.70km s?1.This could be low velocity dispersion material associated with the PDR.Atomic material in the PDR is expected to have a low velocity and thus should be seen in channels corresponding to CO emission.

9.THE ENVIRONMENT OF KR140

9.1.Spontaneous Massive Star Formation

In the context of the Perseus Arm star formation activity,the KR140complex seems to be unique.Figure1shows that KR140is isolated from the massive and violent star formation that is ongoing around it.This isolation is evidence to us that KR140is an example of massive star formation in our Galaxy that is untriggered,at least in the sense used in the context of sequential star formation.In none of our data sets does there appear any evidence for a trigger of the star formation in KR140.The exciting cluster of W4,OCl352,is about60pc away from KR140for a cluster distance of2.35kpc(Massey,Johnson&DeGioia-Eastwood1995).For a sound speed of0.6km s?1(isothermal speed in H2at100K),the time for a signal to reach KR140would be about90Myr,much greater than our estimated age of KR140of a few million years or the age of OCl352.From W3,a signal would take about60Myr,but W3is itself much less than 105years old(Kawamura&Masson1998).The supernova remnant HB3is also in this complex (Normandeau et al.1997),but from Figure1,the edge of the remnant is nowhere near to KR140. Thus,KR140does not seem to be triggered by a neighboring H II region or a nearby supernova remnant,unless the impulse came along the line of sight.This conclusion is complemented by the overall smoothness of the KR140nebula;while the nebula does show density inhomogeneities,it is not far removed from a circular shape.Thus,any perturbation that might have triggered the star formation within KR140must have been a large scale phenomenon with a characteristic scale of～10pc.This kind of triggering might be more consistent with triggering via a spiral density wave(e.g.,Elmegreen1994,1995)or by colliding molecular clouds.Of course,we cannot rule out those kinds of triggers,but the observational evidence would be di?cult to?nd.

It is interesting to contrast KR140to another star forming region that has been studied with multiwavelength data,the Gemini OB1molecular complex(Carpenter,Snell&Schloerb 1995a,1995b).Within the molecular complex,these authors?nd young star clusters(from near infrared data)and a number of dense cores(as identi?ed by CS observations)associated with IRAS point sources.Carpenter et al.suggest that the arc-shaped morphologies of these cores have been formed by swept-up gas from expanding H II regions,and that they would form the massive star clusters in the region.The other lower infrared luminosity sources(most likely belonging to lower mass cores)in the Gem OB1complex are not found to be correlated with any arc-shaped structures or?laments in the molecular gas,and they are not adjacent to any H II regions;in fact they are spread almost randomly around the complex.Carpenter et al.therefore conclude that induced star formation is the prominent mode of formation for massive stars in the Gem OB1 molecular complex.

If KR140is indeed untriggered,then it seems to be an unusual form of spontaneous star formation since current observations suggest that the isolated mode of star formation is generally associated with low mass star forming regions as seen in Gem OB1or even in Taurus.

9.2.IRAS Point Sources and Protostars

As has been seen in other sites in the Galaxy,an H II region can initiate star formation via the expansion of its ionization front(e.g.,Elmegreen&Lada1977).Other stars might also have formed spontaneously.Evidence for other star forming regions within KR140may be sought by examining the IRAS point source catalog.In addition to the IRAS point source02165+6053that is cross-referenced to VES735,there are six other IRAS point sources in the area in and around KR140(see Table3).The crosses in Figure14show the positions of these six IRAS point sources. The circled crosses are the sources discussed here that seem most likely to be associated with the KR140complex.

The source IRAS02160+6057identi?ed with the north-west dust arc has been the subject of two molecular line investigations.Wouterloot&Brand(1989)identi?ed it as a potential star-forming area via its IRAS colors(see their paper for the exact selection criteria),and examined it(along with about1300other IRAS point sources)for CO emission.They found a CO feature in that direction at a velocity of?49.7km s?1(LSR),which corresponds to CO in the associated background molecular cloud,and assigned it the catalog number[WB89]417.This point source was then observed in an H2O line by Wouterloot,Brand&Fiegle(1993),but they were unable to detect any emission.This line of sight has one one of the highest column densities in the KR140H II region,and so it is possible that a protostar could be forming there as a result of the expansion of the H II region.However,examination of the HIRES images shows no point-like features in the dust arc,and none are found in follow-up submillimetre observations with SCUBA (Kerton et al.2000).Therefore,IRAS02160+6057is most likely simply part of the KR140dust shell.Likewise,IRAS02168+6052appears to just be part of the eastern side of the dust shell;it has not been the subject of any molecular line observations.

The source IRAS02157+6053also seems to have been an identi?cation of some part of the dust shell.In the submillimetre there is more structure in this region;this object appears to be a molecular core rather than a protostar(Kerton et al.2000).

The point source IRAS02171+6058is identi?ed with the dust feature to the north of the KR140complex.It was included in a CS(2–1)survey by Bronfman,Nyman&May(1996)of IRAS point sources that have colors characteristic of ultracompact H II regions;however,they were unable to make a detection.Lyder&Galt(1997)observed this source along with other ultracompact H II region candidates in a search for methanol maser emission.Again,they were unable to detect any maser emission from IRAS02171+6058.These non-detections do not rule out the possibility that the source is a protostar.In fact there is a submillimetre detection and all evidence seems consistent with a B5V star(Kerton et al.2000).It is outside the radio contours of H II region and is therefore di?cult to interpret as being triggered.

A photographic survey of the W3and W4region turned up Bright InfraRed Stars(BIRS; Elmegreen1980)which are brighter in I band than in R band.There are?ve of these stars in the KR140region of the sky.Table4gives their coordinates along with their R and I magnitudes and

开关柜及相应电气元器件知识汇总

开关柜及相应电气元器件知识汇总文稿归稿存档编号：[KKUY-KKIO69-OTM243-OLUI129-G00I-FDQS58-

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抽油机维护保养操作规程

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4、检查各润滑点及补充润滑剂。 5、检查胶带的张紧是否适当。在胶带中点用一只手压0.15~~0.20kN力，胶带下垂30~~50mm为合适。V型胶带应保证电动机胶带轮与输入轴胶带轮端面不共平面不大于1mm（用拉线目测）。 6、检查整机噪声。在抽油机运行时，测其噪声，测点分布在减速器输出轴两端，距轴端面1m 远，距抽油机底面1.5m高处，用声级计A声级（慢档）测取算术平均值应符合下表的规定。 7、检查减速器轴承温度。轴承最高温度不得超过 70℃。 8、检测减速器油池温度。油池最高温度不得超过 70℃。 9、检查减速器油池润滑剂质量。如有以下情况应立即更新润滑剂。（a）、减速器内部各表面是否有沉积物。

（b）、油是否被砂子、金属微粒等杂质污染。（c）、油是否乳化。 10、检查齿轮接触斑点、有无擦伤、点蚀或裂纹等损坏情况。 11、检查轮辐、轮缘有无裂纹。 12、检查抽油机所有焊接件本体有无永久变形、开裂现象和焊缝有无疲劳裂纹。 13、检测减速器齿面磨损情况。测量齿轮侧隙是否超标。 14、检查抽油机悬点投影、剪刀差。 15、检查基础稳定性，测量支架振幅指标不得超过要求。 16、更换影响抽油机性能的磨损零部件。

油田抽油机日常维护保养

龙源期刊网 https://www.wendangku.net/doc/5010997715.html, 油田抽油机日常维护保养作者：张骅来源：《科学与财富》2020年第01期摘要：为了提升油田抽油机工程的效率，本文结合实际，在分析抽油机井现场异常状况判断内容的同时，对它的日常维护保养方案进行总结分析。希望相关方案的提出，可以提升油田抽油机的工作效率。关键词：网油田;抽油机;日常维护;保养前言在井下进行作业时，对于施工人员来说，抽油机井的运行参数需要进行严格的检查，并对参数进行核对，保证参数的准确性，通过这样的方式，能够为抽油机井提供更加稳定的安全保障。在井下作业的过程中，许多步骤以及工具都需要进行检查工作，除了机井的油管需要进行分析工作以外，抽油泵也需要对运营状况进行检测工作。抽油机井需要进行维护工作，因此，只有对抽油机井可能出现的故障问题进行深入的理解，才能够保证各项维护工作能够达到预期的目的。 1抽油机井现场异常状况判断 1.1当前抽油机井的分类当前在我国，油田事业正在不断发展的过程中，实际的运行中，主要是通过抽油机和井下油杆泵之间的结合进行的。当前在开采油田时，抽油机与井下油杆泵相结合的方式是最常见的作业方式，除此之外，根据抽油杆的不同，可以进行分类，例如普通钢杆井以及连续杆井等等。抽油机的不同，会产生不同的效果，然而所有抽油机的的原理大致是相同的，油杆通过上下运动进行原油的采取。 1.2抽油机井异常状况分析在对抽油机进行维护时，需要对部分情况进行严格的检查工作，例如电流、流量以及压力的变化等等。对电流情况进行分析，能够对抽油机进行分析，如果出现故障问题能够及时进行故障的分析工作。对流量情况进行分析，能够对抽油机的运行工作进行检查，如果出现流量下降的情况，则需要及时进行处理。 1.3抽油机井抽油泵故障判断

LTE常用参数详解

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1.2、Pa（不含CRS的符号上PDSCH的RE功率与CRS 的RE功率比）功能含义：不含CRS的符号上PDSCH的RE功率与CRS的RE功率比对网络质量的影响：在CRS功率一定的情况下，增大该参数会增大数据RE功率取值建议：-3 1.3、PreambleInitialReceivedTargetPower（初始接收目标功率(dBm)）功能含义：表示当PRACH前导格式为格式0时，eNB期望的目标信号功率水平，由广播消息下发。对网络质量的影响：该参数的设置和调整需要结合实际系统中的测量来进行。该参数设置的偏高，会增加本小区的吞吐量，但是会降低整网的吞吐量；设置偏低，降低对邻区的干扰，导致本小区的吞吐量的降低，提高整网吞吐量。取值建议：-100dBm~-104dBm 1.4、PreambleTransMax（前导码最大传输次数）功能含义：该参数表示前导传送最大次数。对网络质量的影响：最大传输次数设置的越大，随机接入的成功率越高，但是会增加对邻区的干扰；最大传输次数设置的越小，存在上行干扰的场景随机接入的成功率会降低，但是会减小对邻区的干扰取值建议：n8，n10

一般电气元器件介绍

一般电气元器件介绍一、变压器变压器是变换交流电压、电流和阻抗的器件，当初级线圈中通有交流电流时，铁芯（或磁芯）中便产生交流磁通，使次级线圈中感应出电压（或电流）。变压器由铁芯（或磁芯）和线圈组成，线圈有两个或两个以上的绕组，其中接电源的绕组叫初级线圈（一次侧），其余的绕组叫次级线圈（二次侧）。 1、电源变压器的特性参数额定功率：在规定的频率和电压下，变压器能长期工作，而不超过规定温升的输出功率。额定电压：指在变压器的线圈上所允许施加的电压，工作时不得大于规定值。还有空载电流、空载损耗、效率、绝缘电阻等参数不作详细讨论 2、车间一般使用变压器车间采用的变压器的类型比较多，主要有单相380v 变成单相100~110V （基本每台设备上都有该类变压器）、单相380v 变成单相200~220V （主要使用于C 线国产设备）、三相380v 变成三相200V （使用于一些采用200V 三相电机的设备，如曲轴AB 线无心磨床、偏心磨床、活塞倒角机、选配铆接机等） 3、接线方法，下图：输入电源输出电源可选择输入电压

二、开关电源在通常的设备中采用交流电，但在大部分电子设备中都需要使用直流稳压电源。 1、工作原理：下图所示为直流稳压电源的组成框图，实际上它就是表示把交流电转换为直流电的过程。图中各环节的功能如下：电源变压器：将标准交流电源电压变换为符合整流需要的交流电压（一般为降压）。整流电路：利用二极管等整流器件的单向导电特性，将交流电压变换为单向脉动电压。滤波电路：减小整流电压的脉动程度，以适合稳压电路的需要。稳压电路：在交流电源电压波动或负载变化时，使直流输出电压稳定。 2、一般输入电源为交流单相100V （220V ），主要看开关电源标识，也有同时兼容两种电压的开关电源（需作连接跳线处理），而输出电源为直流24V 。替换时主要指标：输入电压、输出电压、定额电流 3、线方法，下图： +24V 输出 0V 输出地线 AC100V 输入电源指示灯输出调整

抽油机维护保养规定

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DATAGUARD配置参数详细解释

DATAGUARD配置参数详细解释 DB_NAME 只需注意DataGuard的主备各节点instance使用相同的db_name即可。推荐与service_name一致。 DB_UNIQUE_NAME Primary与Standby端数据库的唯一名字，设定后不可再更改。注意: 如果主备db_unique_name不一样，需要与LOG_ARCHIVE_CONFIG配合使用 db_unique_name并未规定需要与数据库service_name一致，可以自定义任意名称。 LOG_ARCHIVE_CONFIG 列出主备库上的DB_UNIQUE_NAME 参数。默认情况下，定义该参数能确保主备库数据库能够互相识别对方Primary与Standby端的db_unique_name不一致时如在主备库db_unique_name不一致的情况下未配置LOG_ARCHIVE_CONFIG则会出现如下报错 ORA-16057: DGID from server not in Data Guard configuration 原因:主库没有设置参数log_archive_config 解决方法*.log_archive_config='dg_config=( Primary, Standby)' alter system set log_archive_config='dg_config=( Primary, Standby)' scope=both; Primary与Standby端的db_unique_name一致时

LOG_ARCHIVE_DEST_1 本地归档路径。Primary与Standby需要定义各自的online redo log的归档地址，以系统实际的存放路径为准。格式如下： Primary Site: *.LOG_ARCHIVE_DEST_1='LOCATION=/arch/ VALID_FOR=(ALL_LOGFILES,ALL_ROLES) ' Standby Site: *.LOG_ARCHIVE_DEST_1='LOCATION=/stdby/ VALID_FOR=(ALL_LOGFILES,ALL_ROLES) ' 注意：在LOG_ARCHIVE_DEST_n设置DB_UNIQUE_NAME表示该参数在DB_UNIQUE_NAME指定的数据库上生效，设置为本地的db_unique_name。以priamry端为例，格式如下： *.LOG_ARCHIVE_DEST_1='LOCATION=/archivelog/ VALID_FOR=(ALL_LOGFILES, ALL_ROLES) DB_UNIQUE_NAME=Primary' 这样配置的意义为：在数据库Primary上log_archive_dest_1对主备库上的联机日志都有效，这里的 db_unique_name可以省略 LOG_ARCHIVE_DEST_2 该参数仅当数据库角色为primary时生效，指定primary归档redo log到该参数定义的standby database上。 log_archive_dest_2可以说是dataguard上最重要的参数之一，它定义了redo log的传输方式(sync or async)以及传输目标(即standby apply node)，直接决定了dataguard的数据保护级别。格式如下： Primary Site: *.LOG_ARCHIVE_DEST_2='SERVICE=DR2 lgwr async VALID_FOR=(ONLINE_LOGFILES, PRIMARY_ROLE) ' Standby Site: (switch over后生效) *.LOG_ARCHIVE_DEST_2='SERVICE=DR1 lgwr async VALID_FOR=(ONLINE_LOGFILES, PRIMARY_ROLE) ' 注意： LOG_ARCHIVE_DEST_2参数里定义的service值，比如DR1，是tnsnames.ora文件里定义的Oracle Net名称。

煤油安全技术说明书(新国标格式)

安全技术说明书煤油 ********有限责任公司二〇一四年五月

化学品安全技术说明书产品名称：煤油按GB/T16483-2008、GB/T17519-2013编制修订日期： 2014年5月20日SDS编号： **- 煤油最初编制日期： 2014年5月20日版本： 1.1 第1部分化学品及企业标识化学品中文名：煤油化学品英文名：kerosene 企业名称：********有限责任公司企业地址：****************** 邮编：****************** 传真：****************** 联系电话：****************** 电子邮件地址：************************************ 企业应急电话：****************** 产品推荐及限制用途：可用作燃料、溶剂、洗涤剂、工艺用油。第2部分危险性概述紧急情况概述： GHS危险性类别：易燃液体类别3 急性毒性——经口类别5 急性毒性——经皮类别5

急性毒性——吸入类别5 标签要素：象形图：警示词：警告危险性说明：易燃液体; 引起眼睛刺激; 吞咽可能有害; 皮肤接触可能有害; 吸入可能有害。防范说明： ·预防措施：在得到专门指导后操作。在未了解所有安全措施之前，且勿操作。远离火种、热源。使用不产生火花的工具作业。采取防止静电措施，容器和接收设备接地、连接。使用防爆型的通风系统和设备。密闭操作，全面通风。避免吸入蒸气（或雾）。戴橡胶手套，化学安全防护眼镜，穿防静电工作服。可能接触其蒸气时，戴过滤式防毒面具（半面罩）。紧急事态抢救或撤离时，建议佩戴空气呼吸器。妊娠、哺乳期间避免接触。作业场所不得进食、饮水、吸烟。操作后彻底清洗身体接触部位。污染的工作服不得带出工作场所。 ·事故响应：发生火灾时，尽可能将容器从火场移至空旷处。喷水保持火场容器冷却，直至灭火结束。处在火场中的容器若已变色或从安全泄压装置中产生声音，必须马上撤离。用泡沫、干粉、二氧化碳灭火。 ·安全储存：储存于阴凉、通风的库房。远离火种、热源。库温不宜超过30℃。保持容器密封。应与氧化剂、卤素、强酸碱分开存放，切忌混储。采用防爆型照明、通风设施。禁止使用易产生火花的机械设备和工具。储区应备有泄漏应急处理设备和合适的收容材料。 ·废弃处置：本品或其容器采用焚烧法处置。处置前应参阅国家和地方法规物理和化学危险：易燃，其蒸气与空气可形成爆炸性混合物，遇明火、高热能引起燃烧爆炸。与氧化剂接触发生化学反应或引起燃烧。在火场中，受热的容器有爆炸的危险。蒸气比空气重，沿地面扩散并易积存于低洼处，遇火源会着火回燃。

抽油机保养规程教学内容

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3、按停止按钮，刹住抽油机刹车（停在便于操作的位置），拉下抽油机总开关和闸门。 4、紧固减速箱、底座、中轴承、平衡块、电机等部固定螺丝。 5、打开减速箱检视孔，松开刹车，盘动皮带轮，检查齿轮啮合情况。发现问题及时向部门反馈，以便处理。 6、按各型抽油机的润滑要求和规定进行润滑作业。 7、清洗减速箱呼吸阀。 8、检查、校正电动机皮带轮与减速器皮带轮的平行度，检查皮带松紧程度，不合适进行调整；皮带损坏要及时更换。 9、检查、紧固曲柄销螺母，驻头销连接销，中央轴承座螺栓等关键部件。 10、检查、紧固、调整其他各部件连接。 11、检查刹车是否灵活好用，必要时进行调整。 12、对中轴轴承、尾轴轴承、曲柄销子轴承、驴头固定销子等处加注润滑脂。 13、检查悬绳、有起刺、断股现象应更换；检查后悬挂配重完好。 14、对电气装置进行一级保养。检查电器设备，绝缘应符合规定要求，有接地线，各触点接触完好。 15、清洁抽油机表面油污、泥土。 16、检查抽油机护拦是否完好。抽油机二级保养

电脑参数相关配置

CPU Intel Pentium E5300/盒装450 主板：捷波蓝光X-BLUE P43 399元内存宇瞻2GB DDR3 1333 300 显卡:昂达9600GSO 384M 400元电源航嘉冷静王钻石2.3版本1 ￥200 硬盘希捷500GB 7200.12 ￥350 光驱三星TS-H352 115 显示器:AOC 919SW 760元处理器(CPU) AMD Athlon(速龙) II X2 245 430(原装）主板映泰A770 A2G+ 420 内存金士顿2GB DDR2 800 245 显卡铭瑄极光9500GT高清版-TC512M 375 电源航嘉冷静王钻石2.3版本1 ￥200 硬盘希捷250GB 7200.12 16M（串口/散）1234 ￥280光驱东芝-三星DVD-ROM TS-H352D DVD光驱120 LCD 三星943NW ￥870 硬盘希捷160 300 CPU Intel 奔腾双核E5200(盒) 1 ￥420 主板捷波XBLUE-P43 1 ￥399 内存三星2GB DDR2 800(金条) 1234 ￥156 硬盘希捷250GB 7200.12 16M（串口/散）1234 ￥280 显卡双敏无极2 9600GT金牛版1234 ￥599 LCD HKC S988A 1 ￥690 机箱大水牛A0707（空箱）1 ￥105 电源航嘉冷静王钻石2.3版本1 ￥228 键鼠装灵标经济套装1 ￥35 合计金额：2967 元处理器(CPU) AMD Athlon(速龙) II X2 245 430(原装）主板华擎M3A785GMH/128M 500 内存宇瞻2GB DDR3 1333 300 硬盘希捷ST3500418AS ( 500 GB ) 340 显示器LG GSM4B6F W1942 ( 19.1 英寸) 785（包点）光驱东芝-三星DVD-ROM TS-H352D DVD光驱120 机箱绝尘侠X5 115 电源航嘉冷钻王2.30版本185

抽油机的维护及保养

抽油机的维护及保养（一）抽油机的日检规程抽油机每日应进行下列检查： 1.检查基础是否稳固； 2.检查并禁固地脚螺栓； 3.检查悬绳器和光杆卡瓦是否工作正常，不应有偏斜，不应别别弯光杆； 4.检查驴头销连接是否可靠； 5.检查并禁固连杆固定螺栓； 6.检查并紧固曲柄销； 7.检查并紧固所有构件螺栓； 8.检查各构件是否完好； 9.检查各处轴承，不得有过热现象； 10.抽油机运转不得有异响声； 11.抽油机运转时不得有抖动现象； 12.抽油机运转时减速器不得有震动现象； 13.抽油机运转时三角皮带应张紧适度，不得过紧也不得有打滑现象； 14.抽油机整机应运转平稳。（二）抽油机周检规程在抽油机日检的基础上，每周应增加检查下列项目： 1.检查基础中心线是否正对井口中心，其误差不得超过 3mm；

2.检查底座的水平度，横向偏差不应超过 2/1000 ； 3.检查支架中心与底座标记孔的偏差，不得大于 5mm； 4.检查悬点投影和井口中心的偏差，其偏差圆直径不得大于15mm； 5.检查两连杆内侧到曲柄内加工面对应点的距离，最大差值不得 3mm； 6.检查刹车是否灵活、可靠，刹车间隙是否合适； 7.检查减速器是否需补充润滑油； 8.检查各处轴承是否需要补充加注润滑脂； 9.检查并紧固支架轴承座和游梁轴承座处螺栓； 10.检查抽油机的运转是否平衡，如果抽油机的运转不平衡，则应进行抽油机平衡的调整。 (三) 抽油机的润滑减速器使用润滑油应注意： 1.不准使用不符合标准的润滑油； 2.严禁不同质的润滑油混合使用； 3.季节气温允许的地区，可用另一种牌号的润滑油供减速器全年使用； 4.一种牌号的润滑油不能适应时，应在春秋两季选择合适牌号的润滑油更换； 5.定期取样检查减速器润滑油，当发生下列情况时，应更换

开油水MSDS

开油水（MSDS ）安全技术说明书开油水又名：稀释剂.，是调整油墨粘度，提高印刷适用性。除此之外还能冲淡着色力，增加印刷面积，显著降低成本。开油水：开油水就是稀释油墨或涂料的一种溶剂。在油墨干的时候加入，影响着图案的效果，特别是光滑度和明暗度。 1 化学品及企业标识产品详情商品名称：开油水生产企业名称：邮编： 518010 电话号码： 2 危险概述危险性类别：易燃液体危险有害性信息：易燃液体和蒸汽，口服吞咽有害，接触皮肤有刺激，吸入有害（气体，蒸汽，粉尘，烟雾），对眼睛有刺激。预防措施：远离热源、火花、明火。禁止吸烟。保持容器密闭。只能使用不产生火花的工具。采取防止静电放电的措施。不要吸入粉尘、烟、气体、烟雾、需戴防护蒸气，作业后彻底清洗皮肤。使用本产品时不要进食、饮水或吸烟。手套、防护眼罩以及防护面具。储存：存放在通风良好的地方。保持容器密闭。存放在通风良好的地方。保持低温。 3 成份 / 组成信息成份CAS 号成份重量百分比醋酸乙酯异丁醇防白水醋酸丁酯 4急救措施皮肤接触：用大量清水冲洗。眼睛接触：用清洁的水充分冲洗至少15 分钟并寻求医生的帮助。吸入：脱离现场至空气新鲜处。如呼吸困难，给输氧，就医。食入：如果意外吞下，催吐，就医。

5消防措施灭火剂：二氧化碳，化学干粉，泡沫，水雾。源于此物质或混合物的特别的危害：碳氧化物灭火时可能遭遇之特殊危害：灭火前先停止溢漏，若无法停止溢漏且周围无危险，就让溢漏烧完。若灭火而没有停止溢蒸，可能与空气形成爆炸性混合物而再引燃。特殊灭火程序：在安全情况下将容器搬离火场。用水雾灭火无效，但可用水雾冷却暴露火场的容器。消防人员需着化学防护衣和正压容空气呼吸器( 自携式空气面具) 。消防人员之特殊防护装备：消防人员必须配戴空气呼吸器、消防衣及防护手套。 6泄漏应急处理应急处理：穿戴适当的个人防护装备。环境注意事项：对泄漏区通风换气。移开所有引燃源。通知政府职业安全卫生与环保相关单位。

抽油机现场维护保养及操作流程.doc

抽油机现场维护保养管理及操作流程第一章总则第一条为提高生产设备使用周期，确保生产设备正常运转，做到正确使用、精心维护，使设备长期处于良好状态，结合王沟门采油大队实际情况，特制定本办法。第二条本办法适用于本大队所辖区队、井场设备维护保养管理工作。第二章组织机构及职责第一条大队成立设备维护保养管理领导小组，由分管设备副大队长直接领导。设备维护保养管理分三级管理，即大队、区队、班站。大队设备组和材料组以监督为主、区队维护组负责实施。第二条设备管理目的：安全第一、科学管理、合理使用、高效生产第三条、设备管理措施（一）建立健全设备管理台帐 1、大队、各区队、班站必须建立健全设备管理台帐、档案。

2、设备管理台帐包括：《生产设备维护保养记录》、《设备维修登记》、《井区抽油机档案》、《井区电动机档案》等。 3、设备管理台帐包括内容：（1）设备的名称、类别、数量、规格、型号、统一编号；（2）设备的投产、使用、维护、维修、更换日期；（3）值班人员当班记录，维修、保养、检查、交接班记录。（二）设备的使用管理 1、设备日常管理由所在井区班站长负责；区队设备维护组负责技术指导和维修、监督、检查。 2、各井站职工应熟悉所使用的设备性能、特点和维护、保养要求。 3、所有设备应按照规定要求运行，严禁超负荷运转。 4、所有设备在使用期间要严格按照《设备检修保养规程》的相关规定做好日常维护保养工作，严禁带病运转。 5、设备的检查、保养工作人员应认真填写《生产设备维护保养台帐》及《设备维修记录》。 6、设备管理组应经常检查井区设备维护保养情况及《生产设备维护保养台帐》的填写情况，并做好收集归档工作，要求大队每月不得少于5次，区队每月不得少于10次。 7、新设备安装使用前，必须要有质量检验合格报告。

开油水MSDS

开油水一、化学品名称：开油水二、化学组成：危害物质成分百分比（Percentage for Chemical Ingredient）：≤5% 混合物（Mixing）：丙酮、双丙酮醇三、物理及化学特性: 物质状态：无色液体颜色：无色气味：类似樟脑气味 PH值：5-7 沸点：215℃ 分解温度：50℃—55℃ 闪火点：96℃ 自燃温度：462℃ 爆炸界限：0．84 or 3.84 蒸汽压：0．13kPa(38℃) 密度：0.93（25℃）四、危险性概述: 食入危害健康主要症状：头痛、晕眩、困倦、呕吐、五、急救措施: 皮肤: 立即脱去被污染的衣物，用肥皂和水彻底冲洗去全部外露区域，如果情况加重须立即去医院。使用劳保手套或婴儿油去除已经干燥的胶水。眼睛: 如果胶水进入眼睛或者烟雾引起刺激，立即用水冲洗直至化学物质被消除。如果症状没有减轻，应立即去医院。吸入: 如果症状轻微，转移呼吸新鲜空气；呼吸变得困难，要进行吸氧；呼吸停止，立即进行人工呼吸，实施紧急医疗救助。吞入: 不要诱导自己呕吐。喝水并不要给处于昏迷状态的病人吃任何东西，向中毒控制中心或医生求救。如果无法获得急救建议，应立即将病人和产品送到当地最近的紧急救助中心或医院。六、消防方式: 使用“干冰”或者泡沫灭火。用水可以降温但作为灭火剂可能是无效或低效的。在使用或储存有化学品的地区灭火，消防人员应身穿防火衣服和氧气呼吸器具。火灾和爆炸危险：本品是极易可燃性物体，远离高温和包括电火花、火焰、燃烧的烟头、标灯等热源。避免容器破裂或在高温中暴炸，挥发导致浓度高于空气时应将其转移远离热源和火花处。七、泄漏应急处理：避免无任何防护措施直接接触，避免大量食入；防火、防高温；清理方法：用沙土掩埋后清理。八、储运注意事项：处置：工作区域保持通风良好储存：容器必须紧闭，并处放于5-40℃ 九、废弃处理：

计算机配置参数解释

计算机配置参数解释主频为地频率（如地主频为倍频为外频缓存（）缓存*) 主频外频倍频.也就是倍频是指和系统总线之间相差地倍数，当外频不变时，提高倍频，主频也就越高主频，就是地时钟频率，简单说是运算时地工作频率（秒内发生地同步脉冲数）地简称.单位是.它决定计算机地运行速度，随着计算机地发展，主频由过去发展到了现在地（）.通常来讲，在同系列微处理器，主频越高就代表计算机地速度也越快，但对与不同类型地处理器，它就只能作为一个参数来作参考.另外地运算速度还要看地流水线地各方面地性能指标.由于主频并不直接代表运算速度，所以在一定情况下，很可能会出现主频较高地实际运算速度较低地现象.因此主频仅仅是性能表现地一个方面，而不代表地整体性能. 说到处理器主频，就要提到与之密切相关地两个概念：倍频与外频，外频是地基准频率，单位也是.外频是与主板之间同步运行地速度，而且目前地绝大部分电脑系统中外频也是内存与主板之间地同步运行地速度，在这种方式下，可以理解为地外频直接与内存相连通，实现两者间地同步运行状态；倍频即主频与外频之比地倍数.主频、外频、倍频，其关系式：主频＝外频×倍频.早期地并没有“倍频”这个概念，那时主频和系统总线地速度是一样地.随着技术地发展，速度越来越快，内存、硬盘等配件逐渐跟不上地速度了，而倍频地出现解决了这个问题，它可使内存等部件仍然工作在相对较低地系统总线频率下，而地主频可以通过倍频来无限提升（理论上）.我们可以把外频看作是机器内地一条生产线，而倍频则是生产线地条数，一台机器生产速度地快慢（主频）自然就是生产线地速度（外频）乘以生产线地条数（倍频）了.现在地厂商基本上都已经把倍频锁死，要超频只有从外频下手，通过倍频与外频地搭配来对主板地跳线或在中设置软超频，从而达到计算机总体性能地部分提升.所以在购买地时候要尽量注意地外频. 处理器外频外频是乃至整个计算机系统地基准频率，单位是（兆赫兹）.在早期地电脑中，内存与主板之间地同步运行地速度等于外频，在这种方式下，可以理解为外频直接与内存相连通，实现两者间地同步运行状态.对于目前地计算机系统来说，两者完全可以不相同，但是外频地意义仍然存在，计算机系统中大多数地频率都是在外频地基础上，乘以一定地倍数来实现，这个倍数可以是大于地，也可以是小于地. 说到处理器外频，就要提到与之密切相关地两个概念：倍频与主频，主频就是地时钟频率；倍频即主频与外频之比地倍数.主频、外频、倍频，其关系式：主频＝外频×倍频. 在之前，地主频还处于一个较低地阶段，地主频一般都等于外频.而在出现以后，由于工作频率不断提高，而机地一些其他设备（如插卡、硬盘等）却受到工艺地限制，不能承受更高地频率，因此限制了频率地进一步提高.因此出现了倍频技术，该技术能够使内部工作频率变为外部频率地倍数，从而通过提升倍频而达到提升主频地目地.倍频技术就是使外部设备可以工作在一个较低外频上，而主频是外频地倍数. 在时代，地外频一般是，从Ⅱ开始，外频提高到，目前外频已经达到了.由于正常情况下外频和内存总线频率相同，所以当外频提高后，与内存之间地交换速度也相应得到了提高，对提高电脑整体运行速度影响较大.

抽油机一级保养规范

抽油机一级保养规范注意事项 1.严格按照一级保养周期700～800h进行保养。 2.用手锤检查螺丝紧固时必须敲击螺栓正面，不得敲击棱角，不得逆时针敲击。 3.检查减速箱机油时，一定要注意检查其呼吸阀畅通情况。 4．启抽或停抽时人要侧身操作，操作时戴绝缘手套。 5. 启动抽油机时，严禁抽油机附近站人和有障碍物。 6. 雨雪天要防止电器设备漏电伤人。 7. 检查抽油机曲柄、皮带轮等转动部位防护装置是否齐全、完好，防止曲柄及皮带轮转动部位造成人员伤害。 8. 检查皮带松紧度时严禁手握皮带，以免夹伤手指。 9. 操作刹车锁片时，防止刹车锁片夹伤手。 10.上抽油机操作时不能携带工具，必须用工具袋装好，用绳子吊运，防止工具坠落伤人。 11. 高空作业时脚下要站稳，登高作业超过2m以上必须系好安全带，悬挂高度以高于身高、拉紧为原则，所用工具必须拴牢，以免造成高空坠落事故。 12. 使用工具时要轻拿轻放，避免敲击和碰撞；操作要平稳，规格型号要配套，防止打滑。 13.操作时和操作结束后，必须确认流程正确，观察压力正常后方可进行或离开。

14、一级保养的主要内容就是“十字作业”，即紧固、润滑、防腐、调整、清洁。准备工作 1．劳保用品准备齐全，穿戴整齐 2.工具、用具、材料准备：300mm、375mm、450mm活动扳手各1把，0.75kg手锤1把，试电笔一支，螺丝刀一把，600mm管钳1把，黄油枪1把，安全带2副, 缘手套1副, 钢刷子、毛刷子各1把，工程细线5m，吊线锤1个，清洗桶1只，清洗油适量，黄油适量，机油壶1个，机油适量，测量板一个，直尺一个，棉纱少许，砂纸1张,纸、笔各一。操作步骤 1.检查抽油机刹车锁块是否在行程的1/2-2/3之间，各连接部位完好，灵活好用。 2.用验电笔检测电控柜外壳，确认安全无电。用戴有绝缘手套的手侧身按空气开关停抽，使驴头停在下死点，刹紧刹车，合上刹车锁片。一人系好安全带上驴头，两人配合检查驴头对中情况。 3、打开刹车锁片，缓慢松刹车，当驴头运行到便于操作的位置时（稍高于上死点的位置），刹紧刹车，合上刹车锁片。 4、检查并紧固各部分固定螺栓。检查并紧固基础固定螺栓、检查并紧固支架固定螺栓、检查并紧固平衡块固定螺栓、检查并紧固平衡块锁块固定螺栓、检查并紧固曲柄销子和曲柄轴固定螺栓、检查并紧固连杆固定螺丝、检查并紧固减速箱固定螺栓、检查并紧固刹车装

计算机配置参数解释

计算机配置参数解释 1 主频为CPU的频率（如Intel Core 2 Duo E6300的主频为1.86 倍频为7 外频266MHz L1缓存（KB）32 L2缓存1MB*2) 2 CPU主频= 外频x 倍频。也就是倍频是指CPU和系统总线之间相差的倍数，当外频不变时，提高倍频，CPU主频也就越高主频，就是CPU的时钟频率，简单说是CPU运算时的工作频率（1秒内发生的同步脉冲数）的简称。单位是Hz。它决定计算机的运行速度，随着计算机的发展，主频由过去MHZ发展到了现在的GHZ（1G=1024M）。通常来讲，在同系列微处理器，主频越高就代表计算机的速度也越快，但对与不同类型的处理器，它就只能作为一个参数来作参考。另外CPU的运算速度还要看CPU的流水线的各方面的性能指标。由于主频并不直接代表运算速度，所以在一定情况下，很可能会出现主频较高的CPU实际运算速度较低的现象。因此主频仅仅是CPU 性能表现的一个方面，而不代表CPU的整体性能。说到处理器主频，就要提到与之密切相关的两个概念：倍频与外频，外频是CPU 的基准频率，单位也是MHz。外频是CPU与主板之间同步运行的速度，而且目前的绝大部分电脑系统中外频也是内存与主板之间的同步运行的速度，在这种方式下，可以理解为CPU的外频直接与内存相连通，实现两者间的同步运行状态；倍频即主频与外频之比的倍数。主频、外频、倍频，其关系式：主频＝外频×倍频。早期的CPU并没有“倍频”这个概念，那时主频和系统总线的速度是一样的。随着技术的发展，CPU速度越来越快，内存、硬盘等配件逐渐跟不上CPU的速度了，而倍频的出现解决了这个问题，它可使内存等部件仍然工作在相对较低的系统总线频率下，而CPU的主频可以通过倍频来无限提升（理论上）。我们可以把外频看作是机器内的一条生产线，而倍频则是生产线的条数，一台机器生产速度的快慢（主频）自然就是生产线的速度（外频）乘以生产线的条数（倍频）了。现在的厂商基本上都已经把倍频锁死，要超频只有从外频下手，通过倍频与外频的搭配来对主板的跳线或在BIOS中设置软超频，从而达到计算机总体性能的部分提升。所以在购买的时候要尽量注意CPU的外频。处理器外频外频是CPU乃至整个计算机系统的基准频率，单位是MHz（兆赫兹）。在早期的电脑中，内存与主板之间的同步运行的速度等于外频，在这种方式下，可以理解为CPU外频直接与内存相连通，实现两者间的同步运行状态。对于目前的计算机系统来说，两者完全可以不相同，但是外频的意义仍然存在，计算机系统中大多数的频率都是在外频的基础上，乘以一定的倍数来实现，这个倍数可以是大于1的，也可以是小于1的。说到处理器外频，就要提到与之密切相关的两个概念：倍频与主频，主频就是CPU的时钟频率；倍频即主频与外频之比的倍数。主频、外频、倍频，其关系式：主频＝外频×倍频。

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